Violent Star Formation: From 30 Doradus to QSOs

What impact do massive stellar associations have on their environment through the power of their ionising radiation and mechanical energy? Gathered together in this volume are papers presented at the first IAC-RGO meeting, held in Puerto Naos La Palma, Spain, dedicated to exploring all aspects of this burning question. This volume examines the formation and evolution of new clusters of stars, and explores all the possible consequences in a wide variety of objects where massive stellar bursts have occurred. Thus it presents an alternative model to that of supermassive black holes as the power houses of active galactic nuclei; it analyses the impact of Wolf-Rayet stars, stellar winds and supernovae on their host galaxy; and it provides evidence of massive superassociations and of supersonic velocity dispersions which result from photo-ionisation by violent star formation. For graduate students and researchers, this volume provides a valuable overview and a timely update on all aspects of violent star formation in a host of objects - from 30 Doradus, the largest HII region in the Large Magellanic Cloud, to nuclear starbusts and QSOs.

Violent Star Formation From 30 Doradus to QSOs

Violent Star Formation From 30 Doradus to QSOs Proceedings of the first IAC-RGO meeting, held in Puerto Naos La Palma, Spain November 8-12, 1993 Edited by G. TENORIO-TAGLE Instituto de Astrofisica de Canarias, Tenerife, Spain

CAMBRIDGE UNIVERSITY PRESS

Published by the Press Syndicate of the University of Cambridge The Pitt Building, Trumpington Street, Cambridge CB2 1RP 40 West 20th Street, New York, NY 10011-4211, USA 10 Stamford Road, Oakleigh, Melbourne 3166, Australia © Cambridge University Press 1994 First published 1994 Printed in Great Britain at the University Press, Cambridge A catalogue record of this book is available from the British Library Library of Congress cataloguing in publication data available ISBN 0 521 47277 6 hardback

Contents Participants Preface

xi xiii

Violent Star Formation in 30 Doradus R. C. Kennicutt & Y.-H. Chu

1

The Initial Mass Function of the Center of 30 Doradus E. M. Malumuth & S. R. Heap

13

The Nature and Kinematics of the Emission Nebulae in the Cyg OBI Supershell T. G. Sitnik & V. V. Pravdikova

19

Asymmetry in the Vertical Distribution of Giant Molecular Clouds in the Carina Arm E. J. Alfaro ei al

23

Supersonic Motions in Giant HII Regions C. Munoz-Tunon

25

A Kinematical Study of NGC 604 N. S. P. Sabalisck et al

39

UV Spectroscopy of Giant Extragalactic HII Regions: the Case of NGC 604 L. Sanz Fernandez de Cordoba

41

Evolution of GEHRs: The Effects Caused by Champagne Flows J. A. Rodriguez-Gaspar & G. Tenorio-Tagle

43

Bursts of Star Formation in Central Disks of Galaxies V. Surdin

45

Violent Star Formation G. Tenorio-Tagle

50

Superassociations as Star Complexes with Violent Star Formation Yu. N. Efremov Violent Star Formation Driven by Shock-Shock Collisions A. Chernin & Yu. N. Efremov

61 65

The Search for Hierarchical Structure inside M31 Superassociations P. Battinelli et al

68

A Stochastic PSF Model: Smooth Spirals in Differentially Rotating Disks J. Palous & B. Jungwiert

70

Spatiotemporal Pattern Driven by a Self-Regulating Mechanism of Star Formation A. Parravano

75

Detection of an Age Gradient along the z-Axis in a Star-Forming Region E. J. Alfaro et al

77

Abundances of HII Regions and the Chemical Evolution of Galaxies M. Peimberi ei al

79

vn

viii

Contents

Galaxy Properties in Different Environments: Star Formation in Bulges of LateType Spirals M. G. Pasioriza et al

94

Star Formation in Galaxies in the Bootes Void D. Weistrop

100

Physical Properties of Giant Extragalactic HII Regions A. I. Diaz

105

The Giant HII Region NGC 2363 R. Gonzalez-Delgado et al

117

Photometric Diagrams of NGC 2366 A. Aparicio et al

123

Spectroscopical Imaging of Star-Forming Regions J. M. Mas-Hesse et al

125

A Study on the HII Regions of NGC 4449 0. Fuentes-Masip et al Metallicity Effects on the Properties of Very Young Star Clusters M. L. Garcia- Vargas et al

131 133

The Prototype Starburst Galaxy NGC 7714: Physical Conditions of the Gas and the Stellar Populations R. Gonzdlez-Delgado et al 139 Tracing Violent Star Formation: HI Observations of Nearby Galaxies E. Brinks

145

Massive Star Formation and Supergiant Shells in the Irregular Galaxy NGC 55 D. J. Bomans & E. K. Grebel

156

Galactic Supershells 5. A. Silich el al

162

Violent Star Formation in Dwarf Irregular Galaxies E. Skillman

168

Very Metal-Poor Galaxies and the Primordial Helium Abundance E. Terlevich et al

182

Implications from HI Composition and Lya Emission of HII Galaxies D. Kunth et al

192

The IR- and X-Radiation of the Starburst Dwarf Galaxy UGCA 86 G. M. Richler el al

200

High-Resolution CCD Photometry of HII Galaxies E. Telles & R. Terlevich

202

Formation of Narrow Hell A4686 Emission in HII Galaxies: Link with X-Ray Emission C. Motch et al 208 Environmental Effects in Star-Forming Dwarf Galaxies J. M. Vilchez

214

Contents

ix

Theory of Starburst and Ultraluminous Galaxies B. Elemegreen

220

Colliding Galaxies, Shocked Gas, and Violent Star Formation S. A. Lamb et al

243

Violent Star Formation in Merger Remnants U. Frilze-von Alfvensleben

249

UV Variability of IRAS 13224-3809 J. M. Mas-Hesse et al

256

Star Formation in Polar-Ring Galaxies V. Reshetnikov & F. Combes

258

Infrared Spectroscopy of IR-Luminous Galaxies A. Siernberg et al

263

Application of the Multiphase Model to the Galactic Bulge M. Molld et al

268

Stellar Populations and Population Gradients in Spiral Bulges M. Balcells & R. F. Peletier

270

Implications of Galaxy Alignment for the Galaxy Formation Problem W. Godlowski

275

Annular Structure Analysis of the Starburst Spiral Galaxy NGC 7217 A. M. Varela et al

277

How a Dust Concentration Mimics Dynamical Signatures around the Nucleus of NGC 7331 F. Prada et al 279 UGC 5101: An Ultraluminous IRAS Galaxy with Circumnuclear Star Formation V. Reshetnikov & F. Combes

281

The Stellar Content of Nearby and Distant Starbursts C. Leitherer

283

WR Stars in the Giant HII Region NGC 4236III R. Gonzdlez-Delgado & E. Perez Consequences of High Mass Loss Rates on Wolf-Rayet Populations in Starbursts G. Meynet

289 291

Optical and Ultraviolet Morphology of the Starburst Regions in Wolf-Rayet Galaxies W. D. Vacca 297 The Starburst Nucleus of M83 5. R. Heap

303

Spectrophotometry of Haro Starburst Galaxies 5. Steel et al

309

Starbursts in the Irregular Galaxy VV 523 J. Hecquet et al

317

x

Contents

Long-Slit Spectroscopy of the Central Regions of Starburst Galaxies Henize 2-10 and Markarian 52 H. Sugai & Y. Taniguchi 319 Star Formation in Active Galactic Nuclei: the Cases of NGC 5135, NGC 6221 and NGC 7130 H. R. Schmili et al 325 Metallicity Effects on Starburst M. Cerviiio & J. M. Mas-Hesse

327

From 30 Doradus to QSOs R. Terlevtch

329

Distance Indicators to Low-Luminosity AGN /. Aretxaga & R. J. Terlevich

343

Broad- and Narrow-Band Imaging of the CfA Seyfert Sample A. M. Perez Garcia & J. M. Rodriguez Espinosa

345

Type Transitions in Starburst-Powered AGN /. Aretxaga & R. J. Terlevich

347

Stellar Ionization of Low-Luminosity Active Galactic Nuclei J. C. Shields

353

Line Profiles in Compact Supernova Remnants and Active Galactic Nuclei R. Cid Fernandes & R. Terlevich

365

Composite Galactic Nuclei B. Boer

377

The Nature and Origin of X-Ray Emission in Active Galaxies H. Netzer

379

Starbursts and Compact Supernova Remnants J. Franco et al

387

Broad-Band and Line Emission from Fast Radiative Shocks in Dense Media T. Plewa

396

Study of the Stellar Populations in AGN M. Serole-Roos et al

403

Bidimensional Spectroscopy of Seyfert Galaxies: Offset BLR in NGC 3227 S. Arribas & E. Mediavilla

405

ROSAT Detection of the Most Rapidly Varying Seyfert Galaxy Th. Boiler & J. Trumper

410

QSO Evolution: a Link with Starbursts? B. J. Boyle

413

Evolution of Elliptical Galaxies - a Chemo-Dynamical Model A. C. S. Friaca & R. J. Terlevich

424

Birth of Galaxies at z = 2 or Violent Star Formation at z = 0.4? A. D. Chernin

430

Participants ALFARO, E Instituto de Astroffsica de Andalucia, Spain APARICIO, A Instituto de Astrofi'sica de Canarias, Spain ARETXAGA, I Royal Greenwich Observatory, UK ARMAND, C Laboratoire D'Astronomie Spatiale, France ARR1BAS, S Instituto de Astroffsica de Canarias, Spain BALCELLS, M Groningen University, The Netherlands BALUTEAU, J P Observatoire de Marseille, France BENN, C Royal Greenwich Observatory, La Palma, Spain BOER, B Laboratory for Space Research, The Netherland BOLLER, T MPI fur Extraterrestrische Physic, Germany BOMANS, D J Sternwarte University Bonn, Germany BOYLE, B Institute of Astronomy, Cambridge, UK BRINKS, E National Radio Astronomy Observatory, USA CERVIO, M LAEFF, Spain CHERNIN, A Sternberg Astronomical Institute, Russia CID FERNANDES, R Institute of Astronomy, Cambridge, UK COLINA, L Universidad Autonoma de Madrid, Spain COUPINOT, G Observatoire Midi-Pyrinees, France DIAZ BELTRAN, A Universidad Autonoma de Madrid, Spain DOUGLAS, N Kapteyn Institute Groningen, The Netherlands EFREMOV, Y N Sternberg Astronomical Institute, Russia ELMEGREEN, B IBM Watson Research Center, USA FERLAND, G University of Kentuky, USA FRANCO, J Instituto de Astronomfa UNAM, Mexico FRITZE VAN ALVENSLEBEN, U Universitats Sternwarte, Gottingen, Germany FUENTES MASIP, O Instituto de Astrofi'sica de Canarias, Spain GARCIA VARGAS, M Universidad Autonoma de Madrid, Spain GERRITSEN, J Kapteyn Institute Groningen, The Netherlands GODLOWSKI, W Jagiellonian University Astronomical Observatory, Poland GONZALEZ DELGADO, R Instituto de Astrofi'sica de Canarias, Spain GRIFFITHS, R E John Hopkins University, USA HEAP, S R NASA/Goddard Space Flight Center, USA HECQUET, J Observatoire de Midi-Pyrinnes, France JOUBERT, M Laboratoire D'Astronomie Spatiale, France KENNICUTT, R Steward Observatory, USA KUNTH, D Institute D'Astrophysique Paris LAMB, S University of Illinois, USA LEITHERER, C Space Telescope Science Institute, USA MAIOLINO, R Osservatorio Astrofisico di Arcetri, Italy MALUMUTH, E Goodard Space Flight Center, USA MAS-HESSE, J M LAEFF, Spain MEDIAVILLA, E Instituto de Astrofi'sica de Canarias, Spain MENON, T K Max Planck Institute fur Radioastronomie, Germany MEYNET, G Observatoire de Geneva, Switzerlan MIRABEL, F CEN, Saclay, France MOITINHO, A Universidad de Lisboa, Portugal MOLLA, M Universidad Autonoma de Madrid, Spain MOTCH, C Observatoire de Strasbourg, France xi

Participants

Xll

MUNCH, G MUNOZ-TUNON, C NETZER, H PAKULL, M PALOUS, J PARRAVANO, A PASTORIZA, M PEIMBERT, M PELETIER, R F PEREZ, A PEREZ, E PEREZ OLEA, D PLEWA, T PRADA MARTINEZ, F PRIETO, M REIMERS, D RESHETNIKOV, V RICHTER, G M ROBERTS, W J RODRIGUEZ ESPINOSA, J M RODRIGUEZ GASPAR, J A SABALISCK, N SANZ FDEZ de CORDOBA, L SCHMITT, H R SEROTE, M SHIELDS, J SILICH, S SILLANPAA, A SITNIK, T SKILLMAN, E STEEL, S STERNBERG, A SUGAI, H SURDIN, V TELLES, E TENORIO-TAGLE, G TERLEVICH, E TERLEVICH, R VACCA, W VILCHEZ, J M WAGNER, S WEISTROP, D WHITE, S ZINNECKER, H

Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Tel Aviv University, Israe Observatoire de Strasbourg, France Astronomical Institute, Czechoslovakia Universidad de Los Andes, Venezuela Universidade Federal Rio Grande Do Sul, Brazil Instituto de Astrononna UN AM, Mexico Royal Greenwich Observatory, La Palma, Spain Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Universidad Autonoma de Madrid, Spain Warsaw University Observatory, Poland Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Hamburger Sternwarte, Germany Observatoire de Paris, France Astrophysical Institute Potsdam, Germany Space Telescope Science Institute, USA Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain Instituto de Astrofi'sica de Canarias, Spain LAEFF, Spain Universidade Federal Rio Grande Do Sul, Brazil Observatoire de Paris-Meudon, France Ohio State University, USA Main Astronomical Observatory, Ukrania Turku University, Finland Sternberg Astronomical Institute, Russia University of Minnesota, USA University College Dublin, Ireland Tel Aviv University, Israe National Astronomical Observatory, Japa Sternberg Astronomical Institute, Russia Institute of Astronomy, Cambridge, UK Instituto de Astrofi'sica de Canarias, Spain Royal Greenwich Observatory, UK Royal Greenwich Observatory, UK University of California, Berkeley, USA Instituto de Astrofi'sica de Canarias, Spain Landssternwarte, Heidelberg, German University of Nevada, USA Institute of Astronomy, Cambridge, UK University Wurzburg, Germany

EDITOR'S PREFACE

Violent Star Formation from 30 Doradus to QSOs was the most recent international conference organized jointly by the Institute de Astrofisica de Canarias (IAC) and the Royal Greenwich Observatory (RGO). The meeting took place in Puerto Naos, La Palma (Canary Islands, Spain), in November 1993. This volume contains most of the invited talks and papers presented at the conference. It deals with observations, analysis and theory of violent star formation in its full dimension: from giant HII regions, dwarf and HII galaxies, to starburst galaxies, IRAS ultraluminous galaxies and interacting pairs, all the way up to violent star formation in active galactic nuclei and QSOs. Several critical reviews look at the implications of violent star formation from a variety of angles: ionization structure, interaction with the ISM, hydrodynamics, triggering mechanisms, chemical evolution, luminosity functions and the starburst model for AGNs. On behalf of the organizing committee I would like to thank all the staff and colleagues at the IAC and the RGO who contributed to the organization of the meeting. We are particularly grateful to Monica Murphy (IAC), Judith de Araoz (IAC), Silvia Figueroa (RGO-La Palma) and Rachel Miles (RGO-La Palma) for their efficiency and their keen collaboration to warrant a very pleasant and successful conference. I also want to express my gratitude to Terence Mahoney for his help with the preparation of each of the papers in this volume. The Local Organization was eased by the definite commitment from the Local Government through their Cabildo Insular de La Palma, and the Cabildo Insular de Santa Cruz de Tenerife. It is a pleasure to acknowledge sponsorship from "SOL Hoteles" and the "Grupo SOL", and the local savings bank "Caja General de Ahorros de Canarias", as well as the generous financial support from the Spanish General Directorate for Scientific and Technical Research (DGICYT) and the International Science Foundation (ISF) without which the celebration of the meeting and the attendance of many young colleagues would have been impossible.

Guillermo Tenorio-Tagle Instilulo de Astrofisica de Canarias, Tenerife Spain, April 1994

xin

Violent Star Formation in 30 Doradus By R. C. KENNICUTT 1 AND Y.-H. CHU2 Steward Observatory, University of Arizona, Tucson, AZ 85721, USA 2

Department of Astronomy, University of Illinois, Urbana, IL 61801, USA

The 30 Doradus nebula is the nearest example of a giant HII region, and as such it offers a unique laboratory for studying in detail the structure, stellar content, and dynamics of a starburst region. We begin with an overview of the 30 Doradus region on scales of 0.1-1000 pc, and then discuss two current problems of particular relevance to this conference, the stellar content and IMF in 30 Dor, and the violent dynamics of its interstellar medium.

1. Introduction It is a pleasure to open a conference where 30 Doradus defines the bottom end of the star formation scale! The 30 Doradus region offers a most appropriate starting point for a conference on star formation in galaxies. It is the nearest example of a bona fide giant extragalactic HII region (GEHR), and it is the largest star forming region in the Local Group. It is large enough to exhibit many of the properties of the most luminous starbursts, yet close enough so that its physical structure and stellar content can be studied in detail. As such 30 Dor and other nearby GEHRs provide several crucial pieces of information about starbursts in general. They are the only regions where the embedded stellar population can be resolved and studied directly; this provides a unique stellar census of a starburst, which can be used to test the synthesis models which must be applied to more distant, unresolved GEHRs and starbursts. Imaging and spectroscopy of the gas provides a detailed picture of the physical conditions in the interstellar medium (ISM), hence these regions are the ideal laboratories for studying the complex dynamical interactions between massive stars and the surrounding ISM. Such feedback processes shape the large-scale structure of the ISM, and may well regulate the global evolution of galaxies. 30 Doradus offers all of the ingredients—a highly resolved concentration of thousands of massive stars in a dense, bright, dynamic, ionized medium—to serve as a "starburst Rosetta Stone" (Walborn 1991). Excellent reviews of 30 Doradus have been published recently by Melnick (1987), Walborn (1991) and Meylan (1993), and we refer the reader to those references for a comprehensive review of recent work. In this paper we begin with a brief overview of the 30 Dor region (Section 2), then discuss two specific areas which are especially relevant to this conference, the stellar content and IMF in 30 Dor (Section 3) and the dynamics of its violent ISM (Section 4). 2. 30 Doradus as a prototype giant HII region A useful definition of a giant HII region is any region which contains a sufficient number of massive stars so that its integrated spectra, physical properties, and evolution are dictated by the composite properties of its stellar population (e.g. total stellar mass, IMF, metallicity), rather than the individual properties of its most massive stars, or the local initial conditions in its parent star forming region. By this criterion 30 Doradus easily qualifies as a GEHR, because it contains so many stars that its IMF is fully populated to >100 M@. In this section we describe the 30 Dor region on different spatial scales, and compare it to other nearby star forming regions. 1

Kennicutt k Chu: Violent Star Formation in 30 Doradus

>

•

•

.

I*

•* • . *

FIGURE 1. 30 Doradus and the star cluster NGC 2070. The top panels show a 23' (345 pc) region in Ha and red continuum. The bottom panels are expanded photographs of the central star cluster. Boxes indicate the sizes of subsequent panels. From Kennicutt & Chu (1988).

2.1. Overview of the 30 Doradus region

Figures 1 and 2 show Ha and continuum images of the 30 Doradus region on various scales. The central HII region (top panels in Figure 1) is roughly 15-20 arcmin or 200300 pc across (assumed distance 50 kpc). Projected within the boundary of the nebula are several star clusters and associations, including the central cluster NGC 2070. Figure

Kennicutt & Chu: Violent Star Formation in 30 Doradus

3

2 is a deeper Ha image of a 1600 pc region around 30 Dor, and it reveals that the nebula is located in the midst of a much larger superassociation. Two large HII/OB complexes, 30 Dor B (visible to the SE in Figure 1) and the supershell 30 Dor C, are located in the periphery of 30 Dor and are often considered to be part of the main HII region. Over 20 other large HII regions, OB associations, and star clusters surround this region, mostly to the south of 30 Dor. The ages of these stars range from 0 Myr (including obscured protostellar regions in N159 and 30 Dor itself) to 10-12 Myr (Melnick 1992). Two large expanding supergiant shells, LMC-2 and LMC-3, surround the eastern and western parts of the region, respectively (Meaburn 1980). A faint halo of Ha emission envelops the entire 1 kpc region (Kennicutt k. Hodge 1986); if observed from a large distance the whole complex might well be treated as a single supergiant HII region. The picture of the 30 Dor region which emerges from these data is quite different from the simple single-age stellar population model which is usually applied to more distant GEHRs. Stellar photometry (e.g. Melnick 1987; Lortet & Testor 1991; Schild k Testor 1992; Parker 1993) shows that the region is actually a complex composite of several distinct stellar populations. Even the central nebula contains subclusters with ages ranging over at least 10 Myr. This means that the integrated spectrum of the region (or any similar GEHR) will be dominated by different stellar components at different wavelengths. While the emission-line spectrum of the photoionized gas will trace the youngest stellar population (<3 Myr), the underlying near-UV and visible continua may well be dominated by evolved supergiants from the older population (Searle 1971). The bottom two panels of Figure 1 show enlarged views of the stars in the nebular core. The lower left panel (5.8 arcmin across) is comparable to the region which has been surveyed by Parker (1993), and contains over 2000 stars down to V = 18, not including the compact core cluster around R136. The bottom right panel shows the central 3 arcmin; the unresolved subcluster in the center is R136. It is interesting to note that at a distance of 15 Mpc the bottom panels in Figure 1 would subtend 1.2 and 0.6 arcsec, respectively! Even with the Hubble Space Telescope these regions are unresolvable beyond the Local Group. This emphasizes the importance of objects like 30 Dor for understanding the nature of starbursts on all scales. 2.2. Integrated properties and comparison with other objects

The integrated properties of 30 Dor were tabulated by Kennicutt (1984), and these provide a more quantitative picture of a prototypical GEHR. The extinction-corrected Ha luminosity within a diameter of 370 pc is 1.5 x 1040 ergs s" 1 , which corresponds to an ionizing luminosity of 1.1 x 1052 photons s" 1 . The Ha flux within R < 4' (60 pc) is roughly 1/3 of this total, and is consistent with the ionizing luminosity expected from the ~2400 OB stars within that area (Parker 1992). The total mass of ionized gas in the nebula is roughly 8 x 105 M©. These values are 2-3 orders of magnitude larger than well-known HII regions such as Orion, and several-fold larger than W49 and NGC 3603, the largest HII regions in the Galaxy (Kennicutt 1984). The difference between the LMC and Galactic populations is consistent with a general trend along the Hubble sequence. Although 30 Dor is unrivaled in the Galaxy, it is similar to the largest HII regions in other Magellanic irregulars of comparable mass, and several times less luminous than the brightest GEHRs in giant Sc galaxies such as M51 and M101 (Kennicutt 1988). In addition to the differences in absolute scale, the physical conditions in the GEHRs are quite distinct from those in smaller, compact Galactic regions. The electron density ranges from ~100-300 cm" 3 in the densest filaments to of order 1-10 cm" 3 in the other parts of the nebula (Kennicutt 1984; Mathis, Chu, & Peterson 1985; Rosa & Mathis 1987). This compares to typical densities of order 102 - 104 cm" 3 in compact HII

Kennicutt & Chu: Violent Star Formation in 30 Doradus

FIGURE

2. Ha image of the region surrounding 30 Dor, from Kennicutt & Hodge (1986).

regions. The visual extinction over most of 30 Dor is low, only 0-2 mag (Parker 1993; Dickel et al. 1994), though a few dusty condensations with protostars have been identified (Hyland et al. 1992). As will be discussed extensively at this conference, the ionized gas in GEHRs is much more turbulent, with supersonic velocity dispersions that are 2-3 times higher than in smaller HII regions. These differences are probably related to the concentration of massive stars, which facilitates the transfer of mechanical energy to the ISM from multiple stellar winds and supernovae. We return to this point in Section 4.

Kennicutt & Chu: Violent Star Formation in 30 Doradus

-14

i

-12

1

1

1

*~ o _»-«>

-10

_=™= 0

a ol *a8 °°—=s=0

-8 -6

_ °fl _

°

. 120 Jio . 85

0 * *

—n —

i

^W

O

4 0 -M.o

o

8" o_

o c

1K

o

top

O

_

o—

0

o

2

°o°°°

0

§

°™

-2

0

OnO "oo

_°p 0 o o

O° Oo ^*o o

— i2 o 4o<^

0

0

O

°^

O o Q •* Jll © ° ~> «*

o

o

=;—r~^-7Ao -

njtic

i °x > -

3^>.

(

0

oo°

0

/-

4.9

Sfi .i/. o °

o

_—

j a w "—

-4

^

OU .410

1

1

4.6

4.3

1

" \ 2.5 M.Q I

3.7

3.4

FIGURE 3. H-R diagram for the central 7' region of 30 Doradus, from Parker & Garmany (1993). Sample includes 1230 stars with low photometric errors.

3. Stellar content and IMF As mentioned earlier the proximity of 30 Doradus provides a unique opportunity to measure the stellar population and the form of the initial mass function (IMF) in a starburst region. Numerous authors have resorted to various forms of biased IMFs— top-heavy IMFs, top and bottom-deficient IMFs, or bimodal IMFs— when attempting to account for the extraordinary luminosities of GEHRs and starbursts. The 30 Doradus region is an ideal testing ground for these biased IMF models. Ground-based UBV photometric surveys have been made by several investigators (Lee 1990; La Pierre & Moffatt 1991; Melnick 1992; Meylan 1993; Parker 1993), and HST observations of the core around R136 have been published by Campbell et al. (1992) and Malumuth & Heap (1994). The most complete analysis to date is the thesis study of Parker. He surveyed a 50 arcmin2 region in the center of 30 Dor, as part of a larger survey of the LMC by the Boulder group. CCD photometry in UBV was obtained for a complete sample of 2395 stars brighter than B = V — 18, and used to produce a reddening corrected H-R diagram. Classification spectra were obtained or compiled for ~140 stars brighter than V = 15.5, to provide accurate effective temperatures and bolometric corrections for the hottest and most massive stars. This provided a complete census of the upper end of the IMF for the

6

Kennicutt k Chu: Violent Star Formation in 30 Doradus

center of 30 Dor, except for a few Wolf-Rayet stars and the crowded region immediately around R136 (see below). The final results of the analysis are summarized in Figure 3, which displays the H-R diagram for 1230 stars with the highest quality observations (Parker k Garmany 1993). Evolutionary tracks are superimposed on the data, and the present-day mass function can be derived simply by binning the stars in the respective mass intervals. Parker k Garmany find that the mass function above 12 MQ is well represented by a power law with slope F = d(\ogN)/d{\ogm) = -1.5 ± 0.2. If the cluster is younger than ~3 Myr, as appears to be the case, this PDMF represents the IMF in the region. The slope is considerably shallower than found for the solar neighborhood (F = —2.2) or the LMC outside 30 Dor (F = -1.9) by Blaha k Humphreys (1989), and it is shallower than the mean slope (F = —2.0) in 14 smaller LMC associations measured by Hill, Madore, & Freedman (1994). RecentffST observations allow this analysis to be extended to the dense stellar core of 30 Dor, around the central object R136 (Campbell et al. 1992; Malumuth k Heap 1994). As reported by Malumuth at this conference, the IMF outside of R136 itself (R > 3'-'3 or 0.8 pc) is consistent with the Parker and Garmany results (F = —1.8), but the slope becomes much shallower in the R136 core itself, with F = —0.9 for R < 0.8 pc (Malumuth k Heap 1994). For details we refer the reader to the paper by Malumuth k Heap in this volume. These results taken together lend some support for a top-heavy IMF in GEHRs relative to smaller star forming regions. However it is interesting to note, as pointed out by Parker k Garmany, that previous measurements of the IMF in 30 Dor from UBV photometry (see references above) range over F = —0.8 to —2.0 in roughly the same region! Many of the existing data sets clearly are influenced by systematic errors, and complete homogeneous surveys are needed to test for systematic IMF variations, even in relatively accessible regions such as the Galaxy and the Magellanic Clouds. Highresolution infrared observations, which could extend the mass coverage from ~10 M© (the completeness level of most current surveys) to <1 MQ would be especially valuable (Zinneker 1994, this volume). As long as there is uncertainty in the IMFs determined for highly resolved GEHRs such as 30 Dor, any claims of anomalous IMFs in more distant, unresolved GEHRs and starbursts should be approached with a healthy degree of skepticism. 4. The Violent interstellar medium in 30 Doradus The chaotic filamentary structure of 30 Doradus testifies to its violent internal dynamics. Early Fabry-Perot measurements revealed a supersonic velocity dispersion of 41 km s" 1 FWHM, much higher than those for other HII regions in the Galaxy or Magellanic Clouds (Smith k Weedman 1972), while consistent with a general correlation between linewidth and nebular luminosity among GEHRs (Melnick 1977). High-resolution slit spectra resolved these motions into expanding shells on scales of 1-100 pc (Meaburn 1984, 1988). The detection of several extended X-ray sources in 30 Dor by Einstein (Long, Helfand, k Grabelsky 1981; Chu k Mac Low 1990; Wang k Helfand 1991a) revealed hot gas which is presumably shocked by stellar winds and supernovae embedded in the nebula. These observations point to a dynamical environment which is profoundly different from that in smaller HII regions. In a typical Galactic disk region massive stars form in relatively small associations, and at any particular location the time interval between successive Wolf-Rayet stars or supernovae is usually longer than the lifetime of an in-

Kennicutt & Chu: Violent Star Formation in 30 Doradus

FIGURE 4. Mosaic of 37 echelle line profiles in Ha, covering the central 9' x 9' in 30 Doradus, from Chu fc Kennicutt (1994). Orientation is N top, E left. Each 9' spectrum is separated by 45" N-S. Velocity range in each spectrum is ±300 km s The continuum source near the center is R136.

dividual interstellar bubble or supernova remnant (SNR). In 30 Doradus the situation is dramatically different, with several hundred or thousand OB stars located within a radius of ~l-20 pc. In such environments the surrounding ISM will be influenced by the collective stellar winds from the association, and subsequent supernovae will erupt within large wind-blown bubbles or previously formed SNRs. Such regions are the likely seed environments for the development of the large interstellar bubbles, supershells, and superwinds in galaxies. The ultimate effect of such a GEHR on its ambient ISM is likely to be quite different from an environment where the same number of massive stars are distributed in small, spatially isolated regions.

8

Kennicutt & Chu: Violent Star Formation in 30 Doradus 4.1. Kinematic structure of 30 Doradus

The 30 Doradus region offers a unique opportunity to study these collective interactions between massive stars and the ISM, and the early development of a supergiant shell. As part of a larger survey of the dynamics of the HII regions and SNRs in the Magellanic Clouds, we have mapped the kinematics of 30 Dor using the echelle spectrograph on the CTIO 4-m telescope in a single-order, long-slit mode (Chu et al. 1992; Chu &, Kennicutt 1994). To illustrate the kinematics of the ionized gas in the central 9' x 9' core, Figure 4 shows a mosaic of 37 E-W observations of the Ha line, separated by 3' in the E-W direction and 45" N-S. The velocity range in each individual spectrum is ±300 km s" 1 . R136 is visible as the bright continuum feature near the center of the mosaic. Similar observations were obtained at 53 other positions outside the core. For several regions of interest, echellograms of the [SII] doublet were observed for shock diagnostics and density measurements. The echellograms in Figure 4 reveal the extraordinary complexity of the violent motions in the center of 30 Dor. The outer parts of the HII region (mostly outside the region shown in Figure 4) are characterized by a relatively smooth velocity field, but with a much higher velocity dispersion than in other HII regions in the LMC. As one approaches the center of the nebula the velocity field becomes more complex, with multiple-velocity components at all positions. Several distinct types of kinematic features have been identified (Meaburn 1984; Chu & Kennicutt 1994). These include slow expanding shells with Wexp < 100 km s" 1 and sizes from a few to ~100 pc, and fast expanding shells with Vea;p = 100-300 km s" 1 and sizes from a few to ~30 pc. There are also discrete high-velocity (AV > 100 km s"1) features; some of them are isolated, while others are concentrated within large slow expanding shells, forming complex expanding networks. Several fast expanding shells and complex expanding networks are evident in Figure 4. The slow expanding shells are commonly seen in other regions in the Galaxy or the Magellanic Clouds. The smaller objects (<15 pc) are probably stellar wind-blown bubbles, since the energetic requirements can be met with winds from one W-R star or a small number of OB stars. The larger slow shells (e.g. 30 Dor C outside the main nebula) can be accelerated by a combination of stellar winds and supernovae. The most massive of these, a 20 pc shell expanding at 40 km s~l from the central star cluster, requires an impressive input energy (~ 1052 ergs), but this is not unreasonable given the concentration of massive stars in this region. The fast expanding shells are the most spectacular kinematic features in 30 Doradus, and are only rarely seen elsewhere. No known wind-blown bubbles can match such high expansion velocities. The shell masses range from 100-5000 MQ, and the kinetic energy in the visible ionized shells ranges over 1049 — 1051 ergs. The energy required to produce these shells, whether from stellar winds or supernovae, is roughly an order of magnitude higher, or 1O50 — 1052 ergs. Stellar winds are a less plausible energy source for the fast shells, as the large energy requirement and short expansion ages (104 —105 yr) are usually inconsistent with the number of embedded stars. These fast shells most likely are SNRs inside 30 Dor. The complex expanding networks are unique to 30 Doradus (within the LMC). Each network consists of a large slow expanding shell with additional high-velocity, shocked material within the shell boundary. This structure can be naturally explained by a winddriven superbubble, with SNRs interacting with its shell walls. Chu & Mac Low (1990) have proposed the same mechanism to explain X-ray bright superbubbles in the LMC. Indeed, bright X-ray emission has been detected within the 30 Dor networks, as well as in the aforementioned fast shells.

Kennicutt k Chu: Violent Star Formation in 30 Doradus

9

4.2. Soft X-ray properties of 30 Doradus and surroundings Soft X-ray observations are useful in tracing regions of fast shocks and delineating bubbles of hot gas in a large HII complex. We have obtained a ROSATPSFC pointed observation of the 30 Dor region (Figure 5). The top panel shows the entire 30 Dor superassociation. The brightest sources are the X-ray binary LMC X-l, the known SNRs 30 Dor B and 0540-69.3, and two Galactic foreground stars. If the LMC were observed from a large distance these sources would be unresolved, and a significant fraction of the total flux would be contributed by LMC X-l and the SNRs. Variable interstellar absorption at these low energies (0.1-2.4 keV) further complicates the analysis. This illustrates the difficulty inherent in interpreting X-ray observations of distant GEHRs and starbursts. Figure 5b (bottom) shows a close-up, smoothed X-ray image centered on 30 Dor itself. Within 30 Dor, two types of X-ray sources are visible: point sources coincident with the stellar groupings R136 and R140, and extended diffuse emission associated with expanding shell structures. The fast expanding shells and networks identified in the echelle data have high X-ray surface brightness, while the shells without high-velocity features are much fainter in X-rays. The diffuse X-ray sources in 30 Dor can be explained by SNRs interacting with preexisting superbubbles (Chu & Mac Low 1990; Wang & Helfand 1991b). One such example is 30 Dor C, a superbubble around the OB association LH90; it has a limb-brightened shell structure in the X-ray image (Figure 5a). The SNR activity in 30 Dor C has been unambiguously confirmed by its nonthermal radio emission (Mathewson et al. 1985). A superbubble with SNRs interacting with its shell walls can emit 1035 to several times 1036 ergs s" 1 X-rays (Chu & Mac Low 1990; Chu et al. 1993). The diffuse X-ray emission from 30 Dor, at a level of ~ 2 x 1037 ergs s"1 (Wang & Helfand 1991a), can be decomposed into 5-6 X-ray bright superbubbles (Chu et al. 1994, in preparation). The X-ray observations reveal the prevalent SNR activities in 30 Dor, but also illustrate that the 3-dimensional structure of 30 Dor is dominated by several large shells. 4.3. Integrated kinematic properties We have used the uniform coverage of our echelle data in the central 9' core to estimate the total amounts of energy in the various types of velocity structures, and constrain the physical origin of the supersonic motions. When the echellograms in Figure 4 are summed to produce a single pseudo-integrated profile, the resulting profile has a remarkably smooth structure, as shown in Figure 6. This suggests that the roughly Gaussian velocity structure which is seen in the integrated profiles of more distant GEHRs (e.g. Melnick 1977) is not necessarily the result of virial motions, but rather is produced by the superposition of hundreds of smaller structures which are accelerated by a combination of stellar winds, supernovae, champagne flows, and truly turbulent motions. A point-by-point analysis of the velocityfieldin the 30 Doradus core appears to support this conclusion (Chu & Kennicutt 1994). The background smooth turbulent velocity field is the largest single contributor to the total kinetic energy of the nebula, with ~ 6 x 1051 ergs. The high-velocity shells (Vexp > 100 km s"1) and slower shells each contain approximately 3 — 4 x 1051 ergs, with the shell around R136 alone contributing about 30% of the latter. It is interesting to note that although the high-velocity gas in 30 Dor contains a considerable fraction of the nebular kinetic energy, it represents an insignificant fraction of the integrated line emission. Most of the shells have emission measures of order 100-5000 cm"6 pc, in contrast to 104 — 106 for the foreground nebulosity. The total kinetic energy in the gas, ~ 1052 2 ergs in the central 9', exceeds by at least an order of magnitude the gravitational binding energy of the complex, and suggests that most of 30 Dor, with the possible exception of the core star cluster, is literally blowing

10

Kennicutt & Chu: Violent Star Formation in SO Doradus

star

30 Dor C

FIGURE 5. The 30 Doradus region as observed by the ROSAT PSPC. The top panel shows a 80' x 60' field. The bottom panel is a smoothed expanded plot of 30 Dor itself, covering 42' x 32'. The point sources near the center are coincident with R136 and R140.

Kennicutt k Chu: Violent Star Formation in 30 Doradus 1000

-1—i—1—r

i—i

n—i—r

11

.800

.2 CO o 600 3

I" 400

200

0 -300

, , I -200

-100

0 Velocity (km/s)

100

200

300

FIGURE 6. Ha velocity profile of the central 9' core of 30 Doradus, derived by summing the 37 echellograms in Figure 4. The dashed line shows the profile expanded by a factor of 20, to show the line wings. The peak at -270 km s"1 is geocoronal Ha.

itself apart. On the other hand these energy requirements are in rough accord with the power available from stellar winds and supernovae, as expected if the latter are ultimately responsible for most of the gas motions. For that reason we remain skeptical whether gravitational motion can account for the observed kinematics of GEHRs, though it may contribute to some extent to the lower velocity "turbulent" motions. The mechanical energy in 30 Dor itself is still factors of 5-10 lower than that required to produce the largest supergiant shells in the LMC, such as the large supershell LMC-4 in Shapley Constellation III (Dopita, Mathewson k Ford 1985). However the supergiant shells have characteristic dimensions of 1 kpc (Meaburn 1980), and if one includes the wind and supernova power from the entire 30 Dor superassociation (Figure 2) over a time scale of ~ 107 yr, the energies are comparable (~ 1053 ergs), suggesting that the 30 Dor region will likely form a similar supershell over the next 10-30 Myr. We are very grateful to Eliot Malumuth and Sally Heap for providing preliminary results prior to publication, and to Joel Parker for permission to reproduce Figure 3. RCK is supported in part by the U.S. National Science Foundation through grant AST9019150. YHC acknowledges the support of NASA Grants NAG 5-1900, NAG 5-2112 and GO-4497.02-92A.

BLAHA,

C. &

HUMPHREYS,

REFERENCES R. M. 1989 Astron. J. 98, 1598.

12

Kennicutt & Chu: Violent Star Formation in 30 Doradus

B. ET AL. 1992 Astion. J. 104, 1721. CHU, Y.-H. & KENNICUTT, R. C. 1994 Astrophys. J. In press. CHU, Y.-H., KENNICUTT, R. C , SCHOMMER, R. A. & LAFF, J. 1992 Astion, J. 103, 1545. CHU, Y.-H. & MAC LOW, M.-M. 1990 Astrophys. J. 365, 510. CHU, Y.-H., MAC LOW, M.-M., GARCIA-SEGURA, G., WAKKER, B. & KENNICUTT, R. C. 1993 Astrophys. J. 414, 213. CAMPBELL,

DICKEL, J. R., MILNE, D. K., KENNICUTT, R. C , CHU, Y.-H. & SCHOMMER, R. A. 1994

Astron. J. In press. DOPITA, M. A., MATHEWSON, D. S. & FORD, V. L. 1985 Astrophys. J. 297, 599. HILL, R. J., MADORE, B. F. & FREEDMAN, W. L. 1994 Astrophys. J. In press. HYLAND, A. R., STRAW. S., JONES, T. J. & GATLEY, I. 1992 Mon. Not. Roy. Astron. Soc. 257, 391. KENNICUTT, R. C. 1984 Astrophys. J. 287, 116. KENNICUTT, R. C. 1988 Astrophys. J. 334, 144. KENNICUTT, R. C. & CHU, Y.-H. 1988 Astron. J. 95, 720.

R. C. & HODGE, P. W. 1986 Astrophys. J. 306, 130. LA PIERRE, N. & MOFFAT, A. F. J. 1991 In The Formation and Evolution of Star Clusters (ed. K. Janes). ASP Conference Series, vol. 13, p. 155. LEE, M. G. 1990 PhD thesis, University of Washington. LONG, K. S., HELFAND, D. J. & GRABELSKY, D. A. 1981 Astrophys. J. 248, 925. LORTET, M.-C. & TESTOR, G. 1991 Astro/i. Astrophys. Suppl. 89, 185. MALUMUTH, E. M. &C HEAP, S. R. 1994 Astron. J. In press. KENNICUTT,

MATHEWSON, D. S., FORD, V. L., TUOHY, I. R., MILLS, B. Y., TURTLE, A. J. & HELFAND,

D. J. 1985 Astrophys. J. Suppl. 58, 197. MATHIS, J. S., CHU, Y.-H. & PETERSON, D. E. 1985 Astrophys. J. 292, 155. MEABURN, J. 1980 Mon. Not. Roy. Astron. Soc. 192, 365. MEABURN, J. 1984 Mon. Not. Roy. Astron. Soc. 211, 521. MEABURN, J. 1988 Mon. Not. Roy. Astron. Soc. 235, 375. MELNICK, J. 1977 Astrophys. J. 213, 15. MELNICK, J. 1987 In Observational Evidence of Activity in Galaxies (ed. E. Ye. Khachikian, K. J. Fricke & J. Melnick), p. 545. Springer. MELNICK, J. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto k F. Sanchez), p. 253. Cambridge Univ. Press. MEYLAN, G. 1993 In The Globular Cluster-Galaxy Connection (ed. G. H. Smith & J. P. Brodie). ASP Conference Series, vol. 48, p. 588. PARKER, J. W. 1992 PhD thesis, University of Colorado. PARKER, J. W. 1993 Astron. J. 106, 560. PARKER, J. W. & GARMANY, C. D. 1993 Astron. J. 106, 1471. ROSA, M. & MATHIS, J. S. 1987 Astrophys. J. 317, 163. SCHILD, H. & TESTOR, G. 1992 Astron. Astrophys. Suppl. 92, 729. SEARLE, L. 1971 Astrophys. J. 168, 327. SMITH, M. G. & WEEDMAN, D. W. 1972 Astrophys. J. 172, 307. WALBORN, N. R. 1991 In Massive Stars in Starbursts (ed. C. Leitherer, N. R. Walborn, T. M. Heckman & C. A. Norman), p. 145. Cambridge Univ. Press. WANG, Q. & HELFAND, D. J. 1991a Astrophys. J. 370, 541. WANG, Q. & HELFAND, D. J. 1991b Astrophys. J. 373, 497.

The Initial Mass Function of the Center of 30 Doradus By ELIOT M. MALUMUTH 1 !, AND S A R A R. HEAP 2 'Astronomy Programs, Computer Sciences Corporation 2

The Laboratory for Astronomy and Solar Physics, Goddard Space Flight Center

We report on new Planetary Camera observations of the central region of 30 Doradus. These PC images are the first "deep" HST exposures of 30 Doradus that have appropriate photometric calibration. With R136a at the center of the PC6 CCD chip, the image reveals over 800 stars in a 35"x35"area, and over 200 stars within 3'.'3 of the center of R136a. We used the PSF-fitting method of Malumuth et al. (1991) to measure the magnitudes of all detected stars on the PC6 chip. We used these new B magnitudes, along with U and V magnitudes derived from archived PC images, to derive the luminosity function, mass density profile, and Initial Mass Function of the 30 Doradus ionizing cluster. We find that the mass distribution is like that of a King model, with a core radius of O'/96 (0.24 pc), and a total mass of 17,000 solar masses. Both the luminosity function and the IMF show evidence for mass segregation in the sense that the central region has a higher fraction of massive stars than the outer region of 30 Doradus.

1. Introduction 30 Doradus is one of the most interesting and important objects in the nearby part of the Universe. Walborn (1991) goes as far as to call the 30 Doradus region of the LMC a Rosetta stone for the interpretation of similar, more distant regions. It is no coincidence that 30 Doradus was among the first objects observed with the Hubble Space Telescope. And from the very first HST images, our view of the 30 Dor ionizing cluster began to change. The dense inner core of the cluster, R136a, once thought to be a single supermassive star (Savage et al. 1983), was seen to be a densely packed cluster of several hundred bright, but essentially normal, stars (Malumuth et al. 1991; Weigelt et al. 1991). While much ground based work on the 30 Dor cluster has been done recently (e.g. Parker & Garmany 1993), it is only with HST observations that the nature of the central region can be probed. Indeed, recent papers by the members of the WFPC team (Campbell et al. 1992) and the FOC team (De Marchi et al. 1993) have been published or are awaiting publication. This work is based on a PC image obtained in January 1992. This image is the first deep image of 30 Dor to have a proper flat-field calibration. 2. Photometry The data were obtained using the PC6 chip of the PC and the F439W (hereafter, B-band) filter. Four 180-s and one 30-s images were obtained and combined into a single deep image. The final image is five times deeper than any of the images reported by Campbell et al. (1992). Deep (/-band (F368M) and V-band (F547M) images were obtained as part of the Science Assessment Observation (SAO) program in August 1990. These image were obtained prior to the first WFPC UV flood procedure, so there are no appropriate in-orbit flat-field images. A comparison of our photometry with that of Campbell et al. (1992) demonstrates that it is important to have a deep image in t Address: Code 681, Goddard Space Flight Center, Greenbelt, MD 20771. 13

Malumuth & Heap: The IMF of the Center of SO Doradus

14 -2P

a. o

:

Inside

R136

Q.

- 1 j-•• >

I

•• •

•

• •

. .

•

o

o

i

1 1 1 M 1 1 1

Q_

E

:

*.

,I

1

*

•

•

o "5 -o

•

\ 12

14

16

18

20

V (This paper) FIGURE

1. Comparison of our V-band photometry to that of Campbell et al. (1992) for stars inside R136 (central 7'.'3 x 8'.'7).

order to do accurate photometry in the densest regions, even for the brightest stars. For stars outside of R136 there is no systematic difference, except for a small 0.05-mag color difference, in our photometry of the SAO V-band image and the Campbell et al. V-band photometry. However, for the stars within R136, a 7'.'3 x 8'/7 area, there is a strong systematic difference in the sense that the Campbell et al. magnitudes are smaller (see Figure 1). This difference grows with decreasing brightness, and is what would be expected from an image which is not very deep. Campbell et al. overestimate the brightness of the stars they do detect because of the light contributed by the fainter stars they do not detect. Our photometry suffers from a similar effect; however experiments with test images at the same signal-to-noise ratio as our 5-band image indicate that this bias will be small for stars brighter than about 16.5 mag. For that reason we will use 16.5 mag as a faint end cut-off for determining the luminosity and initial mass functions.

3. Luminosity function In order to compare the properties of the dense inner region of 30 Dor with the less dense outer region we divided our sample in two. All of the stars which are within 3'/3 of R136al comprise the inner sample, while those further away make up the outer sample. For the 5-band image, there are 213 stars in the inner sample and 629 stars in the outer sample. Figure 2(a) shows the luminosity function for the stars in the outer sample. A power law, slope= -0.83, has been fitted to the luminosity function between B =1 2.5 mag and B = 16.5 mag. The fit is displayed in Figure 2(o) as the straight line. The luminosity function for the inner sample is shown in Figure 2(6). Again a power law has been fitted to the data between B = 12.5 mag and B = 16.5 mag, and is displayed as a straight line. The slope of the luminosity function for the inner sample is flatter (slope= -0.66) than that of the outer sample, indicating that there is a larger percentage of bright stars in the inner sample than in the outer sample. Using the

Malumuth & Heap: The IMF of the Center of 30 Doradus 30 Doradus Outer Luminosity Function

15

-0.83

CD

9 1

30 Doradus Inner Luminosity Function -0.66

cn 9 1

10

FIGURE

22

2. Comparison of the B-band luminosity function (a) for stars outside of Rl36a (17'.'5 > R > 3'.'3), and (6) for stars within R136a (R <3'.'3).

Kolmogorov-Smirnov test, which has the advantage that the data are not binned, we can reject the hypothesis that these distributions are the same at the 98% confidence level.

4. Initial mass function Our goal of this work is to determine the initial mass function (IMF) for the center of the 30 Doradus cluster. Since it is not feasible to obtain spectra of more than a few stars in R136, we must rely on evolutionary models of the cluster. First, we convert the observed B magnitudes to absolute magnitude, MB, using a distance modulus m — M = 18.57 (Panagia et al. 1991) and an average reddening E(B — V) = 0.44 (Parker k Garmany 1993). Next, we assume a cluster age of 3 Myr. The presence of Wolf-Rayet stars both within and outside R136a indicate an age greater than 2.8 Myr, while the lack of late O-type supergiants indicate an age less than 5 Myr. Finally, we use the program CLUSTMOD (Landsman 1991) to calculate MB for stars with a range of zero-age masses at 3 Myr using Maeder's (1990) evolutionary tracks for Z = 0.005 and Kurucz' (1991) flux distributions for stars with low Z. Each star is assigned a zero-age mass based upon its magnitude MB . The seven Wolf-Rayet stars (3 in the inner sample and 4 in the outer

Malumuth k Heap: The IMF of the Center of 30 Doradus

16

7.5 T = -1.82

5.5

H

1

1

1

Outer IMF

1

1

T = -0.90

Inner IMF

u a. 7.0 o

6.5 1.2

1.4

1.6 Log Mass (Mo)

1.8

2.0

FIGURE 3. Comparison of the IMF for (a) stars outside Rl36a, and (6) stars within Rl36a (R <3'.'3). In each figure, the solid line represents a power-law fit to the data between B = 12.5 mag (~70 MQ) and B = 16.5 mag (~20 MQ)- The slope of the IMF within Rl36a is much flatter than outside. sample) are put on the high-mass branch, while all of the other stars are placed on the lower-mass branch. For stars which are not Wolf-Rayet stars, the upper mass cutoff is about 70 M Q , however there are few stars more massive than 60 M Q . The resulting IMFs for the outer and inner samples are shown in Figures 3(a) and (6), respectively. The IMF, denoted by £, is denned as the number of stars per logarithmic mass interval per unit area. The slope of the IMF is denned as

where the standard (Salpeter 1955) IMF has a slope, T = -1.35. The slope of the IMF for the outer sample is T = -1.82±0.41. This is in keeping with the slope found by Parker & Garmany (1993), F = - 1 . 5 , for stars in 30 Dor more massive than 12 M 0 , using ground

Malumuth & Heap: The IMF of the Center of 30 Doradus

-1.0

•0.5

0.0

17

0.5

Log r (pc) 4. The radial profile of the surface mass density for the central 17'.'5 of the 30 Doradus cluster. For comparison, the solid line shows a King model with a core radius of 0.24 pc.

FIGURE

based data. The slope derived for the inner sample is much flatter, T = -0.90 ± 0.38. There is a larger percentage of very massive stars within R136a. Again the KolmogorovSmirnov test indicates that we can reject the hypothesis that these distributions are the same at the 95.4% confidence level. Most of the excess massive stars are located right at the very center of the cluster. The IMF for stars in the central ^.'O x 2/.'O area is essentially flat (Heap et al. 1992). Either massive stars are more likely to be formed in the high-density center of the cluster or mass segregation due to dynamical evolution has taken place within the 30 Dor cluster.

5. Mass distribution Given the estimate of the initial mass of each star, we can determine the surface mass density profile by summing the masses of the stars in rings 20 pixels (0/.'88) wide and dividing by the area of the ring. Figure 4 is a plot of the log of the surface mass density as a function of log radius. The line is the projection of a King model with a core radius of 0//96 (0.24 pc). Our core radius is a factor of 4 larger than that derived by Campbell et al. (1992) from their PC images. The large difference is most likely due to their fitting of a King model to the surface brightness profile, instead of the surface mass density profile. As we saw in §3 and §4, the luminosity function and the IMF are both different near the center of the cluster than in the rest of the cluster. Furthermore, the optical luminosity of an OB star does not increase linearly with mass, but instead goes as M 2 2 . As a consequence, the luminosity profile is more sharply peaked than the mass density profile. While 21% of the total B-band luminosity is included in the central l'.'O, only 10% of the total mass is. The question of what is the correct core radius becomes important when we consider the dynamical state of the cluster.

6. Dynamical state of the 30 Doradus cluster One of the most interesting findings of the study is that there is an increase in the percentage of very massive stars near the center of the cluster. The question of why

18

Malumuth k Heap: The IMF of the Center of 30 Doradus

this is so is like the nature vs. nurture questions of human behavior. On the one hand (nature), there may be some density or position-dependent star-formation process that favors the formation of massive stars in the center of clusters or in very dense regions. On the other hand (nurture), the effects of dynamical evolution and mass segregation may cause the most massive stars to sink to the center of the cluster. Campbell et al. (1992) suggest that the relaxation time in the core of 30 Dor is given by the following equation (Spitzer & Hart 1971). = 8.7 x 1 0 6 - ^ _ (6 .2) G mp\nk mlnA Where a — ic\/4irGp0/9, rc is the core radius, m is the mass of a typical star, and p0 is the central density. For the Campbell et al. values of rc = 0.06 pc, p0 = 5 x 104 M© pc~ 3 , m = 10 M 0 , and lnA = 1.4, they get ( r = 3 x 104 yr. The cluster would be about 100 relaxation times old. However, if we use our values for the core radius, rc = 0.24 pc and central density, p0 = 9.5 x 1O3M0 pc~ 3 , we get a relaxation time of tr = 2.6 x 106 yr. The cluster is only one relaxation time old. Since core collapse happens at 12-19 median or half mass relaxation times (Binney & Tremaine 1987), Campbell et al. conclude that the cluster has undergone core collapse while we conclude that it has not. The relaxation time is small enough, however, that dynamical relaxation may be important. tr = 0.34

a

2

REFERENCES J. & TREMAINE, S. 1987 Galactic Dynamics. Princeton University Press, Princeton. CAMPBELL, B. et al. 1992 Astron. J. 104, 1721. DE MARCHI, G., NOTA, A., LEITHERER, C , RAGAZZONI, R. & BARBIERI, C. 1993 Astrophys. J. Accepted. HEAP, S. R., EBBETS, D. & MALUMUTH, E. M. 1992 In Science with the Hubble Space Telescope (ed. P. Benvenuti & E. Schreier), p. 347. European Southern Observatory, Garching. KURUCZ, R. 1991 Private communication. LANDSMAN, W. 1991 Private communication. MAEDER, A. 1990 Astron. Astrophys. Suppl. 84 139. MALUMUTH, E. M., NEILL, J. D., LINDLER D. J., & HEAP, S. R. 1991 In The First Year of HST Observations (ed. A. L. Kinney & J. C. Blades), p. 249. Space Telescope Science Institute, Baltimore. PANAGIA, N., GILMOZZI, R., MACCHETTO, F., ADORF, H.-M., & KIRSHNER, R. P. 1991 Astiophys. J. Lett. 380, L23. PARKER, J. W. & GARMANY C. D. 1993 Astron. J. 106, 1471. SALPETER, E. E. 1955 Astrophys. J. 121, 161. SAVAGE, B. D., FITZPATRICK, E. L., CASSDMELLI, J. P. & EBBETS, D. C. 1983 Astrophys. J. 273, 597. SPITZER, L. & HART, M. H. 1971 Astrophys. J. 164, 399. WALBORN, N. 1991 . In Massive Stars in Starbursts (ed. C. Leitherer, N. R. Walborn, T. M. Heckman, & C. A. Norman), p. 145. Cambridge University Press, Cambridge. WEIGELT, G. et al. 1991 In The First Year of HST Observations (ed. A. L. Kinney & J. C. Blades), p. 208. Space Telescope Science Institute, Baltimore. BINNEY,

The Nature and Kinematics of the Emission Nebulae in the Cyg OBI Supershell By T. G. SITNIK AND V. V. PRAVDIKOVA Sternberg Astronomical Institute, Universitetsky Prospect 13, Moscow, 119899, Russia Detailed Ha line investigations of the gas kinematics in the supershell around the Cyg OBI association were carried out. The supershell contains nebulae and CO-cavities around WR and Oj stars which form a hierarchical system of mutually embedded gaseous dust shells. The nebulae around WR 134, 135, 141 and 142 and the SNR G73.9+0.9 are shown to be located at the far edge of the parent molecular cloud at Vjsr ~ 5 — 10 km s" 1 . We found high negative velocities up to 70 km s" 1 and [OIII]-^Q emission stratification typical for shocks. Both could be associated with stellar wind and SNe. The collective wind and ionizing radiation of the Cyg OBI stars (especially WR) and supernova explosions must play some role in forming the supershell. There are some reasons to suppose that the gas at the sound velocity Viar ~ 12 km s"1 isflowingdowmstream of the ionization front.

Six years ago Lozinskaya & Sitnik, 1988 discovered a hierarchical system of mutually embedded gaseous-dust shells in the Cygnus arm (73 < / < 78°, - 0 * 5< 6 < 3° ). In the sky plane this system consists of several small-size shells around WR and Of stars (NGC 6888 among them) inside the supershell around Cyg OBI association (Figure 1). The supershell (diameter about 100 pc) and inner shells of different sizes are seen as optical ring nebulae, radio-shells, CO-cavities (Lozinskaya & Sitnik, 1988) and IR supershells (Lozinskaya & Repin,1991; Saken et.al, 1992). Analyzing radio and optical data, Lozinskaya & Sitnik, 1988 supposed that the hierarchical system might represent an integral formation created by the collective action of the winds of the Cyg OBI stars. The supershell has been shown to be located at the far side and above the giant molecular cloud. However, at that time we had only a fragmented knowledge about the region near Cyg OBI and we considered the results as preliminary until they could be confirmed by a detailed investigation of the complex's kinematics. Since then we have been carrying out an interferometric Ha survey using the FabryPerot interferometer and an image intensifier with a focus reducing camera mounted on the 125-cm and 60-cm telescopes. The results of these observations can be summarized as follows: (1) the whole region is dynamically very active. The FWHM of the Ha line reaches 100-120 km s" 1 in some positions inside the supershell; the mean value is about 50-100 km s" 1 ; (2) the Ha line emitted in the bright nebulae in the northern, eastern and western parts of the supershell is comparatively narrow and shows little variation; the mean radial velocity of the ionized hydrogen corresponds to V;sr ~ 9-12 km s" 1 , FWHM ~ 50 km s" 1 ; (3) in the southern part of the shell which we have studied in detail the radial velocities of the ionized hydrogen change over a wide range, from —70 km s" 1 to 45 km s" 1 . We have reported several velocity components of the ionized hydrogen in the southern nebulae that we have investigated so far: in the shell around WR 134, SNR G73.9 + 0.9 and near WR 141 and 142 (Lozinskaya, Pravdikova k Sitnik, 1993a, 1993b, 1994). To understand the origin of the different motions we have to clear up the following 19

20

Sitnik & Pravdikova: Cyg OBI supershell

FIGURE

1. The supershell around the Cyg OBI association.

questions: which velocity component corresponds to an undisturbed parent material and which one characterizes the peculiar motions. Let us answer the question using the ring nebula WR 134 as an example (Figure 2). This shell is located in the south-western part of the supershell. The ionized hydrogen here has the following velocity components: 5-7 km s"1 in the bright filaments, -42 km s" 1 to the north from the WR 134 and 12-14 km s"1 almost everywhere inside the shell. If we assume circular Galactic rotation, the interstellar clouds in the Cyg arm towards WR 134 would have radial velocities in the range 0-10 km s" 1 . Indeed, HI and CO observations (Dame et.al, 1988) confirm that the undisturbed interstellar clouds in this direction have the typical velocities from 0 to 10 km s" 1 . Thus the component Vj,r ~ 5-7 km s" 1 which dominates in the brightest filaments corresponds to the radial velocity of the galactic rotation in the shell's location. The CO emission distribution (Dame et.al, 1988) also indicates that the ring nebula is located in a cavity observed at Vhr ~ 0-5 kms" 1 . The fact that the same velocities are found out in the filaments and CO-cavity allows us to suppose that the bright filaments are formed in the dense molecular clouds by the ionizing radiation of the WR 134. The negative velocity component Viar ~ -40 km s" 1 can be connected with the stel-

Sitnik k Pravdikova: Cyg OBI supershell

FIGURE

21

2. The small-size shell around WR 134.

lar wind of WR 134. This velocity component cannot be associated with the Perseus arm, since similar negative velocity was found out in the interstellar absorption lines in the spectrum of WR 134 (St-Louis k Smith, 1991; Nichols-Bohlin k Fesen, 1993). So it is reasonable to suggest that the CO-lagoon has been formed by ionizing radiation and by the stellar wind of WR 134, and that the wind of WR 134 is responsible for the negative velocity component. The following arguments seem to confirm this suggestion. High negative velocities are seen to the north of the star. This is just the region where according to CO measurements and our Ha observations one can expect a higher wind-blown bubble expansion velocities because of the lower ambient density. There is another manifestation of the wind's action. The monochromatic observations of the shell show an [OIII] filament located further from the star WR 134 than the Ha filament (Lozinskaya, Pravdikova k Sitnik, 1993a). Such a stratification of optical emission is typical for the cooling gas behind a shock front formed here by stellar wind. The origin of the third velocity component (Vjsr ~ 12-14 km s"1) is not completely clear yet. It may be connected with the downstream gasflowingaway from the ionization front formed by WR 134. Indeed this velocity is also observed in the absorption spectrum of WR 134 (St-Louis k Smith, 1991; Nichols-Bohlin k Fesen, 1993) indicating that the gas motion takes place in front of the star and is directed toward the star. Moreover this velocity is equal to the sound velocity typical for a downstream flow. In a similar way we have studied the kinematics in the HH-regions near WR 141 and WR 142 and that of the possible supernova remnant G73.9 + 0.9. Thus, the results we have got at present for the supershell and especially for its southern part can be summarized as follows: (1) The Cyg OBI supershell is observed as a shell-cavity at optical, radio and IR wavelengths. The southern part of the supershell is formed by individual objects of a different nature, namely wind-blown shells and supernova remnants. (2) In molecular clouds we discovered small-size cavities around WR stars. They are observed at the same velocities as the bright filaments at the periphery of HII regions. The

22

Sitnik & Pravdikova: Cyg OBI supershell

cavities and bright filaments are traces of the interaction between the ionizing radiation, stellar wind and molecular cloud. (3) In the southern part of the supershell we have found gas motions of a different nature - galactic rotation, peculiar motions at a high negative velocity and at sound velocity. (4) High negative velocity motions up to 70 km s" 1 can be associated with SNR and winds of WR stars. The combined wind and ionizing radiation of these stars and supernova explosions must play some role in forming the supershell. Some manifestations of the stellar wind have been found in the southern part of the Cyg OBI supershell. (5) Gas motions at a positive velocity Visr ~ 12-14 km s" 1 might be caused by the downstream gas flowing away from the "combined" ionization front (like the ring nebula around WR 134). This velocity component is observed in different parts of the supershell as well as in the absorption spectra of the WR and 0 stars of Cyg OBI and Cyg OB3 (St-Louis & Smith, 1991; Nichols-Bohlin & Fesen, 1993). Thus these positive velocities cannot be associated with the background gas or with the receding side of the supershell. But this problem needs further investigation. We consider these results to be a step towards understanding the global and local interstellar structure and of the origin of the the supershell around Cyg OBI. Our investigations of the region are continuing.

REFERENCES DAME, T. M. ET AL. 1987 Astrophys. J. 322, 706.

T. LOZINSKAYA, T. LOZINSKAYA, T. LOZINSKAYA, T. LOZINSKAYA, T.

A. &c SITNIK, T. G. 1988 Soviet Astron. Lett. 14, 100. A. & REPIN, S. V. 1991 Astron. Zh. 67, 1152. A., PRAVDIKOVA, V. V. & SITNIK, T. G. 1993a Astron. Zh. 70, 24. A., PRAVDIKOVA, V. V. & SITNIK, T. G. 1993b Astron. Zh. 70, 469. A., PRAVDIKOVA, V. V. &; SITNIK, T. G. 1994 Astron. Zh. In press. NICHOLS-BOHLIN, J. & FESEN, R. A. 1993 Astron. J. 105, 672.

LOZINSKAYA,

SAKEN, J. M.,

SHULL, J. M.,

GARMANY C. D., NICHOLS-BOHLIN, J.,

Astrophys. J. 397, 537. ST-LOUIS, N. & SMITH, L. J. 1991 Astron. Astrophys. 252, 781.

FESEN, R. A.

1992

Asymmetry in the Vertical Distribution of Giant Molecular Clouds in the Carina Arm ByEMILIO J. ALFARO1'2, JESUS C ABRER A-C ANO 13 AND ANTONIO J. DELGADO1 1

Instituto de Astrofisica de Andaluda, (CSIC), P.O. Box 3004, Granada 18080, Spain 2

Astronomy Department, Boston University, MA02215, USA

3

Universidad de Sevilla, P.O. Box 1045, Sevilla 41080, Spain

The vertical distribution of molecular complexes along the Carina-Sagittarius arm has been studied on the basis of the giant molecular cloud (GMC) data compiled by Myers et al. (1986). The analysis indicates that the CO complexes are preferentially located below the formal Galactic plane. A separation of the sample into two groups: (a) GMCs associated with HII regions, and (b) GMCs without associated HII regions, establishes that group (a) shows, in average, larger z departure and mass than group (b). This result seems to suggest that the star formation activity in this major arm displays a vertical asymmetry which opens up interesting questions about the triggering mechanisms of star formation in spiral arms. The density and location of young stars in major spiral arms and the relation to their parent molecular clouds are important to the understanding of how molecular clouds evolve and form stars in our Galaxy. In previous work (Alfaro et al. 1992, 1993) we analyzed the vertical structure of young open clusters (YOCs) along the optical segment of the Carina-Sagittarius arm, and its connection with the density of YOCs as representative of star formation activity. The main conclusions of that work can be summarized as follows: 1. A clear correlation between YOC density and ^-departure from the formal Galactic plane is found when this density distribution is compared with the vertical structure. The cores of both supercomplexes are closely coincident with the two minima of the vertical profile, and the regions of lowest star-forming tracers appear associated with the relative maximum of z. 2. The z distribution of molecular complexes, aligned along 40 kpc in the arm, shows that the CO complexes are preferentially located below the Galactic plane. This distribution could be biased by: a) the warp of the Galactic disk, and b) the possible out-of-plane location of the Sun. In order to by-pass any possible bias produced by the warp phenomenon, we limited our analysis to the first Galactic quadrant (Dame et al. 1986) where one would expect the warp to be negligible or operate in the opposite direction. Nonetheless, the z distribution for this quadrant also shows the downward trend. 3. The asymmetry would only be apparent if the Sun turns out to be located above the actual midplane. The effect of this upon the result can however be avoided by comparing the z-location of GMCs with respect to each other and not with respect to an absolute reference plane. With this aim we separated the sample into two groups: one containing clouds with HII regions and another are formed by the GMCs with no HII regions. Our analysis led to the result that both groups display a different vertical behaviour, where the probability of finding a difference of means larger than, or equal to, the experimental value is less than 10%, assuming equal means as the null hypothesis. We try to assess it by performing the same analysis for the sample given by Myers et al. (1986). The catalogue of GMCs in the first Galactic quadrant listed by Myers et al. contains 56 objects, out of which 33 CO clouds present association with HII regions and 23 do 23

Alfaro et al.: Vertical Distribution of GMCs in the Carina Arm

24

lolaculor Clouds wlUi araoalatad HH region*

/

-T—v ! \ \

Itoleaular Claud* -without associated HH r*gl<ma

1 1

\

1. Histograms for GMC samples with and without associated HII regions. Histogram in mass (top), in z (middle), and in a linear combination of both variables (bottom).

FIGURE

not. The histogram in z for the two groups (Figure 1, middle), shows that the difference between both distributions, although present, is not so striking as displayed by the data of Dame et al. (1986). The quantitative two-sample analysis led to the result that the null hypothesis cannot be rejected, and that both samples should be considered as representative of the same parent population. On the other hand, a similar analysis for the mass led to the already known result that the two groups are probably different (Myers et al. 1986) in the sense that the more massive clouds are also those more frequently associated with HII regions of high radio continuum emission (Figure 1, top). We try however to go further in discerning the ability of the z variable to discriminate between both groups. With this aim we will test whether some combination of z and mass variables would improve the separation between groups, beyond those obtained with each single variable. The result of this analysis is shown in Figure 1 (bottom) where a linear z-log(mass) combination (Factor 2) represents the new variable. Visual comparison of this plot with the others shows that the combined effect produces a better separation than individual variables. This furthermore indicates that the z variable contains discriminating information, enhancing the one provided by the mass. REFERENCES ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1992 Astrophys. J. 399, 576. ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1993 In Star Formation, Galaxies and the Interstellar Medium (ed. J. Franco, F. Ferrini &; G. Tenorio-Tagle), p. 364. Cambridge Univ. Press. DAME, T. M., ELMEGREEN, B. G., COHEN, R. S. & THADDEUS, P. 1986 Astrophys. J. 305, 892. MYERS, P. C , DAME, T. M., THADDEUS, P., COHEN, R. S., SILVERBERG, R. F., DWEK, E. & HAUSER, M. G. 1986 Astrophys. J. 301, 398.

Supersonic Motions in Giant HII Regions ByCASIANA MUNOZ-TUNON 1

Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain

There are several questions in the field of the study of giant extragalactic HII regions (GEHRs) that remain controversial. The origin of the observed supersonic motions and the validity of the relationship between size and velocity dispersion obtained on the base of single-aperture observations are still debated. From a purely observational point of view, the spatial extent of kinematical features within the nebula, the sense in assuming Gaussian profiles as representative of the variety of emision lines found in GEHR, and finally how lawful the assignment of a single value (IT, .ft) as a defining kinematical parameter of a GEHR are also questioned. The energy input required to provide the supersonic motions observed in GEHRs has been the subject of several scenarios proposed in the literature and some of the points in favour and/or against them are mentioned in this contribution. Bidimensional spectroscopy with good spatial and spectral resolution, sampling of a particular emision line over the whole emitting area, is shown to be the most suitable observational technique for understanding the global kinematics of GEHRs. In NGC 604 for instance, features like loops and filaments although clearly seen in deep exposures and obviously resulting from stellar winds and SN explosions, do not dominate the bulk of the emission that comes from smaller but much brighter areas. Thus the total flux from split or non-Gaussian profiles that arise from shells and/or filaments is much lower than that obtained from bright knots where line profiles are well fitted by Gaussians. An evaluation and comparison of the emission according to line quality for each region is absolutely necessary in order to know the degree of disruption suffered by the remaining cloud under the action of stellar winds from massive stars. Here a new procedure is proposed for defining the size of each region based on its kinematics. The attained values (
1. Motivation: the state of the art Giant Extragalactic HII Regions have typical sizes ranging from 100 pc to more than 1 kpc. The large amount of ionized gas (MHII « 104 to 106 MQ) is embedded within an atomic and molecular cloud about ten times more massive. The high H<» luminosity (1039 erg s" 1 ) is atributible to several hundred massive (0 and B) stars. There is a dependence of the integrated luminosity on the morphological type of the host galaxy; GEHRs in early-type spirals and low-luminosity irregulars are 10 to 100 times less luminous than the above quoted value, whereas luminous Sc and Irregulars show regions 10 to 100 times brighter. Some of the physical properties of GEHR are reviewed by Shields (1990). The shapes and sizes of these large star-forming regions are diverse, and their internal structure is complex. Studies giving values for each region as a whole provide temperatures ranging between 5000 K to 20000 K. The mean electronic density (n e ) goes from 50 to a few hundred cm ~ 3 with typical r.m.s. values between 1 and 10 cm" 3 . The filling factor (nl/{nrm,)2), an observational parameter characterizing the degree of dumpiness of the region, is then / « 10~2. This confirms what on the other hand can be inferred from spatially resolved pictures: GEHRs present a very rich structure, with holes, tunnels, filaments and loops where a small fraction of the nebular volume seems to be ocupied by the high-density ionized gas. Following the advances in astronomical instrumentation together with the increase in 25

26

Munoz-Tunon: Supersonic Motions in Giant HII Regions

telescope diameters, the number of detailed studies of the structure of individual regions is growing. It has been found that there are significant changes in the physical parameters, i.e. temperature, or electronic density, depending on position on the nebulae. 1.1. Supersonic speed (a) Using photographic image plates, Kennicutt (1984) classified giant regions into different Morphological Categories, such as: Giant HII regions (C), Diffuse giant HII regions (D), Multi-core complexes (M) and Shell or ring-like regions (S). Long before, Sandage and Tammann (1974a) following a parallelism with other astronomical objects, gave a practical (observational) definition for these regions. GEHRs were posed as being systems with a Core/Halo structure based on which their core radius r c o r e was defined. Using single-aperture spectroscopy in 1970 Smith k Weedman reported the supersonic behaviour of the acquired velocity dispersion in some of these regions. This feature later on turned into a characteristic common to all of them and in fact should be part of the definition of GEHRs. The classification as supersonic, based on the sound velocity (c,) in a given medium, depends on the system temperature and for the above-mentioned range of temperatures implies velocities which lie between 9.0 km s" 1 and 17 km s" 1 . This physical characteristic of GEHRs left open an additional problem of how to maintain the observed supersonic motions. Several scenarios have been proposed for keeping the ionized gas in GEHRs moving at supersonic speeds. All of them basically use two physical resources, either the gravitational field or hydrodinamical flows. In the first case, the gravitational energy resulting from the potential well of the entire system (the gas-stars ensemble) would make gas clouds move at a supersonic velocity. The second mechanism, hydrodynamical flows, has been employed in two approaches. One of them uses the density discontinuities in the ISM that would favour the champagne phase of the HII region producing high velocity (supersonic) streams. The other one is based on the mechanical energy transferred to the medium by massive stars. The presence of large numbers of massive stars in GEHRs makes it plausible that their winds and SN explosions play a major role in the kinematics of star-forming regions. 1.2. GEHRs as distance indicators Using single-aperture results on the velocity width of the integrated spectrum and a photometric definition for the radius, Melnick (1977) found a correlation between line width and mean core-halo diameter for a group of giant HII regions. Using a larger sample and a more systematic data base, Terlevich k Melnick (1981) reported a a — Ardre empirical law followed by GEHRs. The similarity of this correlation (hereafter referred to as TSiM correlation) with that found in elliptical galaxies, bulges of spirals and globular clusters led them to suggest that giant HII regions are self-gravitating systems. In the same paper by it is also suggested that GEHRs satisfy the Faber k Jackson law (L oc a4) valid for the same virialized systems. The above results presented a new task to be sorted out within the field of star formation. We should understand not only the presence of supersonic motions lasting the whole life-time of the nebulae but also its physical relation between the observed velocity dispersion, luminosity and sizes of emitting objects. The molecular hydrogen in star-forming areas also follows a a — f{R) relationship similar to the T&M correlation (Sanders et a\. 1985). In Figure 1 the two observational laws are plotted together. Until very recently the only explanation proposed in the literature for the T&M correlation was given by the same authors. Following their argument, gravity would have to play a mayor role in big star-forming regions if they follow the same observational law

Munoz-Tunon: Supersonic Motions in Giant HII Regions • A a

Elbpticol o a l a i m GlotxJor clusters Giant HD regions Linear tit to Es Linear tit toEs ondGCs Linear tit toEs andGCs using core radii Linear til to HD regions

27

6166.

o

s

1

log velocity dispersion a (km/s)

FIGURE 1. Correlation between radius and velocity dispersion for globular clusters, elliptical galaxies and GEHRs (from Terlevich & Melnick 1981). Overlapped a similar plot for Galactic molecular clouds; from Sanders et al. (1985).

as globular clusters, bulges and elliptical galaxies. This model has often been referred to as the Gravitational Model (GM). Under the framework of the GM, gas clouds would move under the potential well of the gas-star system . Upon photoionization, the gas velocity dispersion would reflect the velocity distribution of a virialized system, defined by its total mass. The precise modelling of the GM should, however, run into problems. Clouds moving in such a way would frequently collide dissipating their kinetic energy on a relatively small time scale. Thus, unless clouds have an almost negligible cross section, the system would end up in a large cold cloud within the stellar cluster. Winds from massive stars have often been employed in the literature to explain the kinematics of GEHRs. Dyson (1979) developed the Coeval Winds Model (CWM in what follows) in which the kinetic energy observed in GEHRs is provided by a coeval burst of massive stars. A coeval burst of star formation is required to ionize the cloud. The large shell resulting from the combined action of these massive stars would produce an almost Gaussian integrated velocity profile when observed with coarse spatial resolution. The global gas velocity dispersion (cr) would be supersonic. There is an easy way to probe the CWM making use of spatially resolved observational techniques. The line profiles

28

Munoz-Tunon: Supersonic Motions in Giant Ell Regions

obtained when crossing the nebula should vary showing the maximum line-splitting at those positions near the centre of the bubble. This behaviour has been found for instance in Hubble III (Clayton 1987) and other ring-like nebulae but it is not observed in those GEHRs showing other (not ringed) morphologies. On the basis of the CWM, other scenarios based on stellar wind energy have been proposed in the literature, making only massive stars present in GEHRs reponsible for providing the kinetic energy of the gaseous system (see, for example, Kennicutt & Chu, this volume). In any case, the T&M correlation is not addressed by the wind models. A new idea has recently been proposed in the literature by Tenorio-Tagle et al. (1993. Stellar winds from low-mass stars moving under the potential well of the system (cloudstellar cluster) are the ingredients of the Cometary Stiring Model (CSM). For a given scenario of cluster cluster formation implying a critical temperature for star formation, the collection of low-mass stars with an increased cross-section given by their winds, and moving within the gravitational field of the system, would stir the remaining cloud. The cometary shocks, generated by the stellar wind sources ramming through the left-over cloud produce the cloud agitation. When at a subsequent stage massive stars form, the ionized cloud will show the supersonic sigma acquired from the cometary passage of the wind-driven sources. A detailed description of the model as well as the temporal evolution of the cloud is given by Tenorio-Tagle in this volume. The CSM accounts for the empirical correlation both between scale and velocity dispersion, cr, and between luminosity and a found in star-forming regions, molecular clouds and relaxed spheroidal stellar systems.

2. Observational set-up: a new approach The observational approach followed to obtain the above-mentioned results has been made with available techniques. Imaging (photometry) was used to get the core radii from surface brightness distributions. The dimensions (R) of GEHRs have then been obtained with photometric observations, with the assumption that it is possible to define a core/halo structure. Kinematics, mainly velocity dispersion (
Muiioz-Tunon: Supersonic Motions in Giant Ell Regions

29

depends on the chosen area and/or position in the nebulae. In sum, the issue is whether or not it is correct to associate a single (R, a) to a given GEHR. All this new information makes it necessary to define a new observational strategy and data analysis to match if we want to undertake, from the observational point of view, the task of defining the size and line width of GEHRs. Here we propose an alternative approach; spectroscopic 2-D mapping covering the whole region with good spatial resolution is a relevant present-day observational approach. Sampling of one emission line is enough for our purpose, provided it is measured with sufficiently good spectral resolution. The mapping of a region at seeing-limited spatial resolution in one emission line (e.g. H a ) with good spectral resolution should allow us: 1) to know the spatial trend of non-Gaussian asymmetric or split profiles, and 2) to estimate what the percentage of the total flux coming from different zones is, defined according to the shape of the emission lines. The use of the TAURUS2 Fabry-Perot imaging spectrograph (Taylor k Atherton 1980) at the 4.2-m William Hershel Telescope at the Observatorio del Roque de los Muchachos is most suitable for this sort of study. For a detailed description of this instrument see the Observers' Guide (1984) for the Isaac Newton Group of telescopes. The output from this set-up, after phase correction and calibration (Lewis Ic Unger 1991), is a 3-D data cube with x and y being the spatial plane and z the wavelength (etalon step) sampling. In Figure 2 a sketch of an output file is shown. Typical data-cube dimensions are (256, 256, 100) and the chosen etalon free spectral range is about 15 A, enough to sample H a (A = 6562.8 A). Expossure times of about one hour per run are required. The sum of all z planes would produce an image, similar to a narrow-band H a image, called a collapsed image. For a given pixel value (x, y) moving in z provides a spectrum. In summary, each TAURUS run contains typically a total of about 65536 spectra from an area of about 1.1 arcmin2. 2.1. Data Analysis: MATADOR The main problem with TAURUS data comes precisely from its high capability. It provides such an amount of data that often makes you want to run away before any useful information is in sight. After fleeing for quite a while we decided that the available software was not powerful enough to match the potential of the instrument, and therefore we decided to design software that could handle the problem. The outcome was named MATADORf, which was based on the IDL environment and was planned to resolve all present needs while keeping the option open of further extension in the future. In what follows all the data analysis referred to is done with MATADOR . As a first step, and in order to increase the S/N ratio, a global 3-D smoothing is performed. Gaussian functions are taken with widths chosen on the basis of seeing (in x and y) and calibration lamp-profile values (in z). It is important to remark that the line width, i.e. the a of a Gaussian profile, is defined as:

7 = /oexp(-|i).

(2.1)

In the literature, to characterize an emission line, one often finds expressions in terms of the full width at half maximum (fwhm) or the e-folding width. In order to avoid confusion we shall refer to line-width meaning a where, t MATADOR was developed by Vladimir Gavryusev and Casiana Munoz-Tuiion at the Instituto de Astrofisica de Canarias within the framework of the Grupo de Estudios de Forinacion Estelar (GEFE)

30

Munoz-Tunon: Supersonic Motions in Giant Ell Regions

X f (x, y, z ) = f ' ( x , y ) = "Collapsed Image"

256 plx

FIGURE

2. Sketch of phase-corrected TAURUS2 output file.

fwhm = 2>/2 In 2a - 2.35o-

(2.2)

e - folding = y/2
(2.3)

and Results presented here were taken with an IPCS-CCD, with typical seeing values of 0.7 arcsec and a spatial scale, oversampling the seeing, of 0.26 arcsec pixel" 1 . Spectral scanning AA/step = 0.16 A, equivalent to 7.3 km s" 1 in H a . Instrumental width, taken from a 2-D mapping of the calibration lamp, was ffinst = 1.75 ± 0.15 step. The width value, as directly measured on the observed line, will be called a. The result, after substracting thermal and instrumental widths will then be due to mass motions and is accordingly called
" 2 ="L» + ^ + "L,-

(2-4)

Once a data cube is smoothed two kinds of analysis can be undertaken using MATADOR, either global or detailed. Within the Global Analysis Procedure it is possible to do automatic 2-D fitting with a single Gaussian . A threshold, mask or criterion to set a limiting value to the pixels to be fitted can be defined. For instance, a signal threshold can be used as the starting point in order to avoid the fitting to noisy regions on the image. The output from the global fitting procedure are eight 256x256 maps showing the Gaussian fitted parameters: (1) peak intensity, (2) central wavelength, (3) sigma, (4) background, and (5-8) four more maps with the errors corresponding to fitted parameters. The Detailed Analysis Procedure allows one to choose the spectra to be studied by defining the position on the image and the beam size. It allows for the definition of the analytic function and number of them to be used if one wishes to fit the emission profile. These allow us to classify emission flux according to line quality and to redefine the size of the emitting region with a kinematical criterion.

Munoz-Tunon: Supersonic Motions in Giant Ell Regions

31

3. Some results for N G C 5461 and N G C 5471 Here we present some results on the two brightest GEHRs in M101, at a distance of 7.2 Mpc. The collapsed images in H a of NGC 5461 and NGC 5471 are presented in Figures 3 and 4. In Figure 3(b), for comparison with the TAURUS data, the narrow-band H a image from GEFE f collaboration is also shown. Nodes labelled A, B, C, D and E, as given by Skillman (1984), are marked on the image. The pixel size of TAURUS "images is 0.26 arcsec, equivalent to 9.05 pc linear scale in the galaxy, slightly different from that corresponding to the GEFE image. The comparison of both images in Figure 3 shows the global consistency in morphology as measured by two intrinsically different observational techniques. Despite the good quality of the GEFE image, the higher spatial resolution and seeing values of the TAURUS data allow us better to resolve some features in the emission structure. The collapsed image in H o of NGC 5471 is shown in Figure 4. On the image, knots A, B, C and D as, defined by Skillman (1984), are marked. This region, the brightest in M101, looks like an almost circular halo with several bright knots. For this reason it is considered a multi-core GEHR. The spatial resolution and seeing conditions of the observations presented here allow the identification of two knots, A and A', in what previously was classified as a single structure. 3.1. The kinematics of NGC 5461 From the sigma map of NGC 5461 (not presented here) the main results to be pointed out are: • Knots E and D are not connected with the main region. They present a much lower velocity dispersion within the range of normal (not supersonic) HII regions and are possibly physically separated from the rest of the region. This is also apparent in our Figure 3, although it is not until the velocity dispersion map is obtained that one can definitely establish that: • Knot C is also kinematically detached from the rest of the region. • By a kinematical criterion, NGC 5461 is not a region of lOOOpc size. The supersonic sigma values restrict NGC 5461 to knots A and B only. This reduces the size of this GEHR to about 500 pc in diameter. • "Well behaved Gaussian emission-line profiles extend over a diameter of about 500 pc and it is possible to assign a velocity dispersion value of amm s» 25 km s" 1 to the emission coming from this area. • On the edges of the 500 pc region there are poorly defined profiles. The intensity is much lower than that of the bulk of the region and asymmetries, line splitting and peculiar shapes are found. No detailed analyses have been performed due to S/N limitations. With our experimental approach it is then possible to assign a (D, a) value to NGC 5461. The size D — 500 pc was defined following the kinematics of the region and is smaller than previously reported, as measured using imaging techniques. Areas without welldefined line profiles are found bordering the region, but the bulk of the emission (in terms of intensity and covering area) comes from an area of 500 pc in diameter and can be represented by well-behaved Gaussian profiles with a
Munoz-Tuiion: Supersonic Motions in Giant HII Regions

32

- 1

200

250

500

450

400

350

300

150

200

250

300

350

FIGURE 3. (a) Collapsed image of NGC 5461 H a TAURUS data-cube. lpixel=0.26arcsec=9.05pc. (b) Narrow-band H a image of NGC 5461 from the GEFE collaboration.

Munoz-Tuiion: Supersonic Motions in Giant HII Regions

33

• All emission profiles from the sampled area are supersonic. • A large fraction (« 80%) of the emitting area presents a roughly constant emissionline width with Gaussian profiles characterized by a amm w 20 km s" 1 . • Split line profiles have also been detected in a small area on the border of the nebula. Their normalized peak intensity is « 1/20 of that of the bright regions. • There is a fraction of the emitting area (w 20%) with broader profiles. This includes the above mentioned split lines. A further study of the broader profiles over NGC 5471, results from the global analysis fitting procedure. The emission can be classified into two groups: 1) Broad unsplit lines. These arise from two small areas close to, but slightly displaced from, the two brightest knots (A and B). The broad profiles in knot A seem consistent with the molecular core disruption as proposed by Rodriguez-Gaspar & Tenorio-Tagle (this volume). Only knot B shows clear evidence of multiple components in the line profile. It is clearly non-Gaussian and seems consistent with the SN detected in the same area. 2) Lines not well represented by single Gaussians, coming from low surface brightness zones and identified with the split profiles reported above. Several limitations exist also in 2-D Fabry-Perot spectroscopy. One is the small spectral range which may hide the presence of low-intensity very broad components. In NGC 5471 (Castaneda et al. 1990; Mas-Hesse et al. 1994) have found the very broad (1000 km s"1) low (less than 20% of the peak) intensity component in node C, that we fail to detect in our present data. To sum up, it could be said that the a value measured in NGC 5471 using single aperture includes areas (20% of the total) with wider profiles. These localized broader profiles may be explained by winds and/or champagne flow mechanisms. Most of the emitting area however presents a (rmm(region) ss 20 km s" 1 and covers a radius of > 200 — 300 pc. An explanation other than winds from massive stars is necessary to account for the width of the emission measured in most nebular areas. The features from massive star effects, although detected, are only located on a small fraction of the nebula. 4. The case of NGC 604 NGC 604 in M33 at 720 kpc, ten times closer than M101 has been used extensively as a laboratory to study the ISM. The richness of its morphology, which can be resolved with our spatial sampling, makes this GEHR most suitable for the sort of kinematic study proposed here. Generally speaking, it has a core/halo structure. The halo as seen in H a can be described as a low-intensity background against which various loops and filaments can clearly be detected. The core is far brighter than the structures in the halo and can be resolved into two bright knots about 50 pc apart. Detailed analysis of the spectra arising from several areas on the image provides an important result; emission profiles from the core and the halo are different in shape. Line splitting and poorly defined Gaussian profiles are found in the halo, whereas the core mainly the two bright knots - has well defined Gaussian emission profiles. Figures 5 and 6 present, together with two collapsed images of NGC 604, spectra corresponding to several positions in the nebula. The four spectra shown in Figure 5 correspond to the four boxes drawn on the image. These emission profiles are typical of the halo region. As remarked above, line profiles are far from being Gaussian, displaying line-splitting and strong asymmetries. However, the two spectra in Figure 6, corresponding to the two emission peak maxima, are very similar and absolutely Gaussian in shape.

34

Munoz-Tunon: Supersonic Motions in Giant HII Regions 120 140

100

120

100

80 80

60

40

40

20 20

20

40

60

80

100

120

140

FIGURE 4. Collapsed image of TAURUS data-cube in HQ for NGC 5471. 1 pixel = 0.26 arcsec = 9.05 pc.

In both figures numbers in the grey level intensity scale represent the total number of counts (in detector units) measured within the FSR. In the spectra, "array index" means the z plane in etalon steps and "averaged intensity" is the normalized intensity for each given box. The most important result in NGC 604 is that typical line profiles arising from the core Ipeak(core) are more than 10 times brighter than halo emission lines / pea t(/ia/o); and core emission lines and halo emission lines differ in shape. Single-aperture spectra display an observed velocity dispersion 80%) presents asymmetries and line-splitting with a large variety of line profiles. • "Well-behaved" Gaussian profiles are associated with an area of D « 60 pc and o-mm = 18 km s" 1 .

Munoz-Tufion: Supersonic Motions in Giant HII Regions

35

Spactrum ot x-113 y-226 WnX- 3 bJnV- 5 MnZ» 1

250 '

> * .

:'• - '''''••}>''''.*'•

'

150

200

150

100-5

100

50 50

o I

.

.

.••.-

i

50

'•R-IVI

100

i'l'i

f'jl

150

200

250

FIGURE 5. Collapsed image of NGC 604 in HQ; 1 pixel = 0.26 arcsec = 0.86 pc. Four spectra, taken from several positions in the halo, are also plotted corresponding to boxes marked on the image. Intensity is given in arbitrary units and is normalized to the bin size.

5. Final Remarks The story started with the results obtained after single-aperture observations of GEHRs. The measured supersonic line widths led people to look for a physical mechanism to fuel the supersonic motions in giant star-forming regions. Time went by and we were able to see not only further out but also better. Things that seemed to be hidden before began to appear in our CCD images. Loops, filaments, rings and many other nice structures were disclosed from what before was a boring background. That sorted out part of the problem. We know, or at least we think we know, the necessary physics to be able to understand the interaction of massive stars with the interstellar medium. We are almost able to model the precise way in which these loops, filaments and rings are built. It is possible to recognize the features left by stellar winds and SN explosions in the ISM. Thus, the mechanisms that provide the energy for supersonic motions in quite a large fraction of the emitting area of some GEHRs is already known. Meanwhile, the T&M correlation was found. This relates the size of the brigthest zones in each GEHR and the measured velocity dispersion. Moreover, and despite the increase in spatial resolution, most GEHRs present areas - coincident with the brightest zones - where we fail to recognize the signature features of the massive stars. These areas, that could be named the kinematic cores, still present a supersonic velocity dispersion. Therefore, the TScM correlation and the origin of the observed supersonic motions in the kinematic cores of GEHRs have become the present observational challenge in the field.

36

Mufioz-Tunon: Supersonic Motions in Giant Ell Regions Spectrum ot x« 98 y» 155 X S binY5 binZ- 1 binX

250

150

200

150

100-5 Spectrum at x> 149 y-152 binX= 5 binY= tinZ= 1 2

100

'

Jl ' \

0

50 6

50

4 2

50

100

150

200

250

J

\ 40 Arroy lnd«

80

FIGURE 6. Collapsed image of NGC 604 in H a . Two spectra of the bright knots in the core are also plotted, corresponding to the boxes marked on the image. Intensity is given in arbitrary units and is normalized to the bin size.

Two-dimensional spectroscopic mapping covering the whole region with good spatial resolution is a very adequate observational approach to advancing our understanding of the kinematics of GEHRs. Such observations and the use of the data analysis procedure proposed here allow us to classify emission flux according to line quality and to re-define the size of the emitting region with a kinematical criterion. The kinematics of NGC 5461, NGC 5471 and NGC 604 have been studied and presented here as an example of our proposed approach. We are developing a similar study over a large sample of GEHRs and although clearly we do not yet have statistically significant results here we present a tentative list of conclusions: • After a spatially resolved study, it is permitted to assign a value of (
Muiioz-Tunon: Supersonic Motions in Giant Ell Regions

37

20

* %

/-fyi

~*lV

y&h

»VAJ,

-

^'"jr'kTi

i

J

>•

r

J

?

V

-20

^ '

»

-40

FIGURE

7. Narrow-band Ha image of NGC 588 in M33 from the GEFE collaboration.

the energetics of massive stars. But we are no longer observing the kinematical cores of the regions and thus the measured (
REFERENCES H. O., VILCHEZ, J. M. & COPETTI, M. V. F. 1990 Astrophys. J. 365, 164. CHU Y. H. & KENNICUTT, R. C. 1994 Astrophys. J. In press. CLAYTON, C. A. 1987 Mon. Not. R. astron. Soc. 226, 493. DYSON, J. E. 1979 Astron. Astrophys. 73, 132. CASTANEDA,

38

Munoz-Tunon: Supersonic Motions in Giant HII Regions

KENNICUTT, R. C. 1984 Astrophys. J. 287, 116.

R. C. & CHU, Y. H., 1994 This volume. J.& UNGER, S., 1991 William Hershel Telescope. TAURUS data and how to reduce it. MELNICK, J. 1977 Astrophys. J. 213, 15. KENNICUTT,

LEWIS,

MAS-HESSE, M.,

MUNOZ-TUNON, C ,

VILCHEZ, J. M., CASTANEDA, H. O. & CARTER,

D,,

1994 This volume. G., 1994, This volume. N., TENORIO-TAGLE, G., CASTANEDA, H. O. & MUNOZ-TUNON, C , 1994 This volume. SABALISK, N., TENORIO-TAGLE, G., CASTANEDA, H. O. & MUNOZ-TUNON, C , 1994 Astrophys. J. Lett. Submmitted. SANDERS, D. B., SCOVILLE, N. Z., SOLOMAON, P. M. 1985 Astrophys. J. 289, 373. SANDAGE, A. & TAMMANN, G. 1974 Astrophys. J. 190, 525. SANDAGE, A. & TAMMANN, G. 1974 Astrophys. J. 194, 223. SHIELDS, G. A., Ann. Rev. Astron. Astrophys. 28, 525. SKILLMAN, E. 1984 PhD Thesis. TAYLOR, K. & ATHERTON, P. D. 1980 Mon. Not. Roy. Astron. Soc. 191, 675. TENORIO-TAGLE, G. 1994 This volume. TENORIO-TAGLE, G., MUNOZ-TUNON, C. & Cox, D. P. 1993 Astrophys. J. 418, 767. TERLEVICH, R. & MELNICK, J. 1981 Mon. Not. R. astron. Soc. 195, 839. UNGER, S., BRINKS, E., LAING, R. A., TRITTON, K. P. & GRAY, P. M. 1988 Isaac Newton Group, La Palma. Observers' Guide. RODRIGUEZ-GASPAR, TENORIO-TAGLE,

SABALISK,

A Kinematic Study of NGC 604 By N. S. P. SABALISCK, G. TENORIO-T AGLE, H. O. CASTANEDA and C. MUNO Z-TUNON Instituto de Astrofisica de Can arias, E-38200 La Laguna, Tenerife, Spain We present a kinematical study of NGC 604 in the emission lines of Ha and [OIII] A5007 A, based on data obtained with TAURUS-2, a Fabry-Perot imaging spectrograph. The main result is that the global supersonic velocity value is dominated by emission coming from the two maximum brightness peaks covering an area of about 12 arcsec2. NGC 604, located in M33, is one of the most prominent giant extragalactic HII regions (GEHR). It was observed with TAURUS-2 at the 4.2-m William Herschel Telescope (La Palma), with the Image Photon Counting System (IPCS) as detector. Two data cubes, each with an integration time of 3600 sec, were obtained in the emission line of Ha and [OIII] A5007. The 125-/im etalon was used, which gave a free spectral range for Ha of 17.6 A (806.6 km s"1) and 10.0 A (596.1 km s - 1 ) for [OIII]. The velocity resolution (a of the calibration line) was 23.1 km s" 1 for Ha and 19.6 km s" 1 for [OIII]. In order to determine the origin of the supersonic mass motion in NGC 604, the kinematics of the region were analyzed by means of individual spectra, applying automatic fitting routines with single Gaussian fits. The measured velocity dispersions were corrected for instrumental and thermal broadening. The thermal velocity dispersion correction was calculated assuming Te = 10,000 K. The integrated spectrum of the whole region shows a global heliocentric radial velocity of —244 km s" 1 and a velocity dispersion (
Sabalisck et al.: A kinematic Study of NGC 604

40

NCC604 - Ha

:

•

I

n.

V

. .r 0

10

20

30

40

VELOCITY DISPERSION

SO

SO

0

(km/s)

i 10

NGC604 - [0llI]5007k

20

30

40

VELOCITY DISPERSION

SO

SO

(km/s)

1. Histogram of velocity dispersion versus the sum of the integrated flux for (a) Ha and (b) [OIII] A5007 A. The sum of the integrated flux is taken at 2 km s" 1 intervals of a. Velocity dispersion values are corrected for thermal and instrumental broadening. FIGURE

NGC 604 - Ha emission line (BO B )

100

NGC 604 - Ho emission line (15.2")

200 -

77

a 100 13.3 E a = J7 5 k n/i gtoDai a

pixel

160

200

km/s 100 pixel

200

12

FIGURE 2. (a) - Ha flux maps of NGC 604, used to divide the region into 49 squares of 60.8 arcsec2 each (marked with a thin dotted line). We obtained the flux and a of each part to calculate the global a in different flux ranges, identified with different kinds of lines. In (b) we have plotted the most intense area, divided into squares of 15.2 arcsec2. Only the zones with more than 70% of the maximum flux values are shown here. Their corresponding global a is 16.5 km s" 1 . The X-Y axis of these maps are in units of pixels (1 pixel = 0"26).

The result implies an intrinsic mechanism agitating the high-density gas left over from the violent star forming process. We gratefully acknowledge partial support from DGICYT Grant PB91-0531(GEFE) and the NATO Grant CRG920198 for collaborative research.

REFERENCES HIPPELEIN, H. & FRIED, J. W. 1984 A. & A. 141, 49. MELNICK, J. 1977 Ap. J. 213, 15. ROSA, M. & SOLF, J. 1984 A. k A. 130, 29.

UV Spectroscopy of Giant Extragalactic HII Regions: the Case of NGC 604f ByLOURDES SANZ FERNANDEZ DE CORDOBA Lab. Astroffsica Espacial y Fi'sica Fundamental, Instituto Nacional de Tecnica Aeroespacial, Apartado 50727, 28080-Madrid, Spain The UV spectrum of a giant extragalactic HII region is compared here with the synthetic spectrum generated by a fitting procedure which uses a combination of stellar spectra to match the observed spectrum.

1. Observations The ultraviolet observations of giant extragalactic HII regions provide very valuable information on the stellar population embedded in these complexes. The UV spectra of HII regions in the IUE short- wavelength range, SWP (1152-1950A), show a stellar continuum which rises steeply towards the shorter wavelengths. Superposed over this continuum, emission lines and P-Cygni profiles are often observed, as an indication of the presence of hot luminous early-type stars with mass loss. NGC 604, the brightest HII complex in the nearby galaxy M33, is, because of its size and distance, a good candidate for a detailed study of a giant extragalactic HII region in the UV based on WE data. NGC 604 has been observed with IUE on different occasions in low resolution (6 A mm"1) with the large aperture slit (10x20 arcsec2). A total of 16 SWP and 8 LWR, or LWP (1950-3200 A), has been obtained. The spectra correspond to observations of different zones within the region. The general characteristics of the IUE spectra of NGC 604 in the short wavelength range are similar to the ones mentioned above. In the near UV, the observed continuum is practically constant, no emission lines are detected, there is no 2200-A absorption feature, and only interstellar absorption lines are present. 2. Analysis of the data For the detailed qualitative study of the stellar population of this giant extragalactic HII region in the UV, a program for reproducing the UV observed spectrum with a certain combination of individual stellar spectra was developed. This fitting program is based on a least squares fit to a non-linear function. An iterative process minimizes x2 using a combined method of gradient search plus an analytical function. Input data are • the UV spectrum of NGC 604, corrected only for foreground galactic extinction, E(B - V) = 0.03, and • UV spectra of luminous stars of different spectral types (between O3-B5 with UV luminosity classes s and g, and WR stars, N and C types), selected from the IUE-Archive and corrected for distance and reddening. The output result is a composite spectrum, the optimized fit for the set of standard f Based on observations from the International Ultraviolet Explorer satellite. 41

42

Sanz Fernandez de Cordoba: UV Spectroscopy of Giant HII Regions

20.0

0.0

1400.0

1600.0

1B0O.O

Wavelength |A]

FIGURE 1. Comparison between an observed UV spectrum of NGC 604 (solid line) and the optimized composite spectrum obtained by the fitting process (dotted line). stars being used as input. The quality of the resultant fit is a function of the selected set of standard stars.

3. Results Figure 1 is a graphical comparison between a UV spectrum of NGC 604 and the composite spectrum obtained at the end of the fitting process. In this case, the synthetic spectrum results from a combination of O5d, 05s, O7d, O9d, WN7 and WC7 stars selected as input. As shown by this figure the observed spectra of NGC 604 can be satisfactorily fitted to the synthetic spectra, not only in the continuum but also in the observed stellar emission lines and P-Cygni profiles. With this systematic fitting procedure the IUE UV spectra of NGC 604 were analized and compared in detail in order to obtain the luminous and young stellar population content in the complex (Sanz Fernandez de Cordoba 1986). A more detailed presentation and discussion of the results obtained is to be published.

SANZ FERNANDEZ DE CORDOBA,

REFERENCES L. 1986, PhD Thesis, Univ. Granada, Spain.

Evolution of GEHRs: The Effects Caused by Champagne Flows By J. A. RODRIGUEZ-GASPAR AND G. TENORIO-TAGLE Instituto de Astrofi'sica de Can arias, E- 38200 La Laguna, Tenerife, Spain We analyzed the possible evolution of a high surface brightness GEHR into an extended diffuse region due to champagne-produced outflows. We calculate the typical time-scales and the final radius and density for a variety of initial parameters. In addition, a numerical calculation which allows visualization of the evolution is shown. The Giant extragalactic HII regions (GEHRs) in nearby spiral, irregular and blue compact galaxies have widely been studied during the last years due to their remarkable properties, such as their large luminosities, diameters and supersonic gas motions (see review by Shields 1990). Kennicutt (1984) analyzed the structural properties of GEHRs and classified them morphologically as "classical", "high-surface-brightness", "diffuse", "multiple core", "shell" and "supergiant". Here, we suggest a time sequence from the "high-surface-brightness" to the "diffuse" class by the effect of champagne flows. Let us assume a spherical cloud that includes an extended region with radius rc and constant density n c and a central core with a power-law density stratification, '-\niO(r/rior

for ri0 < r < ric

(UA

>

where rj c = r,o (»iio/>»c) , is the external radius of the core. Outside the cloud there is the ISM with a constant density noThe cloud is ionized by a powerful cluster of O5V stars located at the center and with a Lyman continuum photon luminosity of Ni(O5 V) = 5 • 1049 sec" 1 each one. Eventually, the whole cloud becomes ionized and the gas in the core and at edge of the cloud is set into motion because of the pressure gradients. The inner core is disrupted, producing a strong shock whose evolution is given by Franco, Tenorio-Tagle k. Bodenheimer (1990) as r,(<) = r
=

- u;)]

(0.2)

At the edge of the cloud the usual conditions for producing a champagne flow are present. The pressure discontinuity develops into a strong shock that moves into the ionized ISM and a rarefaction wave (RW) that moves inwards into the cloud. The shock strength or Mach number M, is given by Bedjin k Tenorio-Tagle (1981) as »c/»io = Mj exp \{Mj — 1)/M,]. The RW causes a continuous streaming of material out of the cloud and the high-density cloud expands into a more extended and "diffuse" region. The time necessary for this is set by the speed of the RW, tc — rc/c, during which the radius of the HII region grows to rcj = r c ( l + Ma), and the average density falls to n c no/nc + M, . = (0 3) 3

"•'

MT(I

+

MS) -

'

The time necessary to dissipate small condensations inside GEHRs is smaller than the 43

Rodriguez-Gaspar & Tenorio-Tagle: Evolution of GEHRs

44 CORE

Core Disruption

.\

i-

•I

Cloud expansion

.^ !

CLOUD

r

i

•*•

/ '

'•

1

1 1

I

:

1

J

1

ISM

!

o.l

-10

H-'M.U-^.^I O" 2

1.2x10

A

I.OXIO4

i "I

\ i

1

8.0X103 RodH (pc)

Rodii (pc)

FIGURE 1. Hydrodynamic models and surfaces brightness for H(v in erg cm" sec ' sr ' . The initial parameters are: u,o = 106 cm" 3 , nc = 10J cm" 3 , no = 1 cm" 3 , r,o = 2 • 1016 cm , Tc = 3.86 • 1019 cm and NL = 1051 sec"1 (w 20 05 V stars). The evolutionary times are: 0 (initial condition which is represented by a solid line), 3.8 • 104 (dashed line), 1.8 • 105 (dotted line) and 6.3 • 105 (dash-dotted line) yrs.

GEHRs' evolution time. Thus, the "high-brightness" regions having small and denser condensations, unresolved observationally, could evolve into "diffuse" regions with a nearly uniform density. For example, for 30 Doradus, tc = 8.5 • 106 yr (we assume nc « 10 cm" 3 and rc w 100 pc; Kennicutt 1984) while the time neccesary to dissolve a dense core (n,o = 103 cm" 3 and TJO = 1 pc) is 3.7-105 yr, about a 30% smaller. However, this value is only a lower limit to the disruption time of the core. A hydrodynamical simulation using the procedure described by Tenorio-Tagle et ai. (1986) showed excellent agreement with previous estimates. Figure 1 shows the plots of density, velocity, temperature and Hcv brightness at selected evolutionary times. This work is supported by DG1CYT (grant PB91-0531GEFE) and NATO (grant CRG920198 for collaborative research).

REFERENCES BEDJIN, FRANCO,

P. J. &; TENORIO-TAGLE, G. 1981 Astron. Astrophys. 98, 85. J., TENORIO-TAGLE, G. & BODENHEIMER, P. 1990 Astrophys. J. 349, 126.

KENNICUTT, R. C. 1984 Astrophys. J. 287, 116. SHIELDS, G. A. 1990 Ann. Rev. Astron. Astrophys. 28, 525. TENORIO-TAGLE, G., BODENHEIMER, P., LIN, D. N. C. & NORIEGA-CRESPO, A. 1986

Not. R. Astron. Soc. 221, 635.

Mon.

Bursts of Star Formation in Central Disks of Galaxies By VLADIMIR SURDIN Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899 Moscow, Russia Central gaseous disks around the nuclei of flat galaxies continually increase their mass due to spiralling giant molecular clouds (GMCs) under the action of dynamical friction. The radius of the disk depends on a tidal condition in the central parts of a galaxy equal to GMC tidal disruption distance. A central part of the disk can become molecular and be able to undergo a subsequent spontaneous burst of star formation when the mean surface density of the disk becomes larger than the critical UV-opacity column density. In the outer parts of the disk, formation of gaseous clumps and stellar aggregates can be self-consistent with a characteristic clump mass of 107 MQ and new-born stellar masses of 106 MQ in each clump. This scenario is a good approximation to the observable characteristics of central molecular disks of normal galaxies like ours. However, the interaction of galaxies must modify the maximum cloud and stellar aggregate mass up to ~ 108 M Q and lead to stimulated bursts of star formation.

1. Introduction Bursts of star formation are frequently observed in the nuclear regions of galaxies. Observations have revealed compact nuclei surrounded by an extended disk, or disk-like HII regions with radii of about 200-500 pc (Wilson et al. 1991). These regions are, as a rule, the brightest HII regions in such galaxies. One of the main reasons for such intense starbursts is the high star formation efficiency (SFE) in galactic nuclear regions (Planesas et al. 1989). Given the high SFE in the elementary cells of star formation there are dense clumps (n ~ 104 cm"3) in GMCs - it is not surprising if huge analogs of the clumps, the central molecular disks, are located in the central region of some galaxies and have masses of a few 108 MQ. We know examples of this kind of disk in our Galaxy (Bally et al. 1988), M82 (Sofue et al. 1992; Wild et al. 1992), and several more. External parts of the molecular disk are consistent due to the presence of cold neutral gas. In some cases these regions share a few cloudy complexes with active star formation, and mimic multi-nuclei galaxies. In this paper I will try to find answers to the following questions: - What physical mechanisms lead to the particular parameters and structure of the central molecular disk (CMD) and the central gaseous disk (CGD)? - Why does the process of star formation near galactic cores occur as a large-scale outburst? - Why is the SFE of the process so high? 2. Formation of central disks There are some ideas towards the exploration of star formation bursts in isolated galaxies such as, for example, NGC 253 (Rieke et al. 1980). It may have been an externally stimulated process through the accretion of intergalactic gas clouds fueling the nuclear region and probably triggering star formation bursts (Sofue &; Wakamatsu 1991). But in isolated galaxies one should consider the angular momentum problem, and the main cause of a starburst must be inner dynamical processes, such as gas viscosity in a differentially rotating disk (Lesch et al. 1990), dynamical friction of GMCs (Surdin 45

46

Surdin: Bursts of Star Formation

1980a, 1980b, 1981; Stark et al. 1991), and angular momentum loss in the potential of a central bar (Hernquist 1989). Of the three mechanisms dynamical friction is the only one that works without any requirements regarding the particular form of the galactic potential. This mechanism provides a natural explanation for the absence of HI and CO emission between R = 3 — 4 kpc and R = 1.5 kpc in the Galactic disk. When GMCs spiral into the galactic center under the action of dynamical friction, they are destroyed by the tidal gravitational field of the galaxy at some distance from the galactic center. The condition for tidal stability requires that the mean density of a gravitationally bound cloud (pc) exceeds the galactic mean density inside the cloud's orbit (pG) (Spitzer 1987): Pc>ApG,

(2.1)

where A is a constant depending on the galaxy mass distribution: A = 3 for a Keplerian rotation curve, and A = 2 for constant rotational velocity (Rastorguev k Surdin 1978). If equation (2.1) is used for the galactocentric distance at which a GMC is disrupted, one can obtain the radius of the central gaseous disk (RCGD) in the simple case of constant circular velocity Vo in the galactic disk: RCGD = 1 ( . . . , ° _ t ) ( i n " p

M C

_3 )

kpc,

(2.2)

V200 kms V \100 H2 cm A J where UGMC is the mean GMC density. For our Galaxy Vo — 200 km s" 1 and UGMC — (50 — 200) cm" 3 , therefore RCGD = (0.7— 1.4) kpc, in good agreement with observations of the HI circumnuclear disk of our Galaxy. Intensive spontaneous formation of molecular clumps inside the CGD can develop only if the mean surface density of the CGD is larger than the critical high-opacity column density required for molecule formation:

300 (^

J MQ pc" 2 ,

(2.3)

where Z is the metallicity of the gas. From (2.2) and (2.3) we can obtain the maximum mass of the CGD just before starburst:

" (f

10-

The local SFE value depends first of all on the mean density of a cloud. It is because the star formation process is interrupted by the energy output of massive stars. If the cloud's free-fall time, tjj ~ (Gpc)" 1 / 2 , is greater than the pre-main sequence evolution time of massive stars (tpms ~ 106yr), then the stars of a first generation would have enough time to destroy the cloud by their energy output. That is why GMCs as a whole with tff « 7 • 106 yr have SFEs as low as 0.1 - 1% (Surdin 1989). But the cores of GMCs with tjf « 4 • 105 yr have SFEs as high as 20 - 60%. A possible explanation of nuclear starbursts is the delay of clump formation under the tidal action of a galactic nucleus. If the condition of high SFE is taken into account: tff < tpms

(2.5)

and tidal condition (2.1), we can obtain the radius of the central molecular disk where a starburst can occur: RCMD = 0.4 (

^

\200 km s

_,)

kpc

/

(2.6)

Surdin: Bursts of Star Formation

47

and estimate the mass of the starburst disk:

(f)«0.

(2.7)

We may also evaluate the CMDs mean star formation rate (SFR): SFR = SFE- McMDT~\

(2.8)

and luminosity (L): L = io^.(^iA

T

\

MQ

(2.9)

)

where r is the duration of the starburst. If we assume that all massive stars are born at the same time, then r is the main-sequence lifetime of the stars: r m ,(30 MQ) = 5x 106 yr. With a value of SFE ~ 100% we obtain 4Q yr" 1 \zuu Km s

/

(2.10)

\ ^ /

and

L~W42( ^ T) f^ergs" 1 . \ 2 0 0 k m s - V \ZJ&

(2.11) '

v

These formulae give us a good approximation to the observed characteristics of CGDs, CMDs and starbursts in the circumnuclear regions of disk galaxies. 3. Self-induced formation of gaseous clumps and stellar aggregates In principle, the star formation process in the central part of a galaxy may be selfstimulated. It is well known that in molecular rings almost each GMC would create a young stellar aggregate which will evolve into an OB association and/or an open cluster. Threse newborn stars inject sufficient radiative and mechanical energy (stellar wind) to destroy their parent molecular cloud and create a rapidly expanding interstellar shell. As a result of the shell expansion a dense banana-shaped envelope forms in differentially rotating galactic disks (Tenorio-Tagle & Palous 1987). A large part of the interstellar matter swept up during the envelope evolution accumulates at the tips of the "banana remnant" and leads to the formation of GMC, which can begin a new cycle of star formation. In this picture, star formation and the formation of GMC are a self-connected process. It can take place not only in the molecular ring of a galaxy but in the central gaseous disk too. We can evaluate typical masses of clumps (C) and stellar aggregates (S) in a process of consequential formation. The evolution of supershells embedded in differentially rotating disks was analyzed by Tenorio-Tagle & Palous (1987) and Palous et al. (1990). For disks with flat rotation curves Palous et al. (1990) obtained a simple formula for the mass of a molecular cloud formed on the tip of banana-like envelope: M c = 6 x

lcm

where h.4 is the half-thickness of the disk, n0 is the density of interstellar matter, and E is the total input of energy. On the other hand, let us consider a gravitationally bound cloud of mass Me, and radius Re- Its binding energy is \Uc\ = GMQ/RCHOW many stars can be created in the cloud before it is destroyed by the stellar energy input? We can write the energy

48

Surdin: Bursts of Star Formation

input as Es=n(3E0Ms, (3.13) where EQ is the output energy of one massive star, \i is the fraction of this energy conserved in the hot bubble of gas, /? is the number of massive stars per unit mass of young stars, and Ms is the total star mass. Star formation will cease under the natural condition (Surdin 1989): \UC\ = ES. (3.14) From this equation, for EQ = 10 51 erg, fi = 0.1, and /? = 10~ 2 (Larson 1974), we have:

J^^OOHTcm-3)173'

<3-15>

where nc denotes the mean density of the clump. In CGD we can take nc ~ 103 — —104 cm"3, and finally we have:

Taking into account the equality fiE — Es, one can construct, from (3.12), (3.13), and (3.16), a system of two equations with two unknowns Ms and Me-

Ms

\

( Mc

\5/3 (

Mc

0.44 /

\ 0.54 /

F

\ 0.66

where agas is the total column density of interstellar matter perpendicular to the disk. The solution of the system is: Mc = 7 x 103 Mo (r^—)

( ,/ga'_2) 3.4

(J.lOj

Because the described proccess of cloud formation is stimulated, the value of Me is the maximum mass of a cloud in their population. Let us consider, for example, the central part of our Galaxy (R < 1 kpc). For this case agas ~ 150 MQ pc~ 2 and ha ~ 40 pc (Sanders et al. 1984). From (3.18) we can predict the maximum mass of a cloud: Me — 2 x 107 MQ. It is only a few times bigger then the mass of the Sgr B cloud. The solution (3.18) also predicts a rather high SFE (= Ms/Mc ct 15%) in the CGD of the Galaxy.

4. Discussion As it has been shown that the main characteristics of central gaseous disks of galaxies can be explained by the tidal instability of spiralling GMCs, the critical value of the UV-opacity column density, and the delay of clump formation in the disk under the tidal action of a galactic nucleus. This last mechanism provides the formation of clouds with a mean density of the order of 104 cm" 3 only. The free-fall time of such clouds is less than the pre-main sequence evolution time of massive stars. This may be the main reason for a high value of SFE during nuclear starbursts. We have also considered a self-consistent stimulated process of cloud and stellar aggregate formation. In a stable

Surdin: Bursts of Star Formation

49

regime the process produces low-mass clouds (Me ~ 107 MQ in our Galaxy) and not very abundant stellar aggregates (Ms ~ 106 MQ). Finally, I would like to mention the dependencd of Ms and SFE from the vertical scale of the disk: Ms oc hil

(4.19)

and SFEoc/i 16 . (4.20) The gravitational interaction of galaxies may lead to an increase in velocity dispersion and the vertical scale of gaseous disks by large factors; for example, in, numerical experiments of Elmegreen (1993), up to 2 — 3 times. It can also lead to an increase in the value of Ms of up to ~ 108 MQ and in SFE up to 80%. In other words, it can trigger a burst of active star formation. Participation at the Conference was made possible in part by Travel Grant No 1005/4 from the International Science Foundation.

REFERENCES A. A., WILSON, R. W. & HENKEL, C. 1988 Astrophys. J. 324, 223. B. 1993 This conference. HERNQUIST, L. 1989 Nature 340, 687. LARSON, R. B. 1974 Mon. Not. Roy. Astron. Soc. 169, 229. BALLY,

J.,

STARK,

ELMEGREEN,

LESCH, H., BIERMANN, P. L., CRUSIUS, A., REUTER, H. P., DAHLEM, M., BARTELDREES, A.

& WiELEBlNSKl, R. 1990 Mon. Not. Roy. Astron. Soc. 242, 194. J., FRANCO, J. & TENORIO-TAGLE, G. 1990 Astron. Astrophys. 227, 175. PLANESAS, P., GOMEZ-GONZALEZ, J. &; MARTIN-PINTADO, J. 1989 Astron. Astrophys. 216, 1. RASTORGUEV, A. S. & SURDIN, V. G. 1978 Astron. Tsirk. 1016, 3. PALOUS,

RIEKE, G. H., LEBOFSKI, M. J., THOMPSOM, R. I., Low,

F. J. & TOKUNAGA, A. T.

1980

Astrophys. J. 238, 24. SANDERS, D. B., SOLOMON, P. M. & SCOVILLE, N. Z. 1984 Astrophys. J. 276, 182. SOFUE, Y., REUTER, H.-P., KRAUSE, M., WIELEBINSKI, R. &: NAKAI, N. 1992 Astrophys. J. 395, 126. SOFUE, Y. & WAKAMATSU, K. 1991 Publ. Astron. Soc. Jap. 43, L57. SPITZER, L. 1987 Dynamical Evolution of Globular Clusters. Princeton Univ. Press. STARK, A. A., GERHARD, O. E., BINNEY, J. & BALLY, J. 1991 Mon. Not. Roy. Astron. Soc. 248, 14P. SURDIN, V. G. 1980a Astron. Tsirk. 1113, 3. SURDIN, V. G. 1980b Astron. Tsirk. 1126, 1. SURDIN, V. G. 1981 Astron. Tsirk. 1178, 6. SURDIN, V. G. 1989 Astron. Nadir. 310, 381. TENORIO-TAGLE, G. & PALOUS, J. 1987 Astron. Astrophys. 186, 287. WILD, W., HARRIS, A. I., ECKART, A., GENZEL, R., GRAF, U. U., JACKSON, J. M., RUSSELL, A. P. G. & STUTZKI, J. 1992 Astron. Astrophys. 265, 447. WILSON,

A. S.,

HELFER,

T. T.,

HANIFF,

G. A. &

WARD,

M. J. 1991 Astrophys. J. 381, 79.

Violent Star Formation ByGUILLERMO

TENORIO-TAGLE

Instituto de Astron'sica de Canarias, E-38200 La Laguna, Tenerife, Spain The quasi-static gravitational collapse of clouds should lead to fragmentation into small Jeansunstable cloudlets with a column density similar to that of the clouds as a whole. This criterion defines the mass and size of newly-formed fragments as the collapse proceeds. One could in principle expect that all collapsing matter should end up in a massive singularity at the centre of the collapsing configurations. However, the onset of star formation is shown here to stop the collapse. This should happen if the newly-formed stars produce winds while ramming through the left-over cloud and in this way cause the stirring required to stabilize the collapse. An inmediate consequence of this is the velocity dispersion generated in the star-forming region, which is supersonic in the case of massive clusters and detectable upon the appearance of massive stars in giant HII regions. The stirring cause by the supersonically moving wind-driven sources is also shown to cause a distinct cloud structure, or filling factor, in excellent agreement with recent observations of regions of violent star formation. The two effects, i.e. the acquired velocity dispersion and the filling factor caused in the parent cloud, allow us to differentiate between high- and low-mass clusters. In both cases, however, the disruptive energy from massive stars ends up erasing the clues stored in the gas during cluster formation.

1. Introduction There are several trivial but definitive conclusions regarding the history of spheroidal stellar systems that one can infer from the present agglomeration of stars in globular clusters and galactic bulges. In both cases for example, it is quite obvious that star formation stopped ages ago. It is also evident that star formation took place over a sufficiently long period of time, or that there must have been such strong and sudden bursts as to convert large fractions of the original cloud masses into stars, allowing the clusters to appear today as bound stellar systems. However, most important of all, is the fact that there must have been a stabilizing agent. Something stopped the collapse of the clouds as a whole and did not allow the formation of supermassive objects at the centre of the configurations. Instead, what seems the standard answer of Nature to such a situation is what actually happened: plain, simple fragmentation! This led to innumerable smaller and smaller self-gravitating cloudlets out of which eventually resulted the ensemble of stars that today seem to have relaxed their spatial and velocity distributions into a virialized state. If one could follow the history of stellar clusters far back in time, one would eventually find the short time interval over which their massive stars were evolving along the main sequence. The left-over gas would then have appeared as giant HII regions while the heating caused by the ultraviolet light would have inhibited any further stellar formation, thereby defining the efficiency of the process. In other words, by then, all matter that had to go into stars, had gone into stars already. The little remaining gas exposed to the activity of the massive stars, to their winds and supernova explosions, was then dispersed, or blasted away, leaving all stellar sources (massive and low-mass stars) moving under the gravitational potential of the newly-formed cluster. It would have been then, during 50

Tenorio-Tagle: Violent Star Formation

51

the lifetime of the massive stars, that the gas velocity dispersion (o-ga,) had been found to be correlated with the size of the HII regions in the same manner as one finds today that the stellar velocity dispersion (cr») correlates with the size (r) of spheroidal stellar systems (a* ~.Ar 1 / 2 ). Going even further back in time, before the appearance of massive stars, the clouds must have been infested by a large number of low-mass stars, the stars that we see today, all moving within the gravitational potential of the collapsing clouds. Many of these sources undergoing winds while ramming through the leftover cloud, caused a continuous stirring of the leftover gas, inhibiting the further collapse of the clouds as a whole, and at the same time providing their distinctive structure, or cloud filling factor. This however, could have happened only once the cloudlets resultant from fragmentation had acquired stellar sizes (M* ~ MQ). Surely before then, fragments were too large and too massive to form stars and thus the collapse of the clouds and their fragmentation unavoidably continued. It has proven difficult to trace the history of clusters. However, during the last few years a major effort has been channelled into understanding all possible clues, particularly those offered to us during one of the most obvious stages, during the giant HII region phase. These, if rightly interpreted, may offer a definite path towards understanding the cluster formation sequence. Giant HII regions have been studied as large-scale events showing us large sections of galactic disks - the threshold column density - capable of undergoing major bursts of stellar formation. Also, as entities demanding a large number of OB stars to account for their large-scale ionization and excitation, and they are indeed the best laboratories to study the mechanical energy deposition from massive stars. Several studies have also centred their attention on the detected supersonic turbulence. Without exception, the line profiles of giant HII regions imply a supersonic velocity dispersion (cr) in the range 15-40 km s"1 , and the nature of these supersonic motions, also termed supersonic turbulence, has received several interpretations during the last twenty years. These have attempted to point to the energy source, the essential ingredient in all models, capable of providing continuous replenishment. This comes from the fact that supersonic motions cause shock waves that rapidly dissipate energy; thus, a continuous injection of energy seems unavoidable if one is to maintain the shock waves and supersonic turbulence, at least during the lifetime of the HII region. Some authors have suggested the action of multiple strong stellar winds from massive stars as the driving energy source (e.g. Rosa & Solf 1984; Dyson 1979; and more recently Kennicutt & Chu 1994). Others have invoked the champagne expansion of the HII regions as responsible for the supersonic flow (Skillman & Ballick 1984; Hunter & Gallagher 1985). There is no doubt that both mechanisms do lead and contribute to supersonic motions within the ionized gas. However, none of these models or interpretations has attempted to offer an explanation of the observed velocity dispersion vs. size correlation. On the other hand, Terlevich & Melnick (1981) have suggested that supersonic turbulence in giant HII regions ultimately ought to have a gravitational origin. This conclusion seems unavoidable if one looks at the positions occupied by giant HII regions in the log r vs. logcr diagram, with globular clusters and the central regions of elliptical galaxies at opposite ends of the correlation. This fact was originally interpreted as a collection of ionized clouds moving supersonically under the gravitational potential of the newlyformed stellar cluster, a model often referred to as the "gravitational model". This is, however, totally ill-founded as ionized clouds should be completely disrupted through champagne-likeflowsin a few sound crossing times, and hence the model cannot explain the persistence of supersonic turbulence throughout the lifetime of giant HII regions. The only way out is to generate these supersonically moving clouds continuously during

52

Tenorio-Tagle: Violent Star Formation

the life time of the exciting stars. The good news however, is that perhaps everybody was right! According to the cometary stirring model (see Tenorio-Tagle et a\. 1993, hereafter referred to as paper I), gravity causes the stellar velocity dispersion, and the stellar winds from the many low-mass stars to agitate the left-over cloud leading to supersonic turbulence and density enhancements and thus to the observed correletation. This possibility surpasses the obvious boundaries of energy deposition by massive stars. It goes beyond ionization and excitation into the true "realm of the nebulae". It goes beyond the search of a threshold column density of matter required to form stars into the physics of gravitational collapse. On the whole, it provides us with a definite insight into the physics of cluster formation.

2. The quasi-steady gravitational collapse Perhaps a simple but valid approach to the problem is to follow the collapse of the original cloud as if it happened through a sequence of quasi-steady stages, all of which could be described by a simple analytical formulation. In such a case, one can define a cloud of radius r c and mass Mc whose gravitational potential (U) is given by

The cloud will have an internal pressure (Pc) given by 217/2 p

~

3 ~

A fragment within the cloud with temperature Tj and primarily thermal support has Pc as an external pressure and a central pressure: nfkTf

^Pc+

GMjfL,

(2 .3)

where nj = pj /m. A fragment is gravitationally unstable only if rj

2mG

K

'

or if Pc > \nfkTf,

(2.5)

and from these the well known result follows that self-gravitating fragments must have column densities comparable to those of the cloud as a whole: r}pj ~ rcPc.

(2.6)

The behaviour of a cloud during its quasi-static collapse includes processes involving the release of magnetic flux and angular momentum, generation of supersonic motions, dissipation of energy by those motions, shrinking of the cloud as a whole, and fragmentation into smaller gravitationally-bound objects. But, on the whole, the quasi-static collapse is supposed to fulfil at all times the virial theorem, and thus the characteristic speed of fragments in the potential well of their parent cloud will be

Tenorio-Tagle: Violent Star Formation

53

This, should be lost upon encountering a similar amount of matter, and thus, given the comparable cloud and fragment column densities, it should be lost in a cloud-crossing time (txing = 2 r c / v / ) . To sense the unavoidable fate of the matter involved, one should realise that as the quasi-static collapse proceeds,

the cloud would shrink further as the potential well deepens and new fragments should move faster despite their losses. The column density would increase and smaller fragments should still form with the same column

density.

Following the above equations (see paper I), one could show that the fragment radii and masses will follow

and

and that, thus, rj, and Mj decrease rapidly as the collapse proceeds and r c gets smaller. 2.1. The day the collapse stopped

One could then in principle continue to cut smaller and smaller fragments out of the cloud almost forever, or for as long as the collapse continues. However there comes a moment, the magic moment, when this is no longer possible, a moment when the collapse of the cloud as a whole is stabilized. We have inferred this from the observations, because we know that the collapsing matter did not disappear into a massive black hole at the centre of the configurations, but that instead thousands of stars formed and clusters present today a core and a tidal radius. Two conditions are needed to halt, or stabilize, the collapse. The first is that fragments should eventually acquire stellar sizes and turn into stars. At that moment, given their negligible cross-section, fragments can fully detach from the collapsing flow and thus are able to continue to move for the first time undecelerated through the left-over gas. The second condition is that the newly-formed suns undergo stellar winds. This latter event drastically modifies the flow as the left-over gas would feel for the first time an extra pressure, that of the winds ramming through it, and thus instead of causing the deceleration of the fragments, now, the newly formed stars sweep, condense and accelerate the gas up to the speeds acquired from the gravitational potential.

54

Tenorio-Tagle: Violent Star Formation

Given the above result for fragment masses, star formation at mass Mf = Af» becomes possible only when the cloud reaches a radius 2mG .,, c

, , .1 /•)

[M Mc)

'* ~ ZkTj * ' at which point both it and the proto-stellar fragments have column densities o

/ o i oi

\

(2J0)

2

The cloud pressure is Pc ~ ^{Pcrcf

(2.12)

and the characteristic velocity of the forming stars in the gravitational well is

Given that the cloud has stabilized, stars will form over a relatively long period of time, during which the characteristic value of v* remains as above. During this same period, fragments within the stabilized cloud will have similar velocities. A consequence of this is that one can identify the stellar and cloud gas dispersion velocities with v*, which, from the above relationships, implies either a critical pressure, or a critical column density, and/or a critical temperature as the cause of the observed
1/2

This implies a value of A ~ 1.1(T//10/<)(M»/MQ) 1 ^ 2 , while the observed correlation leads to A ~1.7. 2.2. The stabilizing agent The stabilization of the cloud collapse implies that the dispersion velocity of the gas and of the stars forming from it will be characterized by the depth of the gravitational potential, a ~ ( G M c / r c ) 1 / 2 . The gas is compressed and accelerated by the collection of bow shocks caused by the (supersonic) motion of multiple stellar wind sources. TTauri winds have already been propossed as a mechanism to account for the energetics, dynamical structure and lifetime of dark molecular clouds (Norman & Silk 1980; Franco 1984). In this view, the interaction of shells of swept up matter leads to a hierarchical clumpy structure with a mass spectrum that defines the evolution of the cloud and out of which an IMF could be inferred for long-lived clouds. However, the supersonic velocity widths cannot be accounted for with the interaction and collision of the wind-driven shells, as dissipative processes would soon erase any trace of the potential well from the system. Thus, newly-formed stars in such scenarios, would not have the velocity

Tenorio-Tagle: Violent Star Formation

55

dispersion required to stabilized the collapse of the cloud. The bow shock hypothesis, or cometary stirring model, could, on the other hand adequately maintain stability if the ensemble of bow shocks reached all parts of the cloud in a cloud-crossing time. If bow shocks have a cross section KR\0W and move with a velocity v*, then TV such shocks will encounter a volume

in one crossing time. In paper I, the TV sources required to maintain the cloud from collapsing any further were estimated and shown to lead to a cluster bolometric luminosity:

The result is that the cluster luminosity is approximately co~A /G, independently of the stellar wind velocity or luminosity, so long as the stars with radiatively driven winds dominate both the bow shock ensemble and the luminosity. The numerical result is = ^ = 0-5 x 10 8 L o ( ° _ t) . (2.17) 9G \ 1 0 km s / In conjunction with our previous result that a ~Arcll2 and equation (14) lead to Lcluster

and, as shown in paper I, this result is in excellent agreement with the observations.

3. The structure of star-forming clouds The cloud filling factor in star-forming regions is one of the long standing and largely ignored puzzles in Astronomy. It is a quantity originally thought to bring into an agreement the root mean square (rms) densities as measured in radio continuum with those determined at optical wavelengths, such as from Oil and/or SII lines. It is supposed to provide us with an idea about the structure of clouds, or at least indicate their degree of dumpiness. The factor of 10 difference between the two densities, typical in giant extragalactic HII regions (Kennicutt 1984), and the even larger values measured in galactic photoionized nebulae as in Orion (Osterbrock & Flather 1959), remain however unexplained. The controversy results, once again, from the well-known fact that photoionized condensations are supposed to be dispersed, or diluted, into the background, in a time scale which is much smaller than the lifetime of the ionizing stars (t* ~ 5 x 106 yr), or the evolutionary time of an HII region. The question is therefore how to explain the measured HII region filling factors, or the higher densities detected at optical frequencies, or otherwise, what replenishes the high-density condensations within the ionized gas throughout the history of the nebula. 3.1. The filling factor of star-forming clouds. The filling factor has been defined as the ratio of the square of the constant density < ne >r m » over the n2e measured at optical wavelengths, and it leads to a narrow range of values, of the order of 10~2 in giant HII regions (Kennicutt 1984). This expression is derived from the Stromgren relationship, which equates the number of ionizing photons

56

Tenorio-Tagle: Violent Star Formation

(Ft) to the number of recombinations in a constant-density (< ne > r m s ) ionized volume; i.e.

F* = 1*^/3 2rms,

(3.19)

which, when compared to the number (N) of cloudlets evenly spaced within the same ionized volume and capable of absorbing the same ionizing photon flux if their higher density is ne, becomes

l

(3.20)

and leads to the volume filling factor (/„) jvr-,3

.-

" 'cloudlet r

_ f

- h -

3

v. 2

_ ^ " e -^rms

HII

—2

/ Q <m

•

{6.ZL)

"e

If the stabilization of star-forming regions against further collapse is derived from the bow shocks and wakes caused by the winds from moving stars, the supersonic passage of the cometary sources must at the same time cause a sudden density enhancement of the swept up medium; i.e. right at the head of the bow shocks the density, if the shocks are isothermal, should reach a compression factor (3.22)

where Mjnach is the Mach number (the shock velocity over the sound speed of the ambient gas). The value of the density will be lower as one measures it away from the symmetry axis, as there only the velocity component perpendicular to the shock will be thermalized while the parallel component will be unaffected by the bow shock. However, even there, away from the bow shock symmetry axis, it will still be larger that the background density nnnThe compression across the cometary shocks would lead to wakes, filaments and/or channels and holes in the cloud, with a density maximum (ne) right behind the heads of the cometary shocks. As soon as this is produced, it would notice its overpressure with respect to the background medium and thus it would begin to expand. Its complete dispersal would takes place along the wakes caused by every comet, which would become wider and less dense as matter acquires values of density similar to those of the background gas. However, at all times the high-density gas would be replenished at the head of the wind-driven sources. If one identifies the background density with the < ne > r m , , then clearly one can also express the volume filling factor as fv

=

774 M

( 3 - 23 )

Mach

and thus conclude that the degree of dumpiness in star-forming regions is directly related to the supersonic passage of stellar wind sources ramming agaist the left-over cloud, and therefore the filling factor, or the structure of star-forming clouds, is directly related to the mass of the newly-formed cluster. A detailed comparison with the observations can be made (see Figure 1) if one sets limits to the temperature of giant HII regions. This, depending on metallicity, usually lies within the range 5000 K < T < 15000 K. Representative examples are the giant HII regions of M51 (Diaz et al. 1991) and low-metallicity HII galaxies (Terlevich et al.

Tenorio-Tagle: Violent Star Formation

57

FIGURE 1. The volumefillingfactor of star-forming clouds. The two solid lines encompass the narrow range of predicted filling factor values derived assuming THII — 5000 and 15000 K. Several well-known giant HII regions are placed (suqares) according to their measured velocity dispersion and their /„ = < ne >j mj /n,, while the filled circles indicate their predicted value according to the Tkn value given in the literature

1991), at opposite ends of this range of values. The temperature range implies sound speeds, (kT(\ + x)/m)ll2, in the range 9 to 15.7 km s" 1 , while the measured gas velocity dispersion for giant HII regions spans from 15 to 40 km s" 1 (see Table 1 of Hippelein 1986 or Arsenault & Roy 1988). Figure 1 shows the corresponding narrow range of predicted filling factor values (/„ = 1/Mj^,,,.,,) across the a range of applicability of these sources. Also shown are some of the classical giant HII regions for which accurate determinations of temperature, and thus sound speed, as well as gas velocity dispersion are available, together with filling factor values derived from the ratio < ne >r m «/ n eNote that prior to the birth of massive stars the compression factor and the dispersal time-scale of the shocked gas are both large, due to the low sound speed (CHQ ~0.3 km s"1) in the cloud. However, the birth of the ionizing sources marks the ensuing evolution (see Figure 2) and determines the structure of the parental cloud depending only on the total mass of the newly-formed stellar cluster. Clouds ionized by a low-mass cluster (Mciuater < 105 M©), with a correspondingly small c (< CHII) value, will cause the cometary sources to become subsonic upon photoionization. In this case the cloud structure ceases to be

58

Tenorio-Tagle: Violent Star

Formation

restored and thus that revealed to an observer is a consequence of the earlier evolution. In other words, photoionization would rapidly wash away, through local champagne flows, any sign of the cloud structure within the ionized volume. However, the original cloud matter (stirred prior to massive star formation) will continuously become apparent as the ionization fronts proceed into the cloud, showing its original structure (see Figure 2). T h e situation is drastically different in the case of massive stellar clusters, for which cometary sources remain supersonic (a* > CHII), even after complete photoionization of the parental cloud. In this case, the earlier cloud structure is rapidly dispersed, but the cometary sources will continuously replenish a higher-density medium causing, until complete cloud disruption, a characteristic cloud filling factor. This happens despite the continuous rapid disruption, through well-localized champagne flows, that restore the cloud low-density medium as the condensations are washed away to approach pressure equilibrium. The unavoidable fate of violent star-forming regions is cloud disruption. It is the action of massive stars through photoionization, stellar winds and their final supernova explosion that end up dispersing the parent clouds. Clearly, during the evolution of the HII region all sections of the cloud affected by the strong energy deposition will stop showing the effects produced by the presence of a stellar cluster. The effective clearance of matter will erase the supersonic velocity dispersion as well as the cloud structure. Therefore, as a function of time, only the best shielded and/or densest sectors of the parental cloud (which one should realise are not necessarily the central parts of the cloud), will be able to show through true Gaussian line profiles their supersonic kinematic nature, the signature of massive cluster formation. A good example of this is NGC 604 whose central and densest condensations (see the contribution by Munoz Tuiion in this volume and Sabalisck et a/. 1994), although separated by about 50 pc, present almost identical Gaussian profiles with the same value of a. Regions between the densest knots show strongly asymmetric line profiles, a sure signature of the many stellar winds presently disrupting the cloud. These lines can be fitted with broader Gaussians, however, of a much lower intensity than the true Gaussian profiles emanating from the densest knots. Therefore, for as long as this continues to be the case, for as long as some sections of the parent cloud survive the action of the massive stars, single aperture observations including the densest region of NGC 604 will show the supersonic crgas signature of the cometary stirring in what is still left from the parent cloud. 4. Conclusions In the cometary stirring model the cloud agitation is caused by the continuous supersonic passage of isothermal bow shocks, or "cometary" shocks, generated by the stellar wind sources ramming through the left-over cloud. These maintain the cloud stable against further collapse while causing supersonic turbulence and at the same time generating a distinct structure of the remaining cloud. All this becomes possible in the initial phases of stellar formation once the fragments resultant from cloud collapse become small enough, acquire stellar dimensions and form stars. This event detaches them from the collapsingflow,while their wind activity leads to the formation of cometary shocks. Such sources, given in sufficient numbers as to overrun the left-over cloud in a crossing time, cause a cluster luminosity shown to be correlated with cr4 and furthermore lead to a cluster size correlated with 105MG) the still supersonic passage of the numerous wind-driven sources will continuously restore the gas

Tenorio-Tagle: Violent Star Formation

59

Halo

^

FIGURE 2. Schematic view of star forming clouds. On the left, and as a function of time, is the sequence promoted by low mass clusters (MC|U,«er < 105AfQ ) able to cause the structure or filling factor of their parent clouds only before the birth of massive stars. Upon photoionization the cometary sources become subsonic and stop generating the cloud filling factor. Therefore, the structure observed in these clouds is that revealed to us as the ionization fronts progress into the leftover HI or H2 gas. On the right-hand side are the equivalent stages for massive cluster cases. Here despite photoionization the supersonic passage of the wind driven sources continuously replenish a cloud structure.

60

Tenorio-Tagle: Violent Star Formation

velocity dispersion and the cloud filling factor. It is then that the correlations between size and luminosity vs. a could be recognized. The input of mechanical energy from massive stars, however, will soon begin to upset the balanced situation by disrupting larger and larger sections of the parent cloud, erasing the clues stored in the gas phase during cluster formation. In this way, the observed gas turbulent motions are related to the total gravitational energy of the star-forming system, and in completed stellar systems the relationships are a relic of the former gas phase correlation. I would like to acknowledge many pleasent discussions on these topics with Casiana Muiioz Tunon, and Guido Munch. I also would like to thank the DGICYT (grant PB910531GEFE), the EEC (grant CI1*-CT91-O935) and NATO (grant CRG920198 for collaborative research) for partial finantial support to carry out this project. REFERENCES ARSENAULT, R. & ROY, J.-R. 1988 Astron. Astrophys. 201, 199. DIAZ, A. I., TERLEVICH, E., VILCHEZ, J. M., PAGEL, B. E. J. & EDMUNDS, M. G. 1991

M.

N. R. A. S. 253, 245. DYSON, J. E. 1979 Astron. Astrophys. 73, 132. FRANCO, J. 1984 Astron. Astrophys. 137, 85. HlPPELElN, H. H. 1986 Astron. Astrophys. 160, 374. HUNTER, D. A. & GALLAGHER, J. S. 1985 Astron. J. 90, 1457. KENNICUTT, R. C. JR. 1984 Astrophys. J. 287, 116. KENNICUTT, R. C. JR. & CHU, Y.-H. 1994 Astrophys. J. Submitted. MUNOZ TUNON, C. 1994 This volume. NORMAN, C. & SILK, J. 1980 Astrophys. J. 238, 158. OSTERBROCK, D. E. & FLATHER, E. 1959 Astrophys. J. 129, 26. ROSA, M. & SOLF, J. 1984 Astron. Astrophys. 130, 29. SABALISCK, N. S. P.,

Astrophys.

TENORIO-TAGLE, G., CASTANEDA, H. O., MUNOZ TUNON, C.

1994

J. In press.

SKILLMAN, E. D. &: BALICK, B. 1984 Astrophys.

J. 280, 580.

TERLEVICH, R. J., MELNICK, J., MASEGOSA, J., MOLES, M. &; C O P E T T I , M. V. F . 1991

Astron. Astrophys. Suppl. 91, 285. TERLEVICH, R. & MELNICK, J. 1981 M. N. R. A. S. 195, TENORIO-TAGLE,

G.,

MUNOZ TUNON,

839.

C. & Cox, D. P. 1993 Astrophys. J. 418, 767.

Super-Associations as Star Complexes with Violent Star Formation By YURI N. EFREMOV Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899, Moscow, Russia

1. Super-associations and star complexes The concept of a "superassociation" was first introduced by Baade (1963) in his Harvard lectures in 1958. He gave this name to a region about 500 pc across around the giant HII region 30 Dor in the LMC which is full of OB-associations. The same region was the first example of a super-association given by Ambartsumian (1964). Altogether, 19 OB-associations and young clusters here form a morphological unit of 1 kpc in diameter with evident hierarchical structure (Efremov 1988, 1989). One may say that a "super-association" is the counterpart of a hydrogen emission nebula (HII region) in B, V etc. broad bands. Wray and de Vaucouleur (1980) have shown that in the B bandpass the continuum-to-emission ratio is always greater than 10:1. Thus in this band one deals mainly with the star population of a super-association. Nevertheless, the diameter of the 30 Dor HII region is only 250 pc and it occupies less than 0.1 of the total area of the super-association, the remaining HII regions here being much smaller. This may well also be the case for extragalactic super-associations - giant HII regions. In bright star cloud NGC 206 in M31, named by Baade (1963) as a real super-association, there are only a few small HII regions, not seen at all on the B plates. In many respects super-associations (SAs) are similar to common star complexes (Efremov 1978, 1993), the main difference being the richness of an SA in HII gas and OB stars that causes the high total luminosity. Whether there exists a clear borderline between these complexes and SA, either in size, luminosity or mass, remains still to be investigated, though it is likely that the average SA is larger than a common complex. As yet, the luminosity function of star complexes is still unknown. Data for the total sample of star complexes in a number of galaxies are necessary in order to attain this goal . Those objects with lower luminosity have usually not been considered at all and only those isolated patches in galaxies were searched and named as SA whose luminosity is above B = —14 (Ambartsumian 1964; Petrosian, Saakyan & Khachikian 1984). In our opinion SAs have the same nature in principle as common star complexes but with an unusually high rate of (massive) star formation over the total area of the complex. Within a super-association there are a dozen or two OB associations and young massive clusters, whereas within a common complex of similar size the number can range from zero to three or four associations (Efremov 1988, 1989). In what follows we accept the point of view that complexes are formed from giant superclouds forming in a gaseous-stellar galactic disk owing mainly to gravitational instability (Elmegreen k Elmegreen 1983; Elmegreen 1987; Larson 1988; Efremov 1988, 1989, 1993; Elmegreen et al. 1994). According to Kennicutt (1989) the threshold gas density for massive star formation corresponds to just the Jeans mass and wavelength in the gaseous disk, which are in fact observable parameters of superclouds (Rand 1993). 61

62

Efremov: Star Complexes and Super-Associations

2. Reasons for violent star formation in super-associations 2.1. Increased velocity dispersion

There must be special conditions for violent star formation to take place over a whole complex. We have in total no more than two or three SAs within the Local Group (OB 78 = NGC 206 in M31, 30 Dor in LMC and NGC 604 in M33), while there are about 270 star complexes in these three galaxies. This ratio means that SAs are not simply young complexes, because the mean ages of SAs and complexes relate as 1 to 10. (One should note that only within the Local Group of galaxies may complete samples of complexes, including those without active star formation, be picked out, and that even for these galaxies the distributions of complexes are known in size only, not in luminosity.) We suggested that one probable condition for violent star formation in a complex is high-velocity dispersion in a parent supercloud. According to Elmegreen (1992), and Elmegreen, Kaufman k. Thomasson (1993), a cloud with higher internal dispersion should produce stars with greater efficiency because of the cloud's greater resistance to selfdestruction under energy input from new-born O stars and SNe. Such clouds should also produce a larger fraction of massive stars owing to more heating of star-bearing gas because of higher star formation efficiency. The former authors gave a number of examples of superclouds with violent star formation in interacting galaxies. Some disturbance of the cloud- bearing interstellar medium should exist that leads to higher velocity dispersion in superclouds resulting from the ISM (Elmegreen 1992). These superclouds are larger than those usually found in a normal undisturbed galactic disk. Super-associations are common within interacting galaxies but in normal galaxies they usually located at some special position in a spiral arm, such as at the beginning of a spur (as in OB 78 in M31) or at the end of an arm (Petrosyan et al. 1984). Within irregular Magellanic-type galaxies SAs are often located at the end of a bar (30 Dor in the LMC, NGC 2363 in NGC 2366 etc). In such locations some kinematical disturbance of the interstellar medium and higher velocity dispersion in superclouds is also possible. Indeed the star velocity dispersion within the 30 Dor complex is 14 km s" 1 , whereas in other concentrations of supergiants in the LMC it is only 6 km s" 1 . This large dispersion in 30 Dor is explained by three overlapping velocity fields (Martin et al. 1984). There are some data on the high dispersion of HII velocities inside super-associations, though whether it is connected with turbulence or gravitational motions is still a controversial issue. This is reviewed, for example, in Efremov (1989). The arguments of Terlevich & Melnick (1981) and Melnick et al. (1987) in favour of the virial nature of these motions seem attractive and we consider their data as one more piece of evidence that star complexes are born as physically connected entities from a self-gravitating supercloud and are not random agglomerations of clusters, stars and clouds. 2.2. Density and pressure enhancement

It is quite possible that other initial conditions might lead also to more or less violent star formation. Density and temperature within the parent superclouds of SAs are probably similar to those in a spiral density wave and a preference for high-mass star formation was argued for by many authors (e.g. Mezger 1987). A drastic increase in the gas density, owing to cloud collisions or to external pressure (such as observed inside galactic spiral density waves) may lead to a burst of star formation (Larson 1987), and in a galaxy with a strong spiral wave (such as M51, M83 and NGC 6946) some star complexes situated along arms may be well called SAs. It is quite probable that in every situation the transition between SAs and common complexes is a continuous one. The best examples of intermediate entities are provided by the Andromeda galaxy (Battinelli et al. 1993).

Efremov: Star Complexes and Super-Associations 2.3.

63

Triggering by shock-shock collision

Shock-shock collision within a supercloud leads to enhanced reflected shock waves that propagate with high velocity and trigger additional star formation (Chernin k Efremov 1993). There must exist within a parent supercloud at least two simultaneously acting star formation centres producing shock waves in the interstellar medium to produce violent star formation under the action of this mechanism. If there are a number of such centres, the resulting pattern of age gradients in a superassociation must be very complicated and much work still needs to be done before a reliable comparison with observations is possible. Nevertheless, within small parts of the 30 Dor super-association, age progression is observed from the star-forming molecular cloud near the HII region NGC 2074 to the older dispersed association NGC 2081 (Lortet & Testor 1988). Farther away the still older blue globular cluster NGC 2100 is located in the same line. Probably there are some signs of an age gradient in a certain direction, compatible with Chernin k. Efremov's (1993) expectations. This mechanism naturally explains the large velocity dispersion but probably not the preferred positions of super-associations. 2.4. Abundance peculiarities inside superassociations? Another important issue is the possible influence of abundance on the star formation rates, efficiencies and especially the initial mass functions inside different complexes. There exists a strong concentration of WN to young star complexes both in our Galaxy (Alfaro, Delgado k Cabrera-Cafio 1992) and in the LMC (to the 30 Dor complex, Moffat et al. 1987), whereas stars of WC subtype do not show such behavior. It is improbable that WC stars are simply older because there are indications that they are more massive than WN stars (Vazquez k Feinstein 1990). The ratio of WN/WC stars is larger in galaxies and in regions of a galaxy with lower metallicity (Maeder 1987; Arnault et al. 1989) and one suspects that brighter (i.e. with higher SFR) complexes may have lower metallicity compared with the bulk of star complexes. This difference in abundance may have something to do with the initial luminosity function and lead to higher resultant star formation efficiency as well, for lower-metallicity protoclusters should be more successful in resisting the disruptive forces from O-stars and supernovae (Elmegreen 1983). Here the same arguments act as in the case of large velocity dispersion. Lower metallicity might also lead (via influence on mass loss or even initial mas function) to a preponderance of WN stars in complexes with more active star formation. Then the more frequent occurrence of SAs at the periphery of a spiral galaxy may be connected with the lower abundance there. By the way, the blue (not the red or WR star) progenitor of SN1987A in the LMC is also more compatible with lower abundance in the LMC and probably especially in the 30 Dor super-association. Moreover, the star is suggested to be a member of a small cluster and it is quite possible that stars in clusters generally have lower abundances than surrounding field stars just because the latter were born in already dissolved - owing to larger metallicity? - associations (Efremov 1991). There remains of course the difficult problem of why some complexes of 1-kpc scale have lower abundances than that of the general field - or those of other complexes. One may guess that the mixing of the interstellar medium within a galaxy disk is not so effective as one usually thinks, and that some relics of primordial or, more plausibly, subsequent merging are alive now. There may be collisions of extragalactic clouds with the gas of the galactic disk leading to star formation accompanied by the deviation of the galactic disk from the mean plane (Alfaro et al. 1991). These clouds should have generally lower abundances than galactic disk gas, because there has been little or no preceding star formation there.

64

Efremov: Star Complexes and Super-Associations

Surely these considerations do not relate to starbursts in galactic centres where the most interesting phenomena probably connected with violent star formation in high metallicity medium have been observed (Terlevich 1993). Unfortunately we have too many explanations for violent star formation phenomena. Is any one of these correct?

REFERENCES ALFARO, E. J., CABRERA-CANO, J. & DELGADO, A. J. 1991 Ap. J. 378, 106. ALFARO, E. J., DELGADO, A. J. & CABRERA-CANO, J. 1992 Ap. J. Lett. 386, L47.

AMBARTSUMIAN, V. A. 1964 In IAU Symp. No. 20: Galaxy and Magellanic Clouds, p. 122.

Canberra, Austral. Acad. Sci.. ARNAULT, P H . H., KUNTH, D. & SCHILD, H. 1989 A. & A. 224, 73

W. 1963 Evolution of Stars and Galaxies. Harvard Univ. Press. P., EFREMOV, YU. N. & MAGNIER, E., 1993 This conference. CHERNIN, A. D. &; EFREMOV YU. N. 1993 This conference. BAADE,

BATTINELLI,

EFREMOV, YU.N. 1978 Sov. Astr. Lett. 4, 66. EFREMOV, YU.

N. 1988 Soviet. Sci. Rev. E: Astrophys. Space Phys. Rev. 7(2), 105. Harwood,

London. N. 1989 Ochagi zvezdoobrazovania v galactikakh (Origins of Stars Formation in Galaxies) Nauka, Moscow.

EFREMOV, YU.

EFREMOV, YU. N. 1991 Pisma Astr. Zh. 17, 404.

EFREMOV, YU. N. 1993 Star Formation, Galaxies and Interstellar Medium (ed. J. Franco et al.), p. 360. Cambridge Univ. Press. ELMEGREEN, B. G. 1983 M. N. R. A. S. 203, 1011. ELMEGREEN, B. G. 1987 Ap. J. 312, 626.

ELMEGREEN, B. G. 1992 In Star formation in Stellar systems, III Canarian Winter School, (ed. G. Tenorio-Tagle et al.),p 381. Cambridge Univ. Press. ELMEGREEN, B. G. & ELMEGREEN D. M. 1983 M. N. R. A. S. 203, 31. ELMEGREEN, B. G., KAUFMAN, M. & THOMASSON, M. 1993 Ap. J. 412, 90.

ELMEGREEN, D. M. ET AL. 1994 Ap. J. 425, in press. KENNICUTT, R. C. 1989 Ap. J. 344, 685.

LARSON, R. B. 1987 In Starbursts and galaxy evolution (ed. Trinh Xuan Thuan et al.), p. 467. R.B 1988 In Galactic and extragalactic star formation (ed. R.Pudritz and M.Fich), p. 459.

LARSON,

LORTET, M.-C. & TESTOR, G. 1988 A. & A. 194, 11.

MAEDER, A. 1987 In Starburstst and Galaxy Evolution (ed. Trinh Xuan Thuan et al.), p. 107. MARTIN, N. ET AL. 1984 In IAU Symp. 108: Magellanic Clouds, p. 137. MELNICK, J., MOLES, M., TERLEVICH, R. & GARCIA-PELAYO, J.-M. 1987 M. N. R. A. S. 226,

849. MEZGER, P.C. 1987 Publ. Astr. Just. Czech. Acad. Sci. 69, 91. MOFFAT, A. F. J., NIEMELA, V. S., PHILLIPS, W., CHU, Y.-H.& SEGGEWISS, W. 1987 Ap. J. 312, 612. PETROSIAN, A. R., SAAKVAN, K. A. & KHACHIKIAN, E. E. 1984 Astrofisika. 21, 57. RAND, R. J. 1993 Ap. J. 404, 593. TERLEVICH,

R. 1993 This conference.

TERLEVICH, R. & MELNICK, J. 1981 M. N. R. A. S. 195, 839.

VAZQUEZ, R. A. & FEINSTEIN, A. 1990 Rev. Mexicana Asti. Astr. 21, 346. WRAY, J. D. & DE VAUCOULEURS, J. 1980 A. J. 85, 1.

Violent Star Formation Driven by Shock-Shock Collisions By ARTHUR D. CHERNIN AND YURY N. EFREMOV Sternberg Astronomical Institute, Moscow University, Moscow 119899, Russia Regions of violent star formation such as supergiant HII regions or superassociations often reveal a binary space structure: they contain two separate components within which very intensive star formation proceeds more or less simultaneously. This observational fact suggests an evolutionary scenario for the phenomenon, in which the key role is played by a strong collision of shock fronts produced by the energy release of the previous generation of massive stars in the region.

1. Binary stuctures and nonlinear gas dynamics There are well observed giant regions of intensive star formation - superassociations that consist of two (or three) components: Per-Cas, Sco-Cen and Car in the Milky Way, OB 78 in M31, Region IV = 30 Dor E + 30 Dor W in the LMC (Efremov 1988, 1989). A dust lane is observed in some cases between the two parts of the region that makes this composite structure especially obvious. In regions like OB 21 in M31, HII clouds give an even more contrasting picture when they concentrate at the two opposite sides of the dust lane. Can this binary spatial structure be a clue to the physical nature of the violent star formation phenomenon? We assume that the answer to this question is positive and present it here in the form of an evolutionary scenario in which shock-shock collisions in the interstellar gas play a key role. In a simple version, the scenario may consist of three basic stages: 1. Two "ordinary" star formation bursts occur in two separate but neighbouring areas of interstellar matter and produce two expanding shocks; 2. The shocks collide and two reflected shocks appear moving in the opposite directions; 3. Extremely intensive formation of massive stars proceeds in the very dense gas (shocked twice!) behind the reflected fronts due to gravitational and hydrodynamic instabilities.

2. The physics of shock-shock collision The initial stage of this picture is similar to triggered star formation as suggested by Elmegreen and Lada 1977 (see also recent reviews by Tenorio-Tagle & Bodenheimer 1988; Elmegreen 1992). The two other stages are characteristic of the shock-shock scenario. The shocks from the neighbouring star formation events can come into effective interaction, if the distance between the centres of star formation is not too large and the time difference in the output of energy from the stars is not too great. This unity of place and time leads to the unity of action. Our study of the formation and evolution of the shock-shock interaction is based on the classic work by Courant & Fridrichs (1948) and also on a set of computer simulations of the process performed at Ioffe Physical-Technical Institute (Barausov, Voinovich & Chernin 1988, 1992; Chernin 1993). 65

66

Chernin k. Efremov: Violent Star Formation Driven by Shock-Shock Collisions

According to these results, the shock-shock interaction gives rise to the formation of a regular dynamical pattern in the zone of the contact collision of the fronts. This structure includes, besides the reflected shocks that are mentioned already, a Mach ring-shaped front and two contact jumps (for the geometry of the pattern see the figures in the references above). The dynamics of the reflected shock which is of special interest here is critically affected by the character of the matter distribution in the area where the shock propagates. This density distribution is prepared by the initial shocks that compress the major part of the gas into two spherical shells behind their fronts. For the reflected shock, it is the medium of decreasing density along the direction of the shock propagation. The shock in the medium of decreasing density is known to move with an increasing velocity (Chisnell 1955; Whitham 1958). In addition, the accelerating shock in this gas distribution proves to be self-collimating (Chernin 1994): the central area of the front moves along the direction of the steepest density decrease, and so its velocity increases in the most rapid way. As a result, the surface of the front becomes more and more convex to the direction of its propagation. As a result, the gas flow behind the front gets stream-like or even jet-like with time.

3. Back to astronomy The reflected fronts drive the gas back to the areas of the initial star formation bursts. This opens a way to two new, more or less simultaneous, and - most importantly - much larger bursts of star formation in the same two areas. In the regions of violent star formation that might form and evolve in this manner, we can expect: A) Two populations of stars in each component of the binary star forming region with an age difference of about 10-30 Myr that is related to the time interval between the initial star formation event and the secondary one behind the reflected shocks; B) Different space distribution and kinematics of the stars of the two populations, reflecting the initial condition in the gas from which they formed in each of the two components: more or less isotropic distribution of the first generation stars and a streamlike distribution of the second generation stars; C) The zone of the shock-shock interaction bounded by the Mach ring and the surfaces of the tangent jumps accumulates dust particles of the face ridge of the initial shocks; there are no outflows here along the line connected the centres of the two initial starbursts, but there are outflows in the directions perpendicular to this line due to the expansion of the Mach ring front; because of this the dust forms a shell here which is observed in the projection as a strip between the two components of the region. Features that are fairly similar to A and C are definitely or probably observed in some supergiant HII regions and superassociations (see Efremov 1982, 1988, 1989, 1993). Features like B may provide an observational test for the mechanism of violent star formation assumed in the scenario. The physics of shock-shock collisions may suggest, perhaps, a basis for a better understanding of other observed violent processes associated with star formation. One example is small-scale supersonic bipolar structures (a few parsecs in size) around young stellar objects. Highly collimated and accelerating jets observed in these objects may be due to reflected shocks produced by shock-shock collisions (Chernin 1994). Another example may be related to the larger space scales. As was first argued by Shklovsky (1984), many active galactic nuclei are most probably the regions of violent star formation (see also Terlevich & Melnick, 1985). It is instructive from the point of view

Chernin & Efremov: Violent Star Formation Driven by Shock-Shock Collisions 67 of the present discussion (for more details see Efremov & Chernin 1994) to consider that active galactic nuclei often or even always display a binary (multiple) spatial structure when the resolution is high enough in the observations (e.g. Heap et al. 1993; Vacca et al. 1993).

REFERENCES D. I., VOINOVICH, P. A. & CHERNIN, A. D. 1988 Preprint No. 1274, Ioffe PhysicalTechnical Institute , Leningrad. BARAUSOV, D. I., VOINOVICH, P. A. & CHERNIN, A. D. 1992 Sov. Astron. Lett. 18(12). BARAUSOV,

CHERNIN, A. D. 1993 A. k A. 267, 315.

A. D. 1994 In press. R. F. 1955 Proc. Roy. Soc. A223, 250. COURANT, R. li FRIDRICHS, K. O. 1948 Supersonic Flow and Shock Waves. Interscience. EFREMOV, YU. N. 1982 Sov. Astron. Lett. 8, 663. EFREMOV, YU. N. 1988 Stellar complexes (Soviet. Sci. Rev. E: Astroph. Space Phys. Rev. 7(2), 105). Harwood. EFREMOV, YU. N. 1989 Ochagi zvezdoobrazovania v galactikakh (Origins of Star Formation in Galaxies). Nauka. EFREMOV, YU. N 1993 This conference. EFREMOV, YU. N. & CHERNIN, A. D. 1994 In press. ELMEGREEN, B. G. 1992 In Star Formation in Stellar Systems. HI Canarian Winter School (ed. G .Tenorio-Tagle et al.), p. 381. Cambridge University Press. CHERNIN,

CHISNELL,

ELMEGREEN, B. G. & LADA, C. J. 1977 Astrophys. J. 214, 725.

S. R. ET AL. 1993 This conference. SHKLOVSKY, I.S. 1984 Sov. Astron. Lett. 11, 163. TENORIO-TAGLE, G. & BODENHEIMER, P. 1988 Ann. .Rev. Astton. Astrophys. 26, 146. HEAP,

TERLEVICH, R. & MELNICK, J. 1985 M. N. R. A. S. 213, 841.

W. D., CONTI, P. S. & LEITHERER, C. 1993 This conference. WHITHAM, G. B. 1958 J. Fluid Mech. 4, 84. VACCA,

The Search for Hierarchical Structure inside M31 Superassociations ByPAOLO BATTINELLI 1 , YURI N. EFREMOV2 AND EUGENE A. MAGNIER3 1 2

Osservatorio Astronomico di Roma, Viale del Parco Mellini 84, 1-00136 Roma, Italy

Sternberg Astronomical Institute, Universitetskii Prospect 13, 119899, Moscow, Russia

3

Astronomical Institute "Anton Pannekoek" and Center for High Energy Astrophysics, Kruislaan 403, 1098 SJ Amsterdam, The Netherlands and Center for Space Research and Departament of Physics, Massachusetts Institute of Technology, USA

The identification and classification of young star groups in other galaxies is still a controversial topic (Battinelli 1991a). Very different estimates of OB association sizes in different galaxies were explained by Hodge (1986) by a difference in linear resolution and limiting magnitude, but the existence of two kinds of resolved star groups is also essential in this issue. The bulk of OB associations are members of larger groups, star complexes (Efremov 1978), of diameter 400-1000 pc that also include individual fainter and older stars, such as Cepheids. There exists a hierarchical, embedded sequence of young star groups: there are associations, aggregates, complexes and supercomplexes (regions). Associations and complexes are the more common ones (Efremov 1988, 1989, 1993). Given sufficient resolution and limiting magnitude, one can see both complexes as well as brighter and smaller associations mainly inside complexes, and such is the case for M31, where Efremov, Ivanov & Nikolov (1987; hereafter EIN) found, by eye, OB associations of typically 80 pc diameter and star complexes of typically 600 pc. The latter complexes are mainly the same large groups of blue stars that were identified by van den Bergh (1964) under the name of OB associations. The issue arises whether one has simply OB associations with a large range of sizes, the smaller ones being younger as suggested by van den Bergh 1964, or if there exist two kinds of star groups of different hierarchical level, with younger ones as constituent parts of larger and older ones, as suggested by Efremov (1978, 1989) and EIN. To tackle this matter, the application of an objective method of isolation of resolved star groups is of primary importance (Battinelli 1991a, 1992). This task was carried out recently by Magnier et al. (1993; hereafter MBL) for M31 stars. They applied a variant of a cluster analysis method (the Path Linkage Criterion elaborated by Battinelli 1991b) to a sample of about 7000 blue stars extracted from a set of about 300000 stars for which BVRI CCD photometry was obtained by Magnier et al. (1992). MBL found 174 associations with contamination of lower than 5% by random clumps, the total number of associations (within central 1 deg2) being 280 and average diameter 90 pc. Thus MBL confirm the value found by EIN for the size of M31 associations, very similar to the sizes of associations in the Milky Way and Magellanic Clouds. In MBL's list of M31 associations, there are also half a dozen larger groups such as NGC 206 (van den Bergh's OB 78) with large diameters. These groups, having larger densities of blue stars, are intermediate between common star complexes and superassociations. Within OB 78 (= OB 13 in MBL, with a size of 580 pc - a genuine superassociation), MBL identified 8 clumps in a second hierarchical level with diameters of about 100 pc. In the present contribution we apply the second level PLC to six other clumps with 68

Battinelli et al.: Superassociations in MSI

69

diameters larger than 250 pc from the MBL list. Internal structure was found in each of them, with mean diameters close to 100 pc and with the following numbers of "real" clumps (i.e. those with a probability under 10% of being random): OB 92 (5), OB 136 (2), OB 118 (3), OB 142 (1), OB 155 (3). One may conclude that hierarchical structure is a typical property of the larger clumps of blue stars in M31, which corroborates the suggestion that these large clumps are association complexes. The complexes isolated by MBL using the distribution of blue stars are evidently younger (or with higher star formation rate) than the bulk of complexes in M31, i. e. those older complexes which are not detected by MBL but seen by eye both by van den Bergh and EIN. Fainter stars and/or a larger scale length than that used by MBL should be used to search for star complexes in the Andromeda galaxy with the PLC. Thus, only a few of the brightest complexes are easily seen in the blue star distribution, and the case of M31 may well be typical. Only in the nearest galaxies is the total sample of complexes generally detectable. The implied ubiquity of star complexes which encompass the bulk of associations in disk galaxies (Efremov 1988, 1989, 1993) is in agreement with top-down scenario of star formation according to which the vast superclouds that produce star complexes are the initial structures (Elmegreen 1979, 1987; Elmegreen & Elmegreen 1983).

REFERENCES BATTINELLI, P. 1991a Mem. Soc.Astron. It. 62, 959.

P. 1991b Astron. Astrophys. 244, 69. BATTINELLI, P. 1992 Astron. Astrophys. 258, 269. BATTINELLI,

ELMEGREEN, B. G. 1979 Astrophys. J. 231, 372. ELMEGREEN, B. G. 1987 Astrophys. J. 312, 626. ELMEGREEN, B. G. & ELMEGREEN, D. M. 1983 Mon. Not. R. Astron. Soc. 203, 31.

N. 1978 Sov. Astron. Lett. 4, 66. N. 1988 Astrophys. Space Phys. Rev. 7(2), 105. EFREMOV, YU.N. 1989 Origins of Star Formation in Galaxies: Star complexes and Spiral Arms (in Russian). Nauka, Moscow. EFREMOV, YU. N. 1993 In Star Formation, Galaxies and the Interstellar Medium (ed. J. Franco, F. Ferrini & G.Tenorio-Tagle), p. 360. Cambridge Univ. Press, Cambridge. EFREMOV, YU. N., IVANOV, G. R. & NIKOLOV, N. S. 1987 Astrophys. Sp. Sci. 135, 119. HODGE, P. W. 1986 In Proc. IA U Symp. 132: Luminous stars and associations in galaxies (ed. C. H. W. de Loore et al.), p. 369. Reidel, Dordrecht. MAGNIER, E. A. ET AL. 1993 Astron. Astrophys. 278, 36. MAGNIER, E. A. ET AL. 1992 Astron. Astrophys. Suppl. 96, 379. VAN DEN BERGH, S. 1964 Astrophys. J. Suppl. 9, 65. EFREMOV, YU. EFREMOV, YU.

A Stochastic PSF Model: Smooth Spirals in Differentially Rotating Disks ByJAN PALOUS !AND BRUNO JUNGWIERT 2 'Astronomical Institute, Academy of Sciences, Bocni II 1401, 141 31 Prague 4, Czech Republic 2

Center for Theoretical Study, Charles University, Celetna 20, 110 00 Prague 1, Czech Republic

The propagating star formation model with anisotropic probability distribution is investigated. In each star-forming site we define the probability ellipse and show that its two parameters, the excentricity and the orientation relative to the galactic rotation, are closely related to the thickness and inclination of the resulting spiral arms. The relative size of a star-forming region with respect to the whole galaxy is also discussed. Simulations are compared to the observed galactic morphologies and we mimic the differences between the two groups of galaxies of types M101 and NGC 7217.

1. Propagating star formation The idea that star formation at one place in a galaxy can initiate star formation in its neighbourhood was first suggested by Opik (1953) and Oort (1954). Since then, a possible chain of physical processes that joins two regions of subsequent star formation has been proposed in which ionizing radiation from massive stars in a cluster leads to the disruption of the parental molecular cloud via supersonic champagne flows halting further star formation. The mechanical energy input from stellar winds and supernova explosions causes the agglomeration of gas in expanding shells. Their fragmentation, the building of molecules in high opacity areas, and large-scale gravitational instabilities may produce molecular clouds, where the next generation of star formation occurs. The star-forming cycle described above is the basis of deterministic PSF models (Palous et al. 1994). However, the physical parameters such as density, metallicity and cooling times of the ISM, are only partly known. At the same time, the detailed model of star formation inside the cores of a molecular cloud is not generally accepted. The role of magnetic fields is also poorly understood and the interaction of expanding shells with the ambient medium is still under discussion. The undetermined ingredients in the star-forming cycle can be approached with probabilistic arguments in in terms of stochastic self-propagating star formation (SSPSF) models, in which star formation propagates from place to place with a certain probability. Such a description was introduced by Gerola k Seiden (1978). The SSPSF is a percolation model: it describes large-scale galactic properties resulting from the temporal and spatial percolation of small-scale events (Seiden k Schulman 1990). The galaxy is globally parametrized with the finite probability Po, which decreases locally over the recovery time r after the last star formation in the region's history is taken into account. Apart from Po and r, other two parameters are involved: the velocity of rotation and the relative size of a star forming region with respect to the whole galaxy. The SFR depends on Po and the model shows the non-linear behaviour typical for a phase transition: for Po smaller than a critical value P c the SFR is zero. For Po immediately higher than P c the SFR rises steeply. Spiral shapes similar to those of observed galaxies are formed if Po is in a narrow interval near Pc. 70

J. Palous & B. Jungwiert: Stochastic Propagating Star Formation

71

2. SSPSF model with an elliptical probability distribution In differentially rotating galactic disks, two major deformations of expanding shells are predicted and observed (Palous et al. 1990): a. The originally round shell acquires an elliptical shape. The major axis of the shell becomes larger and rotates with respect to the galactic center direction from 45° to 90°. b. The mass collected by the shell continuously slips behind the overtaking shock, and accumulates near the tips, causing an uneven column-density profile. For SSPSF models, the uneven ISM distribution along the shell periphery suggests an anisotropic probability, which is highest along the shell's major axis. This idea was first proposed by Chiang & Elmegreen (1982). In our model (Jungwiert & Palous 1994) the anisotropy is introduced by a probability ellipse connected to any star forming site. The effective propagation probability in a certain direction is proportional to the length of the chord in the probability ellipse, which is defined by two parameters: 1. the axial ratio 6 : a, and 2. the inclination / of the major axis relative to the galactic center. Thus our SSPSF scheme involves the above two parameters plus the following ones, which were already present in the models of Seiden & Schulman (1990): 3. the propagation probability Po, 4. the rotation curve of the galaxy, 5. the time step At, 6. the recovery time r, and 7. the relative size of a star-forming region with respect to the whole galaxy. The relation of these quantities to the spiral features should be discussed and compared to the observed galactic morphologies. The role of 6 : a and I, as well as their connection to the ISM properties and to the rotation curves is analyzed by Jungwiert & Palous (1994) and is only briefly summarized here. In this paper the influence of the relative size of a star-forming region is also mentioned.

3. Results 3.1. The rotation curve, I and b : a The rotation curve determines the amount of differential rotation at a certain galactocentric distance, which affects the global shape of spiral arms. With an isotropic probability distribution the SSPSF produces hyperbolic spirals with the pitch angle depending on the rotation speed Vrot- Slow rotators create more open spirals than do fast rotators. The ellipticity of the probability distribution also generates spiral-like structures. In the absence of differential rotation it leads to logarithmic spirals, if / is kept constant throughout the galaxy. With both ellipticity and differential rotation included, the global spiral shapes are neither purely hyperbolic nor purely logarithmic. The effect of I is shown in Figure 1. With I increasing from 45° through 65° to 85" the winding up of the spiral arms increases: low / produces more open and shorter spirals than high /. The effect of b : a is demonstrated in Figure 2. The anisotropy of the probability distribution plays an important role when it is sufficiently large b : a < 1 : 2. Higher excentricities (b : a = 1 : 10 or 6 : a = 0, where the latter means that the ellipse degenerates to the abscissa) produce remarkably organized spirals, which are smoother than for isotropic SSPSF.

72

J. Palous k B. Jungwiert: Stochastic Propagating Star Formation

7 = 65°

7 = 45°

•

• • • • • • • • •

10

- 1 0

-10

10

10

7 = 85°

-10-10

10

1. A sequence of I = 45°, 65°, and 85°. Vtot = 150 k m s " 1 , and i : a = 1 : 10 in all three cases. Star clusters of ten different ages are shown, the symbol radius decreases linearly with age. Axes are labelled in kpc. FIGURE

10

-10 -10

b:a=l:l

10

io

10

10

b:a=0

-10

-10

10

FIGURE 2. A sequence of b : a = 1 : 1, 1 : 10, and 0. Vr<>t = 150 km s""1 in all three cases. / = 50° when 6 : a < 1.

3.2. The relative size of a star-forming region

In Figures 1 and 2 the size of a star-forming region is close to 200 pc and the radius of the galactic disk is 10 kpc. For comparison, we present in Figures 3 an 4 the same experiments for a star-forming region twice as small. The spiral features are thinner but more numerous. This was already discussed by Seiden and Schulman (1990) in relation to the difference between flocculent and grand-design spirals. They argue that galaxies with relatively small HII regions tend to have flocculent structure. We can confirm this point in the case of isotropic probability distribution (Figure 4, left). In simulations with anisotropy, the long and smooth spiral arms are present even when star-forming regions are small.

J. Palous & B. Jungwiert: Stochastic Propagating Star Formation

10

7 = 45°

-10 -10

10

10

7 = 65°

-10 -10

73

10

to -^10

10

FIGURE 3. The same as Figure 1, except for the relative size of a star forming region, which is twice as small.

10

-10 -10

FIGURE

b:a=l:l

10

10

-10

-10

b:a=l:10

10

••• ' • •

10

-10

-10

b:a=0

10

4. The same as Figure 2 except for the relative size of a star-forming region, which is twice as small.

4. MlOl versus NGC 7217 We try to compare the output from SSPSF models with anisotropic probability distribution with pictures of observed galaxies. To mimic robust, open, grand-design spirals of MlOl type, moderately low 6 : a (~ 1 : 10) and low 7 (~ 45°) are needed. A simulation with such a probability ellipse and with the rotation curve of MlOl is shown in Figure 5. Rather conspicuous, long and smooth arms (unattainable with the isotropic SSPSF) can be seen as well as more patchy areas (at T = 6 Gyr note the existence of two similar and symetrically located arms, extending over 270° in galactic longitude). Of course, the resemblance to MlOl can be only qualitative since the model is stochastic and the right choice of b : a and I is unknown. On the other hand, to reproduce the tightly wound and broken spiral arms of NGC 7217 type, extremely low 6 : a (~ 0) and high I (~ 75°) values are plausible. A more detailed discussion with figures of both galactic groups is given by Jungwiert and Palous (1994).

J. Palous & B. Jungwiert: Stochastic Propagating Star Formation

74

= 2Gyr

T = 4 Gyr

10

10

10

T = 6 Gt/r

. v. -10 -10

10

10

-10

10

FIGURE 5 . T i m e evolution of a galaxy w i t h 6 : o = l : 1 0 , 7 = 45° a n d t h e r o t a t i o n c u r v e of M101.

5. Conclusions We conclude from Figures 1 and 2 that the anisotropic probability distribution produces long and smooth spiral arms or ring-likefilamentsdepending on 6 : a and /: higher anisotropy leads to more organized and thin spirals, whereas higher I yields less steep spirals or ring-like filaments. Broader description of these results are given in Jungwiert k Palous (1994). We also conclude that with small sized star-forming regions the spiral arms are thin and smooth. They keep the overall shape, length and regularity provided the anisotropy is strong. This is not in agreement with Seiden and Schulman (1990), who relate the small size of HII regions to the feathery (flocculent) appearance of galactic disks. The project was supported by the grant No. 205/93/0090 of the Grant Agency of the Czech Republic. JP thanks organizers for the support at the Conference in Puerto Naos.

REFERENCES W. & ELMEGREEN, B. G. 1982 IBM Research Report RC 9726. H. & SEIDEN, P. E. 1978 Astrophys. J. 223, 129. JUNGWIERT, B. & PALOUS, J. 1994 Astron. Astrophys. Submitted. OORT, J. H. 1954 Bull. Astron. Inst. Nether. 12, 177. OPIK, E. J. 1953 Irish Astron. J. 2, 219. PALOUS, J., FRANCO, J. &; TENORIO-TAGLE, G. 1990 Astron. Astrophys. 227, 175. PALOUS, J., TENORIO-TAGLE, G. & FRANCO, J. 1994 Mon. Not. R. astron. Soc. Submitted. SEIDEN, P. E. & SCHULMAN, L. S. 1990 Advances in Physics 39, 1.

CHIANG,

GEROLA,

Spatiotemporal Pattern Driven by a Self-Regulating Mechanism of Star Formation By ANTONIO PARRAVANO Programa de Postgrado en Astronomia y Astrofi'sica, Facultad de Ciencias, Universidad de Los Andes, Merida, Venezuela.

The spatiotemporal pattern of the regions of active star formation in galaxies is the result of cooperative effects that in many cases produce very organized patterns such as spiral arms. The spiral wave theory has been commonly evoked to explain the spiral pattern, where the enhanced density in the arms promotes the coalescence and condensation of clouds and provokes a strong increment in the star formation rate. Nevertheless, the presence of a self-regulating mechanism of the star formation acting on a large scale can also produce an ordered spatial pattern. Here, we include a self-regulating mechanism based on the sensitivity of the condensation and evaporation of small cool clouds upon the radiation density in the 912-1100 A band (Parravano 1987, 1988) produced by massive stars. Since the UV radiation, produced by short living stars, can affect the condensation and evaporation of distant cool clouds, active star formation in adjacent regions tend to occur out of phase (Parravano, Rosenzweig & Teran 1990). This non-local mechanism of self regulation is applied here in order to study its effect on the spatiotemporal pattern in a one-dimensional array of cells which represents a galactic disk ring. The model also includes several local mechanisms of mass exchange between stars, clouds, warm gas, and hot gas. In each cell, the mass-balance of the self-regulated system is represented by: •

>

.

=

•

>

+ «1 >

•

i

/

+ AW ,

•

»

= >

—*0 >

1.

=(l-or)—

+ Ac,

where £_ = a£s £ c +b£'., and £IT = PEw /(«..)- r E,»(»..)' The variables E a , EC) E w , E/,, E r , and £; are, respectively, the mass surface density of massive stars, clouds, warm gas, hot gas, remnants, and long living stars in the cell. Additionally, t* is the average lifetime of massive stars, a is the fraction of mass ejected by the massive stars during its evolution, a and b denote the coefficients for the induced and the spontaneous star formation rates, and (3 and F are the rate of condensation and evaporation of small clouds per unit mass of warm gas and cloud mass, respectively. The term E^(t —
Parravano: Spatiotemporal pattern

76

0 FIGURE

time

1500

Myr

1. Spatiotemporal pattern of the surface density of massive stars.

respectively. The regulating functions f(uuv) and g(uuv) control the mass transition between the warm gas phase and the cloud phase (Parravano et al. 1990). In order to study the effect of the regulating mechanisms in a multizone model, we assume an array of TV cells in a ring. The cells do not exchange mass but are coupled radiatively via the regulating functions f(uuv) and g(uuv) (Parravano et al. 1990). The radiation density uuv in a cell depends not only on the number of massive stars in the cell but also on those in other cells. Figure 1 shows the spatiotemporal pattern of the surface density of massive stars, E, (»,.?), for an array of N = 50 cells (vertical direction; each 1 kpc long), and 500 equally spaced time moments (horizontal direction; each instant 3 x 106 yr apart). At the initial time, all the cells, except for eight, have the state E, = E* = E h = 0 and £ c = 2 x 106 M 0 kpc" 2 . The cells 1, 21, 24, and 40 have the state E9 = S c = Eu, = 0, and EA = 2 x 106 M 0 kpc~2, and the cells 2, 22, 25, and 41 have the state E, = Ec = EA = 0, and E«, = 2 x 106 M© kpc~2. The different grey tones in the spatiotemporal pattern linearly correspond to black for E s = 0 and to white for £j = 3 x 104 MQ kpc" 2 . Even here, where the model excludes dynamical effects such as density waves, the radiative interaction between the cells tends to force the star formation rate in the cells to oscillate out of phase. The resulting spatiotemporal pattern looks like that of a spiral pattern on a ring. Moreover, when a schematic spiral pressure wave is added to the model (not shown here), the expected pattern only appears when the "dynamic" and the "self-regulated" characteristic times are similar. Therefore, the processes included in the present study could cooperate with the dynamical processes to reinforce and maintain the spiral pattern in models of disk galaxies.

REFERENCES PARRAVANO, A. 1987 Astion. Astrophys. 172, 280. PARRAVANO, A. 1988 Astron. Astrophys. 205, 71. PARRAVANO, A., ROSENZWEIG, P. & TERAN, M. 1990 Astrophys. J. 356, 100.

Detection of an Age Gradient along the z-Axis in a Star-Forming Region B y E M I L I O J. ALFARO 1 ' 2 , JOSE FRANCO 3 , E D M U N D O MORENO 3 AND J E S U S C A B R E R A-C ANO 1 - 4 'Instituto de Astrofisica de Andaluci'a, (CSIC), PO Box 3004, Granada 18080, Spain 2

Astronomy Department, Boston University, MA02215, USA 3

Instituto de Astronomi'a, UNAM, Mexico D.F., Mexico

4

Universidad de Sevilla, PO Box 1045, Sevilla 41080, Spain

The age and location of stellar clusters and Wolf-Rayet stars in the third Galactic quadrant are analyzed. The cluster sample has been divided into three age groups: 1) younger than 107 yr, 2) between 10r and 3 x 107 yr and 3) between 3 x 107 and 10* yr. The mean z-locations of these samples in the central region of the Big Dent display a well defined z-age stratification. The existence of such an age gradient seems to corroborate previous hypotheses suggesting that the star formation activity was, probably, triggered by the same strong perturbation which originated the depression. A model in which the Big Dent originated by the collision of a high-velocity cloud with the Galactic disk is able to reproduce the observed gradient. The analysis of the three-dimensional spatial distribution of a sample of young open clusters within 3 kpc around the Sun led Alfaro et al. (1991) to the discovery that the four nearest supercomplexes, previously detected by Efremov & Sitnik (1988), appear to be located below the formal Galactic plane. The one placed in the third Galactic quadrant (labelled III in ES88) shows the largest and deepest ^-displacement and has been called the Big Dent. The mechanisms able to explain the observed large z-departure as well as the different stages of star formation in the Big Dent need a source of mechanical energy able to inject in the Galactic disk a considerable amount of energy and momentum along the z-axis. The halo-disk interaction, via high velocity clouds (HVCs), has been envisaged as a firm candidate to this power source (Tenorio-Tagle et al. 1986; Franco et al. 1988; Alfaro et al. 1991; Comeron k Torra 1992). Obviously, realistic modelling of HVC impacts relies on the physical constraints derived from the analysis of the structures suspected to originate by this mechanism. The time evolution of the star-forming complexes then becomes a notable clue to the understanding of how this process should work. Here we deal with the history of the star formation in the Big Dent. We analyze the existence and possible origin of a z-age gradient in this supercomplex and compare it with that predicted by a simple analytical model, including the effects of interstellar drag and radiation pressure, of the evolution of a cloud formed in a HVC-disk collision. The model provides the locations, as a function of time, of this star-forming cloud and their newly formed stars. The stars are not affected by the interstellar drag and can detach from their parent cloud, generating the z-age correlation through a series of different episodes of star formation. The sample of open clusters has been split into three groups of age: 1) younger than 107 yr, 2) between 107 and 3 x 107 yr and, 3) between 3 x 107 and 108 yr. The mean age for each group is l ) 4 . 3 ± 3 . 5 x 106 yr, 2) 22.9± 5.9 x 106 yr and 3) 60.7 ± 19.4 x 106 yr. The WN stars have been considered to be coeval with the youngest open clusters. The subsamples were further analyzed to get averaged z values for two different positions inside the area covered by the Big Dent. Our analysis yields the results shown in Figure 77

78

E. J. Alfaro et at.: Age Gradient along the z-Axis

o

i

o

o o

_, ^ — %

1

t—1 /

•—€ J-J'

'_

o i

/

• —

1

/

'

c)

I

/ ¥— /

1

1

±I

/ / -/->

WN stars

;

-

WN stars

>

/ 1

.

20

1

1

40 Age (106 yr)

60

.

_

1

80

FIGURE 1. z—Age diagram for the four object samples. Black squares and open circles indicate the average 2-value at two different positions inside the area covered by the Big Dent. The model solution for a gradient originated by a HVC-disk collision is shown as a dashed line.

1. The signal-to-noise ratio estimated for each set allows us to verify the case for a z-age gradient in this star-forming region with a sound statistical significance. We have modelled the kinematics of a cloud generated by a HVC-disk collision, including both the interstellar drag and radiative pressure. The effective radiative force on the cloud is considered to be produced by newly formed stars. These stars are assumed to be formed at a series of different episodes within the cloud. The solution shown in Figure 1 represents the estimated 2-location for the different burst of star formation, assuming that the present time corresponds to 4 x 107 yr after the previous passage through the plane. The model provides an age gradient which is in reasonable agreement with the observed values, and indicates that an HVC-disk collision is certainly able to explain the vertical structure of the Big Dent. Nevertheless, this result does not discard other possible mechanisms for generating this structure but adds new support to previous hypotheses, suggesting that the star formation activity was probably triggered by the same strong perturbation which originated the depression.

REFERENCES E. J., CABRERA-CANO, J. & DELGADO, A. J. 1991 Astrophys. J. 378, 106. COMERON, F. & TORRA, J. 1992 Astron. Astrophys. 261, 94. EFREMOV, YU. N. & SITNIK, T. G. 1988 Soviet Astron. Lett. 14, 347.

ALFARO,

FRANCO, J., TENORIO-TAGLE, G., BODENHEIMER, P., ROZYCZKA, M. & MIRABEL, I. F.

Astrophys. J. 333, 826. TENORIO-TAGLE, G., BODENHEIMER, P., phys. 170, 107.

ROZYCZKA,

M. &

FRANCO,

1988

J. 1986 Astron. Astro-

Abundances of HII Regions and the Chemical Evolution of Galaxies By MANUEL PEIMBERT, PEDRO COLIN AND ANTONIO SARMIENTO Instituto de Astronomia, UNAM, Apdo. Postal 70-264 Mexico 04510 D. F., Mexico We discuss the observational constraints to the chemical evolution of the interstellar medium, ISM, provided by the abundances of HII regions. We present a review of the results derived from these constraints for the solar vicinity and for metal-poor galaxies. It is found that, contrary to previous results, black holes do not play an important role in the chemical evolution of galaxies. Chemical evolution considerations indicate that substellar objects (M < 0.08 MQ) have a mass density smaller than 0.02 MQ pc~ 3 (2
1. Introduction HII regions are excellent probes of the chemical composition of the ISM of the Galaxy and of other galaxies. From the study of bright HII regions it is possible to derive accurate abundances of H, He, C, N, 0 , Ne, S, and Ar for galaxies that are located at many megaparsecs from us. Two important sources of error in the abundance determinations of bright HII regions are the temperature structure and the fraction of heavy elements that is embedded in dust. The observed HII region abundances can be compared with those predicted by models. To produce a chemical evolution model for a given galaxy the following ingredients are needed: a) a set of stellar evolution models for all masses, b) an initial mass function, IMF, c) a star formation history, and d) a mass flow history. In § 2 we will discuss chemical evolution models, presenting some of the most popular, in § 3 we will compare chemical evolution models with observational constraints for the solar vicinity, and in § 4 we will review some of the ideas that have been proposed to explain the chemical evolution of irregular and blue compact galaxies. 2. Chemical evolution models 2.1. Net yield The heavy-element yield is defined by y(Z) = M(Z)/Mt,

(2.1)

where M(Z) is the mass that a generation of stars ejects as newly-formed heavy elements to the interstellar medium and M* is the mass of the same generation that remains locked into stellar remnants and long-lived stars, where we include the low-mass end of the IMF which might comprise objects that do not become stars. Similarly the oxygen yield is defined by 2/(0) = M(0)/M,. (2.2) The yields in equations (2.1) and (2.2), hereafter net yields, have been also defined as true yields or real yields (e.g. Peimbert 1985, Pagel 1987 and references therein). 79

80

Peimbert et al.: Chemical Evolution of Galaxies

To derive the net yield it is necessary to adopt: a) an IMF, and b) a set of stellar evolution models including the whole mass range expected. The net yields can be compared with "observed yields". The observed yields can be computed by means of: a) the determination of the abundances by mass in a star or in an HII region, b) a model of galactic chemical evolution that should include the initial gaseous abundances by mass and the history of gas flows in and out of the representative volume under consideration, and c) the mass that participates in the chemical evolution process, Mj, and the mass that does not participate in the chemical evolution process, Mni,] Mi, includes: the interstellar mass gas, Mg, the stellar mass, and the low-mass end of the IMF; M n j includes the non-baryonic mass and the baryonic mass that does not participate in the chemical evolution process like primordial black holes. Mj and M n j should be evaluated in a representative volume where the stars formed or the HII region is located. 2.2. Basic equations The basic equations of a chemical evolution model in which the instantaneous recycling approximation, IRA, is assumed are the following (Tinsley 1980, Maeder 1992): dMb

f

(oi\

f

-£- = fi-fo, (2-3) where / / and fo are the inflow and outflow rates. Similarly Mg changes according to: ^

fi-fo,

(2.4)

where R is the mass fraction returned to the interstellar medium by a generation of stars, and ip is the star formation rate. On the other hand, the heavy-element mass fraction in the ISM, Z, evolves via: Mg^ = y(Z)(l-R)il> + {Zi-Z)fI,

(2.5)

where y(Z) denotes the "observed" yield and Zi is the Z value of the accreting material. Equation (2.4) can be rewritten in terms of the gas mass fraction, (J, — Mg/Mb, as: ^

(//-/o)(l-/i).

(2.6)

Time can be eliminated by combining equations (2.5) and (2.6) as follows:

dZ

y{Z){lR)^ + {ZIZ)fI

-(1 For most of the heavy elements the IRA is very good because they are produced by stars more massive than 8 MQ. This is the case for 0 , Ne, Ar, and S, elements that are easily observed in bright HII regions. For C, N, Fe, and He, deviations from IRA have to be taken into account. 2.3. Closed-box model When there are no gas flows into or out of the system, / / = / o = 0, the solution to equation (2.7) is given by

v(z) = ^

(2.8)

Peimbert et al.\ Chemical Evolution of Galaxies

81

Notice that frequently equation (2.8) has been used with the total mass of the system, Mr, instead of Mj, a correct procedure only if M n j is negligible. The closed-box model has also been called the simple model. The yield given by equation (2.8) under the assumption that Mn\, — 0 is called the effective yield (e.g. Peimbert 1985, Pagel 1987 and references therein). The yield of agiven heavy element, yi, is obtained by substituting Z for the mass concentration of element i, X{ in equations (2.5), (2.7), and (2.8). 2.4. Infall Under the assumption that the inflow rate is given by / / = a ( l — R)ip, where a is a free parameter, and integrating equation (2.7) for the case where infall and stellar mass loss are balanced by star formation (Zj = 0,/o = 0, and a = 1), we obtain (Larson 1972, Tinsley 1980)

while for other values of a the yield is given by (Matteucci & Chiosi 1983, Peimbert, Sarmiento k Colin 1994)

(210)

Equation (2.10) is valid if a < 1, otherwise it only applies for /i > / i m i n , where Umin = (a - l ) / a ; when ^ -> nmin, Mb -> oo. 2.5. Outflow of well-mixed material Under the assumption that the outflow rate is given by fo = A(l — R)ip, where A is a free parameter and integrating equation (2.7) for / / = 0 we obtain (Matteucci k Chiosi 1983)

In general when inflow and this type of outflow are occurring and a linear variation of the yield with Z is permitted, the solution to equation (2.7) is given by (Peimbert et al. 1994)

-

a Z l t

(2.12)

where we have assumed that y(Z) = j/o + °-Z. Equation (2.12) is valid if (a — A) < 1, otherwise it only applies for fi > fimin = [(a - A) — l]/(a — A). In deriving equations (2.11) and (2.12) it has been assumed that the fresh Z material ejected by the stars to the interstellar medium is completely mixed with the gas before it is ejected to the interstellar medium, the wind associated with this outflow has been called normal or ordinary. Notice that this outflow reduces the yields of the elements but not their ratios; in particular, it does not affect the AY/AZ ratio (under the assumption of IRA). 2.6. Outflow of Z-rich material It has been proposed that outflow of Z-rich material is present, at least in some galaxies, to explain: a) the large helium abundances derived in some irregular and blue compact galaxies, b) the large AY/AZ ratios derived from samples of galaxies, c) the

82

Peimbert et al.\ Chemical Evolution of Galaxies

high Fe/O ratio in the Magellanic Clouds, and d) the Z-Mass relation present in irregular galaxies (e.g. Aparicio, Garcia-Pelayo & Moles 1988; Russell, Bessell & Dopita 1988a,b; Lequeux 1989; Moles, Aparicio, t Masegosa 1990; De Young & Gallagher 1990; Pilyugin 1993, Tosi 1994). Pilyugin has presented quantitative estimates of this effect on the AY/AZ(O) ratio. By assuming that a fraction 7 of the freshly made Z(O) is ejected to the intergalactic medium without mixing with the interstellar gas, and that the decrease of Mg due to this outflow is negligible, it follows that 2/(0) = (1 - j)y(O)SN + UNWIND .

(2-13)

where y{O)SN is the net yield due to SN explosions and y(O) WIND is the net yield due to the mass loss due to stellar winds before the SN explosion of the massive stars. About half of the freshly made He is ejected by stars in the 8-120 MQ range and about half by stars in the 1-8 MQ range, while almost all the 0 is produced by stars in the 8-120 MQ range and most of it is ejected to the interstellar medium by supernova explosions. Stars with main sequence masses in the 1-8 MQ range end their lives as white dwarfs and eject their excess mass in normal stellar winds in the red giant phase and as superwinds in the preplanetary nebula phase, the expanding shells have typical velocities of a few tens of km s - 1 a n d the material ejected is not expected to escape the parental galaxy. Alternatively, stars in the 8-120 MQ range lose part of their mass in stellar winds and afterwards explode as supernovae. The SN ejecta have typical initial velocities of a few thousands of km s" 1 , and a fraction of the SN ejecta is expected to leave the parental galaxy. Under these considerations it follows that AY 1 Jj,

_ y(Yj- 8 ) + (1 - l)y{y%-i2o)sN + y{y%-\2o)wiND ~ (1 - 7)y(O)SN + y(O)WIND - ^

(2-14)

where b is the fraction of the net He yield due to stars in the 1-8 MQ range and by winds of stars in the 8-120 MQ range; (1 — 6) is the fraction of the net He yield produced by SN explosions due to stars in the 8-120 MQ range; c is the fraction of the net 0 yield due to winds by stars in the 1-120 MQ range, and (1 — c) is the fraction of the 0 yield produced by SN explosions in the 8-120 MQ range. From the computations by Maeder (1992) for Z = 0.001 and the high mass-loss rate case, together with the IMF by Kroupa, Tout, & Gilmore (1993) it follows that 6 = 0.56 and c = 0.0, while for the Salpeter (1955) IMF it follows that 6 = 0.42 and c = 0.0. Consequently for low Z values, the c term can be neglected. Under the assumption that the two types of outflow are present, the [AY/AZ(O)]obs will still be given by equation (2.14) but the net oxygen yield would be given by

[c

+ (i _ c ) ( i _ 7 ) ] l n [ ( A + ! ) „ - ! - A] •

(2 15)

-

3. Solar vicinity 3.1. Observational constraints There are four observational constraints that can be used to study the chemical evolution of the solar vicinity: Z(O), AY/AZ, AY/AZ(O), and Z(O)/[Z - Z(O)]. In this discus-

Peimbert et al.: Chemical Evolution of Galaxies

83

sion we will use for Z(0) the solar value, for AY/AZ the M17 value, for AY/AZ(O) and Z(O)/[Z — Z(O)] an average of the solar value and that of HII regions of the solar neighborhood. The adopted values are: Z(0) = 0.0096 ± 0.0009, AY/AZ = 2.5 ± 0.5, AY/AZ(0) = 5.2 ± 1.4, and Z(O)/[Z - Z(O)] = 0.92 ± 0.10 (see Peimbert et al. 1994 and references therein). The Z(O) value for HII regions depends on their temperature structure; it is often assumed that the mean square temperature fluctuation, i 2 , is zero, and in this case the derived Z(O) is two or three times smaller than ZQ(O). Recent results indicate that t2 RS 0.04 for HII regions of the solar neighborhood (Peimbert, Torres-Peimbert k Ruiz 1992; Peimbert, Storey k Torres-Peimbert 1993a; Peimbert, Torres-Peimbert k Dufour 1993b) which increases the Z(O) of HII regions to values very similar to Z © (0). 3.2. IMF A recent determination of the IMF for the solar neighborhood is that by Kroupa, Tout, k Gilmore (1993), hereafter KTG. The KTG IMF is given by r0.035m-13 £ = I 0.019 m - 2 2 [0.019m-27

if0.08<m<0.5, if 0.5 < m < 1.0, if 1.0 < m < 120,

(3.16)

where m is the stellar mass in solar units and £,{m)dm is the number of stars in the mass interval from m to m + dm. The 0.08 < m < 0.5 exponent for the 95 per cent confidence interval is in the —0.70 to -1.85 range and the center is given by —1.3. We will define three IMFs, KTGa, KTGb and KTGc, that are obtained by extrapolating the KTG IMF to zero mass and adopting the center and the two extremes of the 95% confidence interval; the corresponding IMFs are given by £ a (m) = 0.051m- 07 ,

(3.17)

13

6(m) = 0.035m- , £c(m) = 0.025m-

185

(3.18) ,

(3.19)

for 0 < m < 0.5, while for m > 0.5 we have adopted equation (3.16). The KTG IMFs have the same slope for m > 1 as the Scalo (1986) IMF and the slope only differs for rn < 1 values. 3.3. Stellar models and chemical evolution models of the solar vicinity It is possible from stellar evolution models and a given IMF to compute the following net yields and net yield ratios: y(O), y(Y)/y(O), y(Y)/y(Z), and y(O)/y(Z - O). These values will depend on M(BH), the mass above which the stars end their evolution as black holes without enriching the interstellar medium with heavy elements. Black holes have been postulated by many authors to explain low values of the oxygen yield derived for the Galaxy and for other galaxies (e.g. Twarog k Wheeler 1982, 1987; Mallik k Mallik 1985; Maeder 1992, 1993). For different IMFs y(Y)/y(Z), y(Y)/y(O), and y(O)/y(Z - O) can be directly compared with the observational constraints AY/AZ, AY/AZ(O), and Z(O)/[Z - Z(O)], respectively; this comparison is independent of the infall rate if the infall is due to gas with primordial abundances. Alternatively, j/(O) depends on the specific chemical evolution model adopted for the solar vicinity. From four recent models of the solar vicinity it is found that y(O) - (0.7 ± 0.1) Z(O) (Twarog 1980a,b; Pagel 1989; Meusinger k Stecklum 1992; Malinie et al. 1993).

84

Peimbert et al.: Chemical Evolution of Galaxies

0.008-

A.85

0.006-

• \60 i j \ \ 40

Q.

\ 0.004-

\2 7 -5

O.OOfr Kc

- •

onnoi

8

10

12

14

16

y(Y)/y(O)

FIGURE 1. Oxygen yield, y(O), versus helium to oxygen yield ratio, y(Y)/y(O), based on the KTG IMF and the computations by Maeder (1992). The numbers inside the figure denote M(BH) values. The square is the permitted region for the Z(O) and AV/AZ(O) values for the solar vicinity. The narrow column is the permitted region for the Z(O)/[Z-Z(O)] value. The small rectangle bounded by solid lines is the only permitted region that satisfies the three observational constraints.

Peimbert et al. (1994) have found that the Z(O), AY/AZ, and Z(O)/[Z - Z(O)] values for the solar vicinity presented in § 3.1 can only be fitted by the KTGb IMF together with stellar evolution models by Maeder (1992); this fit is made without the presence of black holes. Furthermore, they also find that, regardless of the value for M(BH), the Salpeter IMF cannot satisfy the three restrictions at the la level. In what follows we will consider AY/AZ(0) instead of AY/AZ; AY/AZ(O) is more convenient because the models by Maeder (1992) do not consider the production of C, N, and Fe by intermediate mass stars, IMS, those in the 1 < M/MQ < 8 range; moreover for these stars, deviations from IRA can occur and should be taken into account. From the Z(O), AY/AZ(0), and Z{O)/[Z - Z{O)} constraints presented in § 3.1, the KTG IMF's and the stellar evolution models by Maeder (1992), we have developed Figure 1. From this figure it is found that at the 1
Peimbert et al.: Chemical Evolution of Galaxies

85

Maeder (1992); she also considers the enrichment of C, 0 , and Z by single IMS based on the computations by Renzini & Voli (1981) and of He based on the computations by Maeder (1992); moreover, she considers the enrichment due to Type la and Ib supernovae produced by binary evolution of IMS. Carigi does not reproduce the observed abundance pattern in the solar vicinity for any value of M(BH) by using the Salpeter IMF. The best agreement by Carigi with the observational constraints is obtained with the Scalo IMF and without the presence of black holes. From the KTGb IMF it is possible to derive a mass density of 8.5 x 10~3M© pc~ 3 in substellar objects by integrating equation (3.18). This value is considerably smaller than the 0.1 MQ pc~ 3 estimate of unobserved mass in the solar vicinity by Bahcall (1984) and Bahcall, Flynn k Gould (1992), and implies that if the missing mass amounts to 0.1 MQ pc~ 3 it is not due to substellar objects and should be of the type that does not participate in the chemical evolution process. 3.4. Conclusions Carigi (1994) and Peimbert et al. (1994) find, from independent considerations, that the production of black holes by massive stars is not needed to explain the chemical evolution of the solar vicinity. This result contradicts previous results by other authors who find that objects more massive than 20-28 MQ end their lives as black holes without enriching the interstellar medium with heavy elements. Carigi (1994) finds that while the Scalo IMF is in agreement with the observational constraints, the Salpeter IMF is not; similarly, Peimbert et al. (1994) find that while the KTG IMF is in agreement with the observational constraints, the Salpeter IMF is not. Peimbert et al. (1994) find, from chemical evolution considerations, that the mass of substellar objects in the solar vicinity is smaller than 0.02 M Q pc~ 3 (2
4. Metal-poor galaxies 4.1. Observational constraints A recent discussion of the problems involved with the chemical evolution of metal-poor galaxies has been presented by Lequeux (1993). One of the most important constraints for the study of the chemical evolution of metal-poor galaxies is provided by the Z-/J, diagram. In Figure 3 we present a Z(O)-fi diagram from the Mg and My data by Staveley-Smith, Davies k Kinman (1992), where Z(O) was derived from the (O/H)-L B relationship presented by Skillman, Kennicutt k Hodge (1989). Lequeux et al. (1979) presented the Z-/J, diagram for a group of irregular and blue compact galaxies. They found that y(Z) — 0.004 together with a closed-box model could explain the observations. Notice that this value of the yield is a factor 3.5 times smaller than that derived from the KTG IMF and the Maeder (1992) stellar evolution models (see below). Peimbert & Serrano (1982) tried to explain the distribution of irregular galaxies, blue compact galaxies, the Galaxy, and M83 by suggesting that the yield increases with metallicity as y(Z) = 0.002 + 0.6Z. All the objects but one considered by them fall in the region limited by the closed-box model and an infall model with a = 1. They gave arguments in favor of a change of the low-mass end of the IMF to explain the increase of y with Z. Matteucci & Chiosi (1983) and Chiosi k Matteucci (1984) suggested that in addition

86

Peimbert et al.: Chemical Evolution of Galaxies

to infall, outflow models of well-mixed material could also explain the position of irregular and blue compact galaxies in the Z-/J, diagram. Kumai & Tosa (1992) suggested that the Z-fi diagram could be explained with a closed-box model and a varying fraction of Mni,/Mb. Another important observational constraint is the AY/AZ. As mentioned in § 2 this ratio is affected by outflow of metal-rich material and is not affected by the gas flow behavior of other models. The presence of black holes could also be important in explaining the Z(O)-(J, diagram because they reduce the y(Z) value. Therefore we will also consider the Z(O)/[Z — £(O)] ratio which is very sensitive to M(BH). In what follows we will discuss the different suggestions that have been proposed to explain the Z(O)-fi diagram. We will consider the Z(O)-fi diagram instead of Z-fi diagram because: a) for most objects we only have a Z(0) value, and b) there are some elements, like C and Fe, for which the IRA is not as good as for 0 . The two main characteristics that need explanation are: a) the large range of /z values for a given Z(0) value, and b) the low (Z(0)) associated to the (/z) value. 4.2. Stellar models and black holes

For our discussion we will make use of the Maeder (1992) stellar evolution computations for Z — 0.001 and Z = 0.02 with the high mass-loss rate. From these models the Z production considered is due to stars with M > 8 MQ for which the IRA is excellent, moreover the Z(O)/Z ratio is a strong function of the stellar mass. In Figure 2 we present the y(O)/y(Z — O) ratio based on Maeder's computations and the KTG and Salpeter IMFs, notice that this ratio is independent of the behavior of the IMF for M < 8 M©; this ratio is compared with the observed Z(O)/[Z - Z(O)] value determined for the best observed metal poor HII region, NGC 2363; from this HII region it is obtained that Z(O)/[Z - Z(O)] = 1.48 ± 0.15 (Peimbert, Pena & Torres-Peimbert 1986). This observational restriction is presented in Figure 2. From this figure it follows that there is an excellent agreement between the KTG IMF and Maeder's model without the need to invoke black holes. Alternatively, with the Salpeter IMF, M(BH) = (83 ± 15) M 0 ; by considering the production of C, Fe, and N by stars with M < 8 MQ, the y(O)/y(Z — O) ratio becomes even smaller, strengthening the result that black holes do not play a role in the chemical evolution of metal-poor galaxies. 4.3. Closed-box model

From the computations by Maeder (1992) and the KTG IMF we find that y(Z) = 0.014 and 2/(0) = 0.0083 for Z = 0.001. In Figure 3a-3b we present the closed-box solution for y(O) = 0.0085. It is clear that a closed-box solution, under the assumption that Mnb
Peimbert et al.: Chemical Evolution of Galaxies

20

40

60

80

100

87

120

FIGURE 2. y(O)/y(Z - O) versus M(BH) (in solar mass units) diagram for the IMFs by KTG and Salpeter based on the computations by Maeder (1992). The yields ratio depends only on the IMF for m > 1, consequently there are no differences among cases a, b, and c. The permitted band corresponds to the Z(O)/[Z-Z(O)\ observed value of 1.48 ± 0.15.

and He abundances might be affected by the presence of Wolf Rayet stars (Pagel et al. 1992), but the O/H ratios derived from different HII regions in the SMC and the LMC (e.g. Dufour 1984 and references therein) and for different HII regions in I Zw 18 (Skillman & Kennicutt 1993) indicate that the self-enrichment in O is not very important. Furthermore, the presence of O self-enrichment makes it more difficult to explain the Z(O)-fx diagram, since we need to increase 2/(0) not to lower it. The problems associated with the determination of Mr could be greater. A line profile obtained with a single radio dish cannot be converted into a reliable total mass estimate because the inclination and radius derived from optical images can be misleading (Skillman 1992); moreover, the assumption that the velocity field is due to rotation might not be correct in some cases (e.g. Skillman k Kennicutt 1993), nor the assumption that rotation is completely absent in others. Reliable mass estimates must be based on interferometric observations and only a few objects have reliable total mass determinations. 4.5. Variations in the IMF To explain the Z(O)-(i diagram for low Z(O) values, smaller values of y(O) are needed. This can be done by reducing the fraction of stars in the high-mass end of the IMF (M > 8 MQ) or by increasing the fraction of objects in the low-mass end of the IMF (M < 1 — 2 M Q ), including substellar objects. Massey et al. (1989a,b), and Parker et al. (1992) find that most of the OB associa-

88

Peimbert et al.: Chemical Evolution of Galaxies

tions in the Magellanic Clouds have IMFs with slopes in the —1.6 < T < —1.9 range. Parker k Garmany (1993) find T = - 1 . 5 ± 0.2 for the 30 Doradus region. These results are in agreement with the KTG slope of - 1 . 7 and do not present evidence for substantial changes in the high-mass end of the IMF with metallicity. Moreover, from the study of many extragalactic HII regions in irregular and blue compact galaxies of different Z(O) content, McGaugh (1991) finds that a fixed IMF adequately explains the observed trends in the spectra, so no variation in the high-mass end of the IMF is inferred. Notice that a Salpeter slope for the high-mass end of the IMF predicts a higher heavy element yield, making it more difficult to explain the Z(O)-fi diagram. There is some evidence in favor of an increase in the fraction of low-mass objects with decreasing Z (e.g. Peimbert k Serrano 1982). Furthermore, Elmegreen (1994) has suggested that smaller-mass galaxies could have a lower low-mass end of their IMF due to their smaller velocity dispersion. In Figure 3c-3d we present a series of models with different outflow rates and with 2/(0) = 0.0012 + 0.67Z, or y(Z) = 0.002 + L I Z . This type of yield reduces the need for strong effects due to other mechanisms to explain the Z(O)-fi diagram.

4.6. AY/AZ(0) The best AY/AZ(0) ratio for irregular and blue compact galaxies is that by Pagel et al. (1992) and is equal to 10.3 ± 3.5, this value was derived without considering the fraction of O tied down in dust grains or the presence of temperature variations inside the observed HII regions. From the results by Skillmanet al. (1994) it is found that about 15 % of O is tied down in dust grains, therefore the AY/AZ(O) by Pagel et al. (1992) should be lowered to 8.8. By considering a modest amount of correction due to t2 the value should be lowered to about 6.7. This value is a factor of 2.4 higher than that computed from the KTG IMF and the stellar evolution models by Maeder (1992) and probably implies the presence of Z-rich outflows. This value is also a factor of 3.3 higher than that computed from the Salpeter IMF. Unfortunately not all the objects in the sample by Pagel et al. (1992) have Mr and Mg determinations. On the other hand, not all the objects in the sample by Matteucci k Chiosi (1983) and by Staveley-Smith et al. (1992) have Y and Z determinations. Colin et al. (1994) have used a subsample of the sample by Pagel et al. (1992) with 10 objects and find AY/AZ(0) = 10.3 ± 1.8 without considering the amount of O tied down in dust grains, or the effect off2 in the O determinations. The (Z(O)} and (log/z) values for the Colin et al. (1994) sample are 0.0016 and —0.52, respectively. 4.7. Inflow Inflow of material with pregalactic abundances, Y = Yp and Z — 0, does not change AY/AZ(O), though it is important in reducing Z for values in the 0 < a < 1.0 range. For larger values of a there is a minimum fi value, see Figures 3a-3b and equation (2.10), and this process is not able to cover the regions of low Z(O) and low fi values. The solutions presented by Matteucci k Chiosi (1983) that cover all the Z-\i diagram for a > 1.0 are incorrect because equation (2.10) does not apply for fi < (a — 1)/Q. 4.8. Outflow of well-mixed material It is possible to cover the Z-fi diagram, to the right of the solution for the closed-box model, by varying A. There are two problems with this solution: a) it does not explain the large AY/AZ(O) observed values, and b) for some galaxies, with low Z(O) and low

89

Peimbert et al.: Chemical Evolution of Galaxies

I/ 1 ; 1 ' ' I ' ' ' '

1 - LSB / :

"

^4i

0

i

i

(a) 4 - BCD/.'

i

-0.5

i

•

,-"• i

i

i

-1

i

(b)

i

i

-1.5

0

-0.5

-1 log

-1.5

FIGURE 3. Z(O)-fi diagram for the sample by Staveley-Smith et al. (1992), where BCD and LSB denote blue compact and low surface brightness galaxies, repectively. Figures a-b present a closed-box model for y(O) — 0.0085 and infall models for different values of the infall parameter a (see eq. 2.10). Figures c-d present well-mixed outflow models for a yield increasing with metallicity given by J/(O) = 0.0012 + 0.67 Z(O) and for different values of the outflow parameter A (see eq. 2.12); the 0.0 curve corresponds to the closed-box model.

90

Peimbert et al.: Chemical Evolution of Galaxies

fi, values of A of about 50 are needed, which would imply the presence of haloes with about fifty times more mass in gas than the amount of mass in the form of stars in the main body of the galaxy, we consider this possibility unlikely. This type of outflow does make the prediction that galaxies should be surrounded by massive gas haloes, a prediction that could be tested. In Figures 3c-3d we present models with a yield increasing with metallicity and various values of A. 4.9. Outflow of Z-rich material

To explain an increase in the AY/AZ(O) ratio due to the outflow of Z(O)rich material we can make use of equation (2.14) and the stellar evolution models by Maeder (1992). For an increase of a factor of 2 in the AY/AZ(0) ratio, for an object with Z = 0.001, a value of 7 = 0.7 is needed for the Salpeter IMF and a value of 7 = 0.64 for the KTG IMF, this would raise the yields derived from Z-fi diagram by a factor 1/(1 — 7) which is of the order of three, reducing the importance of the other suggested mechanisms to explain the Z-fi diagram. Moreover, the objects with small \i and small Z might yield even higher AY/AZ(O) ratios, and this process by itself could explain all the points in the Z-fi diagram. In particular, for the sample of Colin et al. (1994), (Z) = 0.0027, (Z(O)) = 0.0016 and (log^i) = —0.52, which, together with equation (2.8), give y(O)ej = 0.0013, and y(Z)ej = 0.00225. Based on the computations by Maeder (1992) interpolated for Z = 0.0027 and adopting the KTG IMF, we obtain b = 0.58 and c = 0.021; under the assumption that AY/AZ(O) has been increased by a factor of 2.4 due to Z-rich outflow we obtain 7 = 0.73, which, together with equation (2.15 for A = 0), give y(O) = 0.0057 and y(Z) = 0.01, still a factor of 1.4 smaller than the predicted yields from Maeder and the KTG IMF that amount to 0.0083 and 0.014, respectively. This result reduces considerably the importance of other mechanisms under consideration, nevertheless the errors in the AY/AZ(O) determination are large and better determinations are needed to evaluate the importance of this effect. 4.10. Dark matter that does not participate in the chemical evolution process

Kumai & Tosa (1992) have suggested that the Z-fi diagram can be explained with a closed-box model with a constant yield combined with Mnj 7^ 0, where each object would have a different Mnt/Mi ratio. Based on the sample by Matteucci k Chiosi (1983), Kumai & Tosa find a good correlation between the dark matter fraction, fo, derived by two independent methods. The methods to determine fo are based on: a) fitting the Z-fi diagram with a constant yield, and b) determining the difference between the dynamical mass and the sum of the gaseous mass plus the stellar mass derived from an adopted M«/L ratio. Of the effects that affect the Z-fi diagram discussed before there is one that is particularly relevant to this suggestion: errors in the determination of Mr could be responsible for part of the correlation derived by Kumai & Tosa (1992) because they will affect in the same manner the fo values derived by the two methods. Nevertheless, from the sample of Matteucci k Chiosi (1983) there is no correlation between M n j/Mj (or fa) and the mass of the galaxy. We would expect fD to increase for less massive objects (e.g. White L Frenk 1991; Persic &; Salucci 1988; Persic, Salucci & Ashman 1993). Persic et al. find for a sample of 56 galaxies, most of them spiral, that the fractional visible mass at the optical radius, Ropt, decreases for less luminous galaxies. From the sample by Staveley-Smith et al. (1992), Colin et al. (1994) find that there is an inverse correlation of Mni/Mb with MT for the subsample made up of the low surface

Peimbert et al.: Chemical Evolution of Galaxies

91

brightness galaxies of the sample, but there is no correlation between Mnb/Mb and for the subsample of blue compact galaxies. This result might imply that Mnj plays a role in the explanation of the Z-fi diagram for the low surface brightness galaxies but not for the blue compact ones. 4.11. Conclusions The closed-box model with the KTG or the Salpeter IMF and the stellar evolution models by Maeder (1992), under the assumption that Mnf, — 0, cannot explain all the points in the Z(O)-fi diagram. The distribution of points in the Z(O)-fi diagram is not mainly due to: a) errors in the Z(O) determinations, b) the presence of black holes, c) inflow, and d) a varying IMF at the high-mass end. Alternatively, the distribution of points in the Z(O)-/i diagram could be due to: a) errors in My, b) a varying IMF at the low-mass end, c) outflow of well-mixed material, d) the presence of dark matter that does not participate in the chemical evolution process, and e) outflow of Z(O) rich material. The first four possibilities cannot explain the high AY/AZ(O) observed values and one or more of them should be combined with the last to reproduce the Z(O)-fi diagram. Alternatively, if the AY/AZ(O) observed value is considerably higher than the value discussed above, then the derived 7 will become even higher and the last possibility alone might be responsible for the distribution of points in the Z(O)-fi diagram. Determinations of Y, Z(O), Mg, and MT of high qualtity are needed to advance in this problem, particularly of objects with low Z(O) and low fi values.

Manuel Peimbert wishes to acknowledge helpful discussions with many of the conference participants, in particular with Evan D. Skillman, moreover he also wishes to thank the members of the Instituto de Astrofisica de Canarias for the invitation to visit the Institute and participate in this conference and for their warm hospitality. The authors acknowledge Luis A. Martinez for his assistance with the handling of the postscript figures.

REFERENCES APARICIO, A., GARCIA-PELAYO, J. M.& MOLES, M. 1988 A.
287,

74,

375.

926.

BAHCALL, J. N., FLYNN C.& GOULD, A. 1992 Ap. J.

389,

234.

L. 1994 Ap. J. In press. COLIN, P., SARMIENTO, A., CARIGI, L. & PEIMBERT, M. 1994. In preparation. CHIOSI, C. & MATTEUCCI, F. 1984. In Stellar Nucleosythesis (ed. C. Chiosi & A. Renzini), p. 359. Reidel. CARIGI,

DE YOUNG, D. S.& GALLAGHER, J. S. 1990 Ap. J.

356,

L15.

R. J. 1984. In Structure and Evolution of the Magellanic Clouds (ed. S. van den Bergh fc K. S. de Boer), p. 363. ELMEGREEN, B. G. 1994, this conference. GARNETT, D. R. 1992 A. J. 103, 1330. DUFOUR,

KROUPA, P., TOUT, C. A. & GILMORE, G. 1993 M. N. R. A. S.

KUMAI, Y. & TOSA, M. 1992 A. & A. LARSON,

R. B. 1972 Nature 236, 21.

257, 511.

262,

545.

92

Peimbert et al.: Chemical Evolution of Galaxies

LEQUEUX, J. 1989. In Evolution of Galaxies-Astronomical Observations (ed. I. Appenzeller, H.J. Habing & P. Lena), p. 147. Springer. LEQUEUX, J. 1993. In Origin and Evolution of the Elements (es. N. Prantzos, E. Vangioni-Flam k. M. Casse), p. 504. Cambridge University Press. LEQUEUX, J., PEIMBERT, M., RAYO, J. M., SERRANO, A. & TORRES-PEIMBERT, S. 1979

A.

& A. 80, 155. MAEDER, A. 1992 A. & A.

264,

105.

MAEDER, A. 1993 A. & A.

268,

833.

MALLIK,

D. C. V. &

MALLIK,

S. V. 1985 J. Astrophys. & Astion.

6, 113.

MALINIE, G., HARTMANN, D. H., CLAYTON, D. D.&; MATHEWS G. J. 1993 Ap. J. MASSEY,

P.,

GARMANY,

P.,

PARKER,

C. D.,

SILKEY,

M. &

DEGIOIA-EASTWOOD

413,

633.

D.K. 1989a A. J.

97,

107. MASSEY,

J.

W M . & GARMANY,

MATTEUCCI, F. & CHIOSI, C. 1983 A. & A. MCGAUGH, S. S. 1991 Ap. J.

380,

C. D. 1989b A. J. 98, 1305. 123,

121.

140.

MEUSINGER, H.& STECKLUM, B. 1992 A. & A.

256,

415.

MOLES, M., APARICIO, A. & MASEGOSA, J. 1990 A. & A. PAGEL, PAGEL,

228,

310.

B. E. J. 1987. In The Galaxy (ed. G. Gilmore & B. Carswell), p. 341. Reidel. B. E. J. 1989 .Rev. Mexicans. Astion. AstroRs. 18, 161.

PAGEL, B. E. J., SIMONSON, E. A., TERLEVICH, R. J.& EDMUNDS M. G. 1992 M. N. R.

A.

S. 255, 325. PARKER J. WM., GARMANY CD.,

MASSEY P. & WALBORN N. R. 1992 A. J.

PARKER, J. WM. & GARMANY C. D. 1993 A. J.

106,

103,

1205.

1471.

M. 1985. In Star Forming Dwarf Galaxies and Related Objects (ed. D. Kunth, T. X. Thuan &; J. T. T. Van), p. 403. Editions Frontieres.

PEIMBERT,

PEIMBERT, M. & SERRANO, A. 1982 M. JV. R. A. S.

198,

563.

PEIMBERT, M., PENA M. & TORRES-PEIMBERT S. 1986 A. & A. PEIMBERT, M., SARMIENTO A. & FIERRO J. 1991 P. A. S. P.

158,

103,

266.

815.

M., TORRES-PEIMBERT, S. & Ruiz, M. T. 1992 Rev. Mexicans. Astion. Astiofis. 24, 155. PEIMBERT, M., STOREY, P. J. & TORRES-PEIMBERT S. 1993a Ap. J. 414, 626. PEIMBERT M., TORRES-PEIMBERT S. & DUFOUR R. J. 1993b Ap. J. 418, 760. PEIMBERT, M., SARMIENTO A. & COLIN P. 1994 Ap. J. Submitted. PEIMBERT,

PERSIC, M. & SALUCCI, P. 1988 M. JV. R. A. S.

234,

131.

PERSIC, M., SALUCCI P. & ASHMAN, K. M. 1993 A. & A.

279,

343.

PlLYUGlN, L. S. 1993 A. & A. 277, 42. RENZINI A. & VOLI M. 1981 A. & A.

94,

175.

RUSSELL, S. C , BESSELL, M. S. & DOPITA, M. A. 1988a. In Galactic and Extragalactic Star Formation (ed. R. E. Pudritz & M. Fich), P. 106. Kluwer. RUSSELL, S. C , BESSELL, M. S. & DOPITA, M. A. 1988b. In The Impact of Very High S/N Spectroscopy (ed. G. Cayrel de Strobel fc M. Spite), p. 545. Kluwer. SALPETER, E. E. 1955 Ap. J.

121,

161.

SCALO J. M. 1986 Fund. Cosmic Phys. 11, 1. SKILLMAN, E. D. 1992. In Elements and the Cosmos (ed. M. G. Edmunds & R. Terlevich), p. 246. Cambridge University Press. SKILLMAN, E. D., KENNICUTT, R. C. & HODGE P. W. 1989, Ap. J. SKILLMAN, E. D. & KENNICUTT R. C. 1993 Ap. J.

411,

347,

875.

655.

SKILLMAN, E. D., GARNETT, D. R., DUFOUR, R. J., PEIMBERT, M., TORRES-PEIMBERT, S.,

Peimbert ei al.: Chemical Evolution of Galaxies

93

SHIELDS, G. A., TERLEVICH E. & TERLEVICH R. J. 1994. In preparation. STAVELEY-SMITH, L., DAVIES, R. D. & KINMAN T. D. 1992 M. N. R. A. S.

TlNSLEY, B. 1980 Fund. Cosmic Phys. TOSI, M. 1994, preprint. TWAROG, B. A. 1980a Ap. J.

5, 287.

242, 242.

TWAROG, B. A. 1980b Ap. J. S.

44, 1.

TWAROG, B. A. & WHEELER J. C. 1982 Ap. J.

261, 636

TWAROG, B. A. & WHEELER J. C. 1987 Ap. J.

316, 153.

WHITE, S. D. M. & FRENK C. S. 1991 Ap. J.

379, 52.

258, 334.

Galaxy Properties in Different Environments: Star Formation in Bulges of Late-Type Spirals 1 2 By M. G. PASTORIZA !, E. BICA 1 M. MAIA 1 1 C. BONATTO AND H. DOTTORI 1

Departamento de Astronomia, IF-UFRGS, CP 15051, CEP 91501-970,Porto Alegre, RS, Brazil 2

Departamento de Astronomia, Observatorio Nacional, Rua Gal. Jose Cristino 77, Rio de Janeiro, 20921-RJ, Brazil

The star formation history in the nuclei of late-type spiral galaxies is compared between a sample in a high galaxy density medium (HDS) and a control sample (CS) of isolated galaxies. We have observed 20 HDS and 18 CS galaxies selected from a larger list generated by the application of a group-finding algorithm to the SSRS survey. Using equivalent widths of absorption lines and the continuum distribution, we determined the nuclear stellar population types, from those dominated by old populations to those containing star formation bursts of different ages and intensities. The HDS and CS stellar population type histograms are similar, suggesting that environmental influences, at least for the present samples, do not substantially affect the nuclear stellar population. However, the nuclear emission lines indicate that, in the BPT diagnostic diagrams, there is an excess of HDS galaxies located within or close to the AGN loci. For 6 HDS and 2 CS galaxies it was possible to determine Oxygen (O/H) and Nitrogen (N/H) abundances. The samples present similar (O/H) values, but in the CS galaxies the (N/O) ratio is lower at equal galaxy luminosity.

1. Introduction Evidence of several environmental effects that affect galaxy properties have been reported recently: (a) the morphology-density relation (fractional increase of early-type galaxies towards regions of high concentration, Dressier 1980; Postman & Geller 1984; Giovanelli, Haynes & Chincarini 1986; Maia & da Costa 1990); (b) the morphologyclustercentric radius relation (Whitmore, Gilmore & Jones 1993). However, the impact of the local density of galaxies on the star formation rate (SFR) is not well understood. Two effects have been suggested to influence the SFR: (1) the SFR could be enhanced by tidal interaction (Bushouse 1986; Kennicutt et al. 1987); (2) in the very high density regions in the cores of clusters, close encounters leave preferentially anemic spirals (Dressier 1984). The W(Ha) of galaxies within groups tends to be larger than those outside groups (Maia & da Costa 1990). In order to understand the dependence of the SFR on the galaxy environment, we have studied the stellar component and properties of the emitting gas for two galaxy samples: High Density Sample (HDS): p > 18galMpc- 3 , and Control Sample (CS): p < 0.0004gal Mpc~ 3 . t Visiting Astronomer at the Cerro Tololo Interamerican Observatory, operated by the Association of Universities for Research in Astronomy, Inc. under contract with the National Science Foundation. 94

Pastoriza et al.\ Galaxy Properties in Different Environments

95

2. The samples The southern Sky Redshift Survey Catalog (da Costa et al. 1988, 1989, 1991) was taken as data base for the sample selection. This SSRS catalog consists of 2028 galaxies selected from the ESO/Uppsala survey, satisfying the conditions: log[D0] > 0.1, 6" < -30°, 6 < -17.5°. The High Density Sample is formed by galaxies that are in groups of 3 or more members generated by the group finding algorithm described by Maia, da Costa & Latham (1989). The groups have a surrounding density contrast (8p/p) > 500, equivalent to p > 18 gal Mpc~3. Groups are formed by accumulation of galaxy pairs with a member in common having a projected separation D12'. D12 = 2{V/H0)am(912/2)(Vy^3, Vo = 400 km s" 1 , and (V) is the selection function for the galaxy Catalog normalized to VF = 1000 km s" 1 . The Control Sample is made up of galaxies that are not assigned to any group in the search procedure for neighbors. This search was performed at (Sp/p) = 0.01, equivalent to p< 0.0004 gal Mpc- 3 . All the selected objects have radial velocity (after correction for the Virgo infall) VR < 8000 km s" 1 . HDS contains 151 galaxies, and the CS contains 179 galaxies. 3. Subsamples of late-type spirals (Sb to Sc) We have selected HDS and CS late-type spirals brighter than m? = 14 mag, and avoided edge-on galaxies, since our spectral analysis deals with the central region within 1.2 kpc. The HDS late spiral sample contains 50 galaxies and the CS, 79. 4. The observations Observations were carried out at the 1.0-m telescope of CTIO (Chile), with the Cassegrain Spectrograph and 2D-Frutti detector; spectral range: AA3600 - 7000 A, resolution 5 A pixel"1 at A5000 A. Reductions were made with IRAF at our University. 5. Absorption line spectra The stellar population analysis is based on the equivalent width of the absorption features. Continuum tracing and spectral windows are those defined by Bica & Alloin (1986). A typical error is a w 0.5 A. The equivalent widths were used to classify the galaxies in terms of templates which span the same properties of the stellar population in normal galaxies. Characteristics of the templates: SI to S3 — red population, metallicity decreasing from 4 to 1 Z©; S4 to S7 — increasing contribution from blue components, younger than lGyr: from 10% to 70% of flux at A5580 A. Most of the HDS and CS galaxies match these templates (see example in Figure 1). No substantial differences are found between the stellar population types for the two samples as shown in Figure 2.

Pastoriza et al.: Galaxy Properties in Different Environments

96

U

4000

5000

6000

MA) FIGURE

1. Top panel: example of a red stellar population; bottom panel: blue stellar population.

6. Emission lines and nuclear activity We have subtracted the stellar population contribution from the calibrated spectra normalized at A5870 A. We found all types of nuclear activity among the galaxies of the samples. The diagnostic diagrams using emission-line ratios show that most of the HDS galaxies are located in the region occupied by active galaxies or transition between HII regions and AGNs, while the CS galaxies have HII region-like spectra.

Pastoriza et al.: Galaxy Properties in Different Environments

97

1O HDS

6

--

4

-

2

-

8

-

N

cs

6 N 4

-

2

-

-

S1

S2 S3 E7 S4 S5 S6 Stellar Population Type

S7

FIGURE 2. The distribution of stellar population does not show any remarkable differences between CS and HDS galaxies.

7. Oxygen and nitrogen abundances Evidence that spiral galaxies in the Virgo Cluster have systematically higher interestellar abundance than the field ones have been presented by Shields, Skillman & Kennicutt (1991). However, Henry et al. (1992) found that Virgo spiral and field galaxies are indistinguishable in terms of abundance properties. In Figure 3 are plotted log(N/O) and 12 + log(O/H) vs. absolute magnitude. Wefindthat independently of galaxy luminosity, HDS and CS galaxies present similar values of (0/H), and, for equal luminosity, CS galaxies present a tendency of lower (N/0) than HDS galaxies. These findings suggest that the environment does not affect (0/H) but may have some influence on (N/O).

8. Conclusions The main conclusions of this paper are: (a) the environment does not play a significant role in the stellar population type of the bulges of the late-type spirals; (b) Seyfert-like nuclear activity is more frequently found in HDS galaxies; (c) the environment does not have much effect on the oxygen abundance, but may have some influence on the nitrogen for HDS galaxies.

98

Pastoriza et al.: Galaxy Properties in Different

Environments

0.0

ocs

r

-0.5

_.

-1.0

HDS

N4303

p CD

-1.5 A N3310

A LMC

-2.0

-2.5

N43O3

9.5

A A MS1

9.O

N3310A LMC

o

IS

AM81

£ 8.5 8.0

7.5

-22

-20

-18

-16

FIGURE 3. Chemical abundance for the 2 samples plotted against Absolute Magnitude.

REFERENCES BlCA, E. & ALLOIN, D. 1986 Astron. Astrophys. 162, 21. BUSHOUSE, H. A. 1986 Astron. J. 91, 255. DA COSTA, L. N., PELLEGRINI, P. S., SARGENT, W. L. W., TONRY, J., DAVIS, M., MEIKSIN, A., LATHAM, D. W., MENZIES, J. W. & COULSON, I. A. 1988 Astrophys. J. 327, 544. DA COSTA, L. N., PELLEGRINI, P. S., WILLMER, C , DE CARVALHO, R., MAIA, M., LATHAM, D. W. & GEARY, J. C. 1989 Astron. J. 97, 315. DA COSTA, L. N., PELLEGRINI, P. S., DAVIS, M., MEIKSIN, A., SARGENT, W. L. W. & TONRY,

J. 1991 Astrophys. J. Suppl. 75, 935.

Pastoriza et al.: Galaxy Properties in Different Environments

99

DRESSLER, A. 1980 Astiophys. J. 236, 351. A. 1984 Astion. Astrophys. 22, 185. GIOVANELLI, R., HAYNES, M. P. & CHINCARINI, G. L. 1986 Astiophys. J. 300, 77. DRESSLER,

HENRY, R. B. C ,

PAGEL, B. E. J., LASSETER, D. F. & CHINCARINI, G. L. 1992 Mon.

Not.

R. Astion. Soc. 258, 321. KENNICUTT, R. C ,

KEEL, W. C , VAN DER HULST, J. M., HUMMEL, E. &; ROETTIGER, K.

A. 1987 Astion. J. 93, 1011. MAIA, M. A. G. & DA COSTA, L. N. 1990 Astiophys. J. 349, 477. POSTMAN, M. &; GELLER, M. J. 1984 Astrophys. J. 281, 95. SHIELDS, G. A., SKILLMAN, E. D. & KENNICUTT, R. C. 1991 Astrophys. J. 371, 82. WHITMORE, B. C , GILMORE, D. M. & JONES, C. 1993 Astrophys. J. 407, 489.

Star Formation in Galaxies in the Bootes Void ByDONNA WEISTROP Department of Physics, University of Nevada, Las Vegas, NV 89154, USA Observations are reported for 16 of the 25 galaxies found in the Bootes void. At least five of the galaxies are spirals, and another five are disk systems. Two interacting galaxy pairs have been definitely identified, and there are additional candidate pairs. Several of the galaxies are luminous in Ha, due in most cases to significant amounts of star formation. The observed galaxies do not resemble the galaxies predicted to inhabit voids. There is evidence for structure in the spatial distribution of the galaxies, in particular half the galaxies are located in a single plane 10 Mpc wide.

1. Introduction The Bootes void, a spherical region with radius 62 Mpc at a distance of 310 Mpc (Ho = 50 km s" 1 Mpc" 1 ), is one of the largest known low-density regions in the largescale distribution of galaxies (Kirshner et al. 1981, 1987). Projected on the sky its diameter subtends an angle of 23°. Identifications of 25 galaxies in Bootes have been published. Almost all of these galaxies were discovered from spectroscopy of samples selected from IRAS and objective prism surveys (see Weistrop et al. 1992 for references). Comparisons with more populated parts of the universe indicate the galaxy density in the Bootes void is about one-third the normal density (Weistrop 1989; Dey, Strauss, & Huchra 1990). This region is therefore a good one in which to investigate the effects of a low-density environment on galaxy evolution, and to compare the nature of the galaxies observed in voids with the predictions of the types of galaxies to be found in voids. Various predictions have been made concerning galaxies in voids. Brainerd k. Villumsen (1992) speculate that voids may be occupied by low-mass galaxies in low-mass halos. Other suggestions include populating voids with giant disk galaxies with low surface brightness (Hoffman, Silk, & Wyse 1992) or the absence of luminous galaxies in voids due to the lack of tidal interactions to trigger star formation (Lacey et al. 1993). We have obtained broad-band and Ha images, spectra, and CO observations of void galaxies, to compare them with galaxies in denser parts of the universe, and to compare our observations with the galaxies predicted to be in voids. We report here principally the Ha results, although we also refer to the broad-band imaging and spectroscopy, and discuss the distribution of galaxies within the Bootes void.

2. Morphology We have obtained imaging for 16 of the 25 galaxies in the Bootes void. Twelve galaxies were imaged in Ha and four additional galaxies in BVRI. Among the 16 galaxies, there are at least five spirals, five flattened, edge-on disks, three galaxies with significant amounts of Ha emission both within and outside the nucleus which could not be further classified, one featureless companion, one slightly flattened system, and one system with peculiar morphology. Two interacting systems have been definitely identified. 1519+5050 consists of a Seyfert 1 plus a one-armed spiral experiencing a considerable amount of star formation 100

Weistrop: Star Formation in the Bootes Void

101

FIGURE 1. Broad-band / image of the central region of 1517+3956. The velocity difference and morphology suggest this object is a pair of interacting or merging galaxies. (Weistrop et al. 1992). The second pair is a galaxy with Ha emission within and outside its nucleus with a featureless companion at a projected distance of 39 kpc (Dey et al. 1990). A third system, 1517+3956, has a peculiar morphology (Figure 1). Preliminary analysis of spectra suggest there may be a velocity difference of 300 km s" 1 between the northern and southern brightness concentrations, suggesting an interacting or merging pair.

3. Ha luminosities and star formation rates Ha fluxes were derived from on- and off-band Fabry-Perot observations. The uncertainty in the flux is estimated to be ±25%, due principally to the uncertainty in the conversion from the count rates. Our Ha fluxes agree within 25 — 30% with fluxes determined from our spectra and fluxes reported by Peimbert & Torres-Peimbert (1992). We find Ha luminosities ranging from 1 x 1041 erg s" 1 to greater than 2 x 1042 erg s" 1 for the Bootes void galaxies (Table 1). No correction for extinction internal to the galaxies was included in the luminosity calculation. Extinction originating in our Galaxy is 1 — 2% for galaxies at these declinations and has been ignored. The luminosity of the brightest normal galaxies is approximately L(Ha+[NII]) = 1 x 1042 erg s" 1 (Kennicutt & Kent 1983) while the brightest starburst galaxies have L(Ha) = 3 x 1042 erg s" 1 (Balzano 1983). These values do not include correction for extinction internal to the galaxies, for comparison with our results. Thus we find that some of the Bootes void galaxies have very large Ha luminosities, although not larger than have been observed in similar galaxies in more densely populated regions of the universe. The galaxy with the largest Ha luminosity, 1510+4727 (Table 1), is not in the complete sample of IRAS galaxies in the direction of the Bootes void studied by Dey et al. (1990), suggesting its infrared emission is relatively weak. The Ha emission from this galaxy is strongly concentrated to the nucleus. There does not appear to be a companion or interaction, but the morphology in the off-band (continuum) image is peculiar. The next brightest galaxy in Ha is the one-armed spiral in the 1519+5050 pair. Star formation rates (SFRs) were calculated from the relationship given by Kennicutt

102

Weistrop: Star Formation in the Bootes Void

Galaxy 1406+4905 1428+5255 1432+5302 1444+4402 1446+4457 1510+4727 1519+5050 1519+5050 B 1530+4332 1530+4332 B 1537+5315 1540+5013

I(Ho)f 41.70 41.72 41.71 41.85 41.33 42.35 42.31 41.47 41.55 41.05 ? weak 41.16

SFRJ LINER 13 13 17 5 55 50 Sy 1 9 3

Note 1 Zw 81 X-ray luminous (Boiler & Dennefeld 1994)

4

t No absorption correction. X MQ yr" 1 for stars with 0.1 < M < 100 MQ. Assumes factor 2.8 absorption. TABLE 1. Ha Luminosities and star formation rates

(1983). In order to compare our SFRs with Kennicutt's, a correction for the internal extinction in the galaxies must be applied. Since we do not have spectra to estimate the internal extinction for most of the galaxies, we adopt the factor 2.8 determined by Kennicutt from the thermal radio emission of galaxies experiencing significant star formation. For comparison, Balzano (1983) finds a correction factor of 3.6 from the Ha/H/? ratio in the spectra of starburst galaxies. The calculated SFRs for stars of mass 0.1 — 100 MQ are shown in Table 1. For two galaxies for which the H a emission is due to nonthermal nuclear activity, the nature of that activity is indicated in the table. Ha observations of 1537+5315 were taken through thin cirrus clouds and are thus uncertain. For normal galaxies Kennicutt finds SFRs up to 2 0 M Q y r " 1 . Thus, several of the Bootes void galaxies have large SFRs, but only 1510+4727 and 1519+5050 have SFRs significantly higher than those found in field disk galaxies.

4. Distribution of galaxies within the Bootes void Models of voids in the large scale structure predict density enhancements within the void boundaries (van de Weygaert & van Kampen 1993; Dubinski et al. 1993). Such density peaks could provide suitable conditions for galaxy formation. The distribution of bright infrared galaxies within the Bootes void is shown in Figure 2. The location of the galaxies is given in a three-dimensional co-ordinate system determined from right ascension (6), declination (1 — ), and distance derived from the redshift. The origin of the coordinate system is at the center of the void, R.A. = 14h50m, Dec. = +46°, r = 310 Mpc (Kirshner et al. 1987). Figure 2 is obtained by rotating the entire galaxy distribution through angles 6 — 275° and = 40°. There appears to be structure in the distribution, in particular about half the galaxies are found within a vertical slab 10 Mpc wide. This structure is found in the distribution of all the galaxies known in the void. However, because these galaxies were identified using a variety of methods, and some of them are from incomplete samples, we show in Figure 2 only the galaxies found in the complete sample defined by Dey et al. The density of galaxies in the vertical slab is 3-4 times the galaxy density in the entire void and is comparable to the density of galaxies outside the Bootes void. This distribution may be consistent with the density enhancements in voids predicted by the models, or it may suggest that Bootes is not a

Weistrop: Star Formation in the Bootes Void

103

Y (Mpc)

-30

0 x 30 X (Mpc)

60

FIGURE 2. Distribution of a complete sample of bright infrared galaxie» * ^ * ^ £ ITat A similar distribution is found for all galaxies known to be in the void. The void center u> (0,0,0). single void Further studies of the distribution of the galaxies in the complete infrared sample may help resolve this issue. We are currently assessing the s t a t i c a l significance of the structure in Figure 2. 5. Summary Ha and broad-band imaging has been used to investigate the properties of galaxies in the Bootes void. Of the 16 galaxies observed thus far, at least ten are d1Sk systems^ Two interacting pairs are known, with a third likely. There may be other interacting and/or merging pairs within the sample. The frequency of interacting pairs may be high compared to the 8% found in emission-line galaxies (Salzer, MacAlpine, k Boroson 1989) aZough higher pair frequencies have been reported for luminous infrared and starburst galaxies (Gallimore & Keel 1993). Several of the galaxies are luminous in Ha, due in most cases to significant amounts of star formation. This is not surprising as the galaxies were identified from objective prism surveys and/or infrared detection by IRAS^ However, pencil-beam surveys found no normal galaxies in this region of space (Kirshner ^ T h e s e s are luminous, and in some cases undergoing star formation induced by interaction with other galaxies. These are not the types of galaxies predicted to be in V

°There appears to be structure in the spatial distribution of the galaxies with about half of them concentrated in a single slab. This structure may originate in the predicted density enhancements in voids. Results reported here were obtained in collaboration with R. Angione, I . Brown K , P. Cheng, S Cruzen, P. Hintzen, C. Hoopes, C. Liu, J. Lowenthal, R. Ohversen, L.

104

Weistrop: Star Formation in the Bootes Void

Sage, and B. Woodgate. The observations were obtained at the University of Arizona 90-inch telescope; KPNO 4-m telescope, part of NOAO, which is operated by AURA under a co-operative agreement with the National Science Foundation; Mt. Laguna Observatory 40-inch telescope, which is operated by San Diego State University; and the Lowell Observatory 31-inch telescope which, under an agreement with Northern Arizona University and the NURO Consortium, is operated 60% of the time as the National Undergraduate Research Observatory. Partial support was provided by the UNS Space Grant Consortium and a Barrick Faculty Development Grant.

REFERENCES BALZANO, V. A. 1983 Ap. J. 268, 602. BOLLER, T. &: DENNEFELD, R. 1994 This conference. BRAEMERD, T. G. & VILLUMSEN, J. V. 1992 Ap. J. 394, 409. DEY, A., STRAUSS, M. A. & HUCHRA, J. 1990 A. J. 99, 463. DUBINSKI, J., DA COSTA, L. N., GOLDWIRTH, D. S., LECAR, M. & PIRAN, T. 1993 Ap. J.

410, 458. GALLIMORE, J. F. & KEEL, W. C. 1993 A. J. 106, 1337. HOFFMAN, Y., SILK, J. & WYSE, R. F. G. 1992 Ap. J. Lett. 388, L13. KENNICUTT, JR., R. C. 1983 Ap. J. 272, 54. KENNICUTT, J R . , R. C. & KENT, S. M. 1983 A. J. 88, 1094. KIRSHNER, R. P., OEMLER, A., SCHECHTER, P. L. & SHECTMAN, S. A. 1981 Ap. J. Lett. 248,

L57. KIRSHNER, R. P., OEMLER, A., SCHECHTER, P. L. &; SHECTMAN, S. A. 1987 Ap. J. 314, 493. LACEY, C , GUIDERDONI, B., ROCCA-VOLMERANGE, B. & SILK, J. 1993 Ap. J. 402, 15. PEIMBERT, M. & TORRES-PEIMBERT, S. 1992 A. & A. 253, 349. SALZER,

J. J., MACALPINE, G. M. &

BOROSON,

T. A. 1989 Ap. J. Suppl. 70, 479.

VAN DE WEYGAERT, R. & VAN KAMPEN, E. 1993 M. JV. R. A. S. 263, 481. WEISTROP, D. 1989 A. J. 97, 357. WEISTROP,

D.,

HINTZEN,

P., KENNICUTT, R. C ,

LIU,

C,

LOWENTHAL,

OLIVERSEN, R. & WOODGATE, B. 1992. Ap. J. lett. 396, L23.

J., CHENG, K.-P.,

Physical Properties of Giant Extragalactic HII Regions By A. I. D I A Z Depto. de Fisica Teorica, Universidad Autonoma de Madrid, 28049-Cantoblanco, Spain A discussion of the determination of the "functional parameters" of Giant Extragalactic HII Regions is made for a wide range of physical conditions of the ionized gas. These functional parameters which are found to be the effective temperature of the ionizing radiation, the ionization parameter and the gas metallicity, are then traslated into physical properties of the ionizing star clusters - age, total mass and metal content - through the use of evolutionary models. These models combine evolutionary synthesis techniques with a well-tested photoionization code, CLOUDY, to obtain emission line intensities that can be compared directly with observations. From this comparison some conclusions about the general properties of the ionizing clusters can be extracted.

1. Introduction The term Giant Extragalactic HII Region (GEHR) usually refers to outstandingly large and luminous HII regions that are therefore easily observable in external galaxies, mainly on the discs of late-type spirals and in irregulars. Their size, characterized by mean diameters measured on Hcv images, can be as large as 1 kpc. Their Ha emission luminosity, which can reach 1041 erg s~', translates into a number of ionizing photons betbeen 1050 and 1053 s" 1 . This high number of ionizing photons cannot originate in a single star, which would provide about 1049 at most, but requires the presence of a young star cluster. In the few instances in which GEHR can be spatially resolved, observations reveal a core-halo structure with the core showing detailed substructure. However, in most cases these regions are observed as single objects. For example, for 30 Dor in the LMC Walborn (1991) distinguishes, in descending order of scale, the 30 Dor region, the 30 Dor nebula, the 30 Dor cluster and the cluster core; but, if placed at a distance of 10 Mpc, the whole 30 Dor region would be 20 arcsec in diameter and the 30 Dor cluster would be inside 1 arcsec. We can therefore say that probably any extragalactic HII region that can be observed from the ground is a GEHR and their physical properties such as mass of the ionized gas, mass of the ionizing cluster, stellar content, age etc., can be determined only indirectly from observations of the entire star-forming region. Photo-ionized gas shows a characteristic emission-line spectrum and, from this point of view, GEHR, as well as galactic HII regions, HII galaxies and even starburst nuclei, form a three-parameter family. The shape of the ionizing continuum, the degree of ionization and the chemical composition of the gas control the features of the emitted spectrum, with the gas density playing only a secondary role. Using simple photo-ionization models and under the assumption of a single ionizing star, the three parameters of a GEHR can be represented by what we call "functional parameters": the effective temperature of the ionizing star, ionization parameter and Z, the abundance by number of all elements heavier than helium, taken in solar proportions. These functional parameters can be related to physical properties of the star-forming regions by applying more reallistic models in which the ionization is provided by an evolving star cluster. 105

106

Diaz: Physical Properties of Giant Extragalaclic Ell Regions

Low-metallicity models

z/zQ

logf/ Ten (K) Z/ZQ

High-metallicity models

logU

Ten (K) logn«

TABLE 1.

0.05 -1.0 35000

0.10 -1.5 37500

0.20 -2.0 40000

-2.5 45000

1.0

1.7

2.0

2.3

-3.0 50000 2.7

0.5

55000

1.0

2.0

-2.0 35000

-2.5 37500

-3.0 40000

-3.5 45000

-4.0 50000

55000

1.0

1.4

1.7

1.9

2.0

2.2

2.3

Computed photo-ionization models

2. Functional parameters of GEHR We will call functional parameters of an HII region those which are determined with the use of photo-ionization models under certain simplifying assumptions. Most photoionization models in use solve for the ionization and thermal equilibrium of a sphere of gas ionized by a central radiation source. The integrated line emission from a number of elements is then calculated by adding the contributions of concentric shells with radii between some specified inner and outer values. The main assumptions underlying the determination of our functional parameters are: (a) the ionization is provided by a single main sequence star placed at the centre of a sphere; (b) the gas surrounding it is located at a given inner radius; and (c) the density of the gas is assumed to be constant. Under these assumptions the shape of the ionizing continuum is given by the effective temperature of the ionizing star, Tejj\ the degree of ionization is represented by the ionization parameter, which is a combination of the density, the filling factor and the ionizing flux; and the metallicity is usually represented by Z, the metal abundance, taken in solar proportions. We have used the photo-ionization code CLOUDY, kindly made available to us by G. Ferland, to compute an extensive grid of models to look for suitable indicators of the functional parameters over a wide range of physical conditions. In these models the ionizing stellar continua are taken from the atmosphere models by Mihalas (1972) for stars of solar metallicity. These models do not include the effects of line blanketing but, on the other hand, relax the assumption of LTE which we have considered to be more important for hot stars. We have also assumed a uniform chemical composition throughout the nebula. The effects of dust have been taken into account only through the depletion of the refractory elements Ca, Si, Fe and Mg by a factor of ten with respect to solar. Due to the different behaviour of models with subsolar, solar and higher metallicity we have divided our models in two metallicity regimes as given in Table 1 and we will dicuss the two sets of models separately.

3. Determination of the functional parameters 3.1. Low melaUicily regions The models we have computed assume a constant density throughout the nebula. Electron densities, which almost coincide with ionized hydrogen densities if the nebula is fully ionized, are easily calculated from optical spectra at moderate resolution by measuring the ratio of the red [SII] lines at 6717 A and 6731 A. The upper limit for the densities

Diaz: Physical Properties of Giant Exlragalactic Ell Regions

107

•3.25

FIGURE

1. The oxygen abundance, as calculated from model emission lines, versus input inetallicity, Z.

calculated in this way is around 50 cm known.

3

. We will assume the electron density to be

3.1.1. Determination ofL Since Z comprises the abundances of all elements heavier than helium, taken in solar proportions, it can be calculated just by determining the abundances of one of these elements. It is customary to use oxygen to represent the mean metallicity of the nebula. The abundance of oxygen is easy to determine from optical spectra since strong lines of [Oil] at 3727, 3729 A and [OIII] at 4959, 5007 A, the two dominant ionization states, are present. For the brighter and hotter HII regions, temperature-sensitive lines, weak in nature, like that of [OIII] at 4363 A are detectable and measurable, thus providing an accurate determination of the electron temperature, knowing this, the abundance of oxygen can be determined under the assumption that log(O/H)~ log(O+ + O + + )/H + . This assumption is well justified since in all the computed models the amounts of 0° and O 3+ are neggligible. Figure 1 shows the abundance of oxygen, O/H, calculated in this way from the emission lines resulting from the models versus the metallicity, Z, in units of the solar metallicity, ZQ, introduced as input in the same models. It can be seen that for metallicities Z < 0.2 Z 0 , the corresponding values of O/H are inside 30% and O/H abundances determined in this way seem to be a good indicator of Z. However, for metallicities as high as 0.5 Z o , the electron temperature is already below 10000 K, the limit of validity of the expressions commonly used for the derivation of oxygen ionic abundances. A better indicator for Z can be obtained by representing the electron temperature, as derived from the [OIII] A4363 A/[OIII] AA4959.5507 A ratio (e.g. Osterbrock 1989) versus the ratio [Oil] AA3727.3729 A + [OIII] AA4959.5007 A (Figure 2). On this plot, lines of constant Z are well defined. This kind of diagnostic seems to be independent of changes in the relative abundances of the elements as evidenced by the square symbols

Diaz: Physical Properties of Giant Extragalaclic Ell Regions

108

1

—

1

•

—i—

^ - " 0.05

. . . - • " 11.2 • *I

" * - ' • • \ •

.l-l.l-lI f l . l ' .

11.1(1

(IS5

0.70

• •

-.'••'"

*

Z=O.5Z..

K

' • '

0 85

1.00

I.I]

1.30

2. The electron temperature, as derived from the [OIII] A4363 A line, versus Z. Squares correspond to models with O/S and N/O higher than solar by a factor of 2.

FIGURE

on the plot, which correspond to models in which the ratio O/S and N/O are higher than solar by a factor of 2. 3.1.2. Determination of the ionizalion parameter, U In our models the ionization parameter, U, is defined as the density of the ratio of ionizing photons to the particle density

where Q(H) is the number of hydrogen-ionizing photons per second, nn is the total hydrogen density (ionized and neutral), c is the speed of light and R is the distance to the centre of the cloud. Since the ionization parameter is directly proportional to the degree of ionization, it is customary to use ratios of two lines of a given element corresponding to consecutive (or just different) ionization states in order to estimate U. To first order, this method should be independent of metallicity. [OII]/[OI1I] and [SI1]/[S1II] have been used by different authors but with only limited success since the hardness of the ionizing continuum also affects these ratios in a significant way. Lines corresponding to a single ionization state have also been used as ionization parameter indicator. In particular, the easily measurable quotient [SII] AA6717+6731 A/Ha has been proposed by Dopita k Evans (1986). This ratio is strongly dependent on metallicity, as can be seen in Figure 3a), and so it should not be used as an independent indicator; however, once the metallicity is known, it can provide relatively accurate values of U. Figures 3b, 4a and 4b show the sensitivity of the relation between the ratio [SII] AA6717+6731 A/Ha and (/ to the different parameters: the hydrogen particle density, the effective temperature of the ionizing radiation and the metallicity. It is easily seen that, at a fixed metallicity, both Tejj and 7iH add a small scatter to the relation. Taking

Diaz: Physical Properties of Giant Exiragalactic HII Regions

109

into account the typical error quoted for [SII] line intensities, logt/ can be determined by this method within ±0.2 dex. 3.1.3. Determination of the effective temperature, Teff Once both the metallicity and ionization parameter are known, the effective temperature of the ionizing radiation can be estimated by looking at the intensities of the [OIII] AA4959,5007 A lines. Figure 5 shows the relation between the ratio of the sum of these lines to H^ and Te// for regions of 100 cm" 3 particle density and Z=0.2 Z 0 . Taking into account the typical errors in line measurements and the uncertainty in the determination of U quoted above, T e // can be estimated only to within an accuracy of ±1000 K. It should also be recalled that these values of Tejj correspond to specific atmosphere models and therefore do not constitute an absolute scale. However, the method outlined can be used to rank in effective temperature different HII regions whose density, metallicity and ionization parameters are known. 3.2. High-metallicity regions Regions with rnetallicities higher than 0.5 ZQ are usually referred to as high-metallicity HII regions. They require more elaborate methods for abundance analysis. The low electron temperatures of the gas in these regions due to the effectiveness of the cooling through the emission lines of the different elements, makes the detection of the weak temperature sensitive lines almost impossible. Empirical calibrations have been proposed in the past in order to overcome this difficulty. Among them, the one involving the sum of the optical [Oil] and [OIII] lines, is probably the most popular. Originally proposed by Pagel, Edmunds & Smith (1980), it relies on the fact that the ratio R23 = [on]AA3727+3729;+jom]AA4959+5007; d e c r e a s e s with increasing metallicity due to the more efficient cooling of the nebula, being relatively insensitive to geometrical factors affecting the degree of ionization. The actual relation between i?23 and oxygen abundance is calibrated empirically at the low-metallicity end, where abundances can be derived directly (see the previous section), and requires the use of theoretical photoinization models to define the high-metallicity end. The calibration has undergone several revisions as better data and more refined photoinization models have become available. For low abundances the Edmunds & Pagel (1984) and Skillrnan (1989) recalibrations almost coincide. However, for high abundances different model sequences give different results depending on the assumptions underlying the models. In fact, this behaviour should be expected, since in this case the assumption that the cooling is dominated by the oxygen abundance through its optical lines is no longer true. At the low electron temperatures of these regions, IR lines become important and other elements besides oxygen also start to contribute substantially to the cooling. Therefore, for moderate-rnetallicity regions (i.e. those with metallicity between 0.5 and 1.0 ZQ) the empirical calibration can be used, bearing in mind that oxygen abundances derived in this way are accurate to ±0.2 dex. For higher-metallicity regions, it is necessary to determine the other two functional parameters, ionization and stellar effective temperature, in order to estimate the global metal content of the region. 3.2.1. Determination of the ionization parameter, U The sulphur ionic ratio S + / S + + was proposed by Matins (1985) as an ionization parameter indicator. For the ionization parameter range usually found in high-metallicity regions, the line ratio [SII] AA6717+6731 A / [SIM] AA9069+9532 A can equally be used. Figure 6 shows the logarithmic relation between this ratio and (/. A good correlation

Diaz: Physical Properties of Giant Extragalactic HII Regions

110 1

1

r-—i

p—i

1—- , i — i — i — i -

•

(a)

i

i i

A

C

z = 0.05 Zo

•

s

P-"

= 0.50 Z o

z = 0.20 Z Q z = o.io z 0

4

i

CD

z

•

I

1

•

i

S -

-

t

!9 C3 []

cD

i

-2.5

—3

-3.5

-1.5

-2

-0.5

-1

log u

•

(b)

T

,

„ = 40000 K

Z = 0-2 Z o

CO

a

logn H = 1.7

A

logn H = 2.0

0

logn H = 2.3

•

logn H = 2.7

\

- _ .

-3.5

.

•

.

1

-3

•

1

.

.

.

.

1

-1.5

-2.5

1

-0.5

log u

FIGURE 3. a) The [S1I]/H« ratio versus ionization parameter for all computed models, and b) the effect of the density on the relation for models with T e // = 40000 K and metallicity 0.2 solar.

111

Diaz: Physical Properties of Giant Extragalactic HII Regions

•

(a)

z=

0.2 Z o

logn „

X — co

= 2

•

.

T,.,, = 55000

to r,.,, = 50000

\

tp^ CO

•

r e(( = 45000

0

Ttfl = 40000

»

-

T t| , = 37500 T.,, = 35000

1

-3

-3.5

.

.

.

.

i

-2.5

.

-2

.

.

.

i

-1.5

-1

-0.5

log u

—i

1 1

—r •,

1

,

r

,

,

1

r— i — r — i — . — | — i —

'

(b)

CO

'

elf = 40000 K l°gn H = 2

T

d \

s \

0.2

17, 6731,

X

-

\

\ \

\

\

,

. z = 0.50 Z o

\

\

d

A Z

= 0.10 Z o

D Z

= 0.05 Z o

\

1

-3.5

o z = 0.20 Z o

\

-3

•

i

-2.5

i

i

, , i ;—v i :

-2

-1.5

-1

-0.5

log u FIGURE 4. a) Same as Figure 3 but showing the efTect of varying the stellar effective temperature on models of 0.2 solar metallicity and 100 c m ' 3 particle density and b) same as Figure 3 but showing the effect of the metallicity for models o(TeJj = 40000 K and 100 c m ' J particle density.

112

Diaz: Physical Properties of Giant Exlragalactic HII Regions

6xlO4 FIGURE

FIGURE

5.5xlO4

5xlO4

4.5xlO4

4xlO4

3.5xlO4

3xlO4

5. The [OIIIJ/H/? ratio versus stellar effective temperature for models of different ionization parameters, as labelled, and 100 cm" 3 particle density.

6. Relation between ionization parameter and the [SII]/[SIII] ratio. Open and solid symbols refer to solar and twice-solar models, respectively.

exists for logf/ < —2.5, independently of metallicity and effective temperature. Typical errors of logf/ determined in this way are of the order of ±0.1 dex. 3.2.2. Determination o/Teff and Z Vilchez k Pagel (1988) have suggested the use of a second parameter besides R23 in order to select the right photoionization model for a given HII region. This second

Diaz: Physical Properties of Giant Extragalactic H1I Regions ,

!

,

,

,

-2.5 -3.0 a

\

,

r-—1

1

1—1—

—i

-4.0

ns •

\

1

-3.5

a

\

1

» n

V

•

35000 K

•

37500 K

Z

-

.

1

0 45000 K * 50000 K • 55000 K

>

m c>

o

.

40000 K

*-0

°

1

113

O

\V — .

"

in

b l

1

.

.

.

.

1

1.5

0.5 log 7)'

FIGURE 7. The abundance parameter, 7223, as a function of »/ for models of solar (upper part) and twice-solar (lower part) metallicity and different effective temperatures and ionization parameters, as labelled.

parameter, i), is related to the shape of the ionizing radiation and involves the quotients of different ionic fractions of oxygen and sulphur. ?/ is not very sensitive to the electron temperature and an equivalent observational parameter, ;/ defined as - _ [OII]AA3727 + 3729A/[OIII]AA4959 + 5007A '' ~ [SlI]AA6717 + 673lA/[SIII]AA69069 can equally be used (Diaz et al. 1991). In a plot of #23 versus ij (Figure 7) it is possible to see how these two observables respond to changes in the three functional parameters. Solar and twice-solar models occupy the upper and lower parts of the diagram respectively evidencing the dependence of #23 on metallicity. Some dependence of #23 on effective temperature can also be seen, #23 being higher for higher values of Tejj, but #23 is almost independent of the ionization parameter, as was originally suggested. The dependence of #23 on Tejj is stronger for the highest metallicity varying by about an order of magnitude from Tejj = 35000 K to Tejf = 55000 K, the whole range of temperatures considered, for twice-solar models. On the other hand, ?/ depends very strongly on effective temperature; but it also depends on the ionization parameter to a large extent. However, it is almost independent of global metallicity, Z. Therefore, once the ionization parameter is found by the means outlined in the previous section, the position of a given H1I region on the #23 versus tj diagram can provide both its metallicity and the effective temperature of the radiation responsible for its ionization. The models represented in Figure 9 correspond to a constant particle density of 10 cm" 3 . An increase of the density from 10 to 100 cm" 3 does not affect the values of rj but increases the values of #23 by a factor of about 2 for the highest-metallicity models. This would translate into an underestimate of the global abundances. Observationally, it is not possible to discriminate among densities lower than 50 cm" 3 .

114

Diaz: Physical Properties of Giant Extragalactic HII Regions

According to our calculations, models with 10 cm" 3 particle density have R23 values higher than models with 50 cm" 3 particle density by a factor of 1.5 for an effective temperature of 37500 K. This is of the same order as the typical errors quoted for the intensities of the oxygen lines for regions of high metallicity. For higher effective temperatures the effect is even lower. Therefore, the errors introduced when deriving metallicities for regions of densities lower than 50 cm" 3 are, in most cases, within observational errors and using the lowest density models provides a metallicity estimate at the low metallicity limit.

4. Physical parameters of GEHRs The methods outlined so far allow the determination of the functional parameters both for low-metallicity and for high-metallicity HII regions and, if applied in a systematic way to homogeneous sets of data, they can provide very valuable information about fundamental relationships. Some physical properties of the regions can also be obtained by combining the derived parameters with complementary information. For instance, if Ha luminosities and sizes are known for the HII regions, this information combined with the derived ionization parameter can yield the actual value of the particle density, the filling factor and the total mass of ionized hydrogen (e.g. Diaz el al. 1991). However, in order to find information about other physical properties the assumption of a single ionizing star has to be relaxed. In fact, the number of photons ionizing GEHRs, as determined from observed Ha fluxes, indicates that they are ionized by star clusters. These clusters need to be young in order to keep the region ionized, but can be sufficiently evolved as to present significant differences in their spectra when compared to main sequence stellar models. For unevolved, zero-age clusters, the functional parameter Tejj probably denotes the temperature of the most massive, and therefore the hottest, stars in the cluster. This is not the case for older clusters in which evolved stars can actually be hotter than the most massive stars still on the main sequence and can dominate the ionization of the nebulae. In order to explore the properties of the ionizing clusters, we have computed evolutionary models according to the following guidelines: (a) The clusters are assumed to form in a single burst with a power-law initial mass function (IMF). (b) The stars in the clusters are evolved along their corresponding evolutionary tracks depending on their initial mass and metallicity. (c) The star clusters are evolved until the number of ionizing photons they supply are insufficient to maintain the nebula ionized, without the inclusion of supernova explosion effects. (d) For each evolutionary stage , the emergent spectrum is synthesized by calculating the number of stars at each point of the HR diagram. We have assigned to this point the closest atmosphere model in effective temperature and gravity. This spectrum has then been scaled to the luminosity of the theoretical star. The emergent spectra resulting from the computation are subsequently used as input for the photoionization code CLOUDY (Ferland 1990) and the emission-line spectra of the ionized regions are obtained as described in Section 2. Models of this kind have been computed by Garcia-Vargas & Diaz (1994) for solar and twice-solar metallicity GEHRs. The main results of this work are: 1. When the cluster evolves, its spectrum becomes harder at about 3.5 to 4 Myr due to the appearance of WR stars, if they are present. Nevertheless, the corresponding optical

Dfaz: Physical Properties of Giant Extragalaclic Ell Regions

115

3 5 Myr

1.5 Uyr

2 Myr

. .

,

0.2E5 Mo 0.6E5 Mo

2.5 Myr

*

1

1

-0.5

I

0.5 Log

FIGURE 8. Same as Fig. 7 but for star clusters as ionizing source. Only clusters with 2 and 6 x 104 Ma are shown. The cluster ages are labelled on the curve corresponding to the cluster of 6 x 10 MQ at twice solar metallicity.

spectra are almost identical for all ages, except around 4 Myr when the contribution of red supergiants becomes conspicuous. 2. For a given cluster, the number of Lyman ionizing photons, Q(H), does not change much during its evolution. 3. The ratio of Q(H) to the number of He ionizing photons, Q(He), changes appreciably with age, having a maximun at about 3.5 to 4.5 Myr. 4. For a given age, clusters with higher total mass have higher luminosities. 5. For a given IMF, the emission-line spectrum is controlled by: (a) the ionization parameter, which changes mainly with the total mass of the cluster; (b) the shape of the ionizing continuum, which changes mainly with the age of the cluster; and (c) the metallicity of the gas. The model results concerning emission lines can be examined by constructing diagnostic diagrams. Figure 8 shows the same diagram presented in Figure 6 but for star clusters of different total masses formed with a Salpeter IMF and lower and upper mass limits of 0.85 and 120 M0respectively. It can be seen that the evolutionary history of the clusters reproduces the diagram rather well and an equivalence between functional and physical parameters can be easily found. The ionization parameter and the stellar effective temperature relate to total mass and age of the cluster, respectively. Other diagnostic diagrams are reproduced equally well (see Garcia-Vargas, Bressan k Diaz in this volume). Therefore, this kind of diagram can be used to estimate the total mass of the ionizing cluster and its evolutionary state. The total masses of the ionizing clusters corresponding to the models with functional parameters in the range derived from observations, are between 2-6 xlO4 M©, for an IMF as above. These models have equivalent widths of H/? (EW(H/?) between 64 and 610 A. However, most of the observed GEHRs have EW(H/?) between 5 and 90 A. This fact points to the presence of an underlying stellar population contributing substantially

116

Diaz: Physical Properties of Giant Extragalaclic HII Regions

to the continuum at H/?. Such a composite population should be studied in more detail with the help of population synthesis techniques. On the other hand, as mentioned before, the few spatially resolved GEHRs show several emitting regions with very different spectra (Skillman 1985; Diaz et al. 1987). The different emitting regions may be different gas clouds with different physical properties, ionized by the central cluster or the ionizing cluster itself could consist of smaller subclusters of different ages, the older ones providing the bulk of the continuum light and the younger ones dominating the ionization. Both alternatives deserve further study. I would like to thank J. Iglesias and R. Jimenez for their work in the computation of the low and high metallicity models respectively and M.L.Garcia-Vargas with whom a great part of this work has been made. This work has been partially supported by the DGICYT through project PB90/0182 and by NATO, through grant CRG920198 for collaborative research.

REFERENCES DIAZ, A. I., TERLEVICH, E., PAGEL, B. E.J ., VILCHEZ, J. M. & EDMUNDS, M. G. 1987

Mon.

Not. R. astr. Soc. 226, 19. DIAZ, A. I., TERLEVICH, E., VILCHEZ, J. M., PAGEL, B. E. J. & EDMUNDS, M. G. 1991

Mon.

Not. R. astr. Soc. 253, 245. DOPITA, M. A. & EVANS, I. N. 1986 Asttophys. J. 307, 431. EDMUNDS, M. G. & PAGEL, B. E. J. 1984 Mon. Not. R. astr. Soc 211, 507. FERLAND, G. 1990 HAZY, a brief introduction to CLOUDY. GARCIA-VARGAS, M. L., BRESSON, A. & DI'AZ, A. I. 1994 This volume. GARCIA-VARGAS, M. L. & DIAZ, A. I. 1994 Astrophys. J. Suppl. In press. MATHIS, J. S. 1985 Astrophys. J. 291, 247. MIHALAS, D. 1972 Non-LTE Model Atmospheres for B and 0 stars NCARTN/STR-76. OsTERBROCK, D. 1989 Astrophysics of Gaseous Nebulae and Active Galactic Mic/e»Mill Valley, Univ. Science Books. PAGEL, B. E. J.,

EDMUNDS, M. G.,

BLACKWELL, D. E., CHUN, M. S. & SMITH, G.

1979

Mon. Not. ft. astr. Soc. 189, 95. SKILLMAN, E. D. 1985 Astrophys. J. 290, 449. SKILLMAN, E. D. 1989 Astrophys. J. 347, 883. VILCHEZ, J. M. & PAGEL, B. E. J. 1988 Mon. Not. R. astr. Soc. 231, 257.

N. R. 1991 Massive Stars in Starbursts (Ed. C. Leitherer, N.R. Walborn & T.M. Heckman), p. 145. Cambridge Univ. Press.

WALBORN,

The Giant HII Region NGC 2363f ByROSA GONZALEZ-DELGADO1,ENRIQUE PEREZ1, GUILLERMO TENORIO-TAGLE1, JOSE M. VICHEZ1, ELENA TERLEVICH2, ROBERTO J. TERLEVICH2, EDUARDO TELLES2, JOSE M. RODRIGUEZESPINOSA1, MIGUEL MAS-HESSE4, MARIA LUISA GARCIA-VARGAS3, ANGELES I. DIAZ3, JORDI CEPA ANDHECTOR CASTANEDA1 1

Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain

2

Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, UK

3

Depto. Fisica Teorica CIX, Universidad Autonoma, Cantoblanco, 28049 Madrid, Spain

"Laboratorio de Astrofisica Espacial y Fisica Fundamental (LAEFF), Apdo. 50727, 28080 Madrid, Spain We present narrow-band Ha imaging and long-slit optical and near infrared spectroscopy of the giant HII region NGC 2363. We have found broad emission lines at 4686 A and at 5810 A attributed to WC stars at 6 arcsec to the East of the brightest core of the region. We confirm the existence of low-intensity broad components in Ha and [OIH] which extend some 500 pc. We have derived the physical conditions and chemical composition of the gas in 15 different zones in the region, and do not find significant variations in the abundances. The Paschen discontinuity has been found in emission. The Pa electron temperatures obtained are significantly smaller than those obtained from the [OIII] and [SIII] emission lines, indicating the presence of large temperature fluctuations.

1. Introduction One of the targets of the GEFE programme is the giant HII region NGC 2363 located in the SW of the irregular galaxy NGC 2366. This is one of the largest extragalactic HII regions with high surface brightness. The object was observed in La Palma in narrow band Ha and long-slit spectrophotometry from [Oil] A3227 to [SIII] A9532 at two positions, at the brighest core of the region (which we call knot A) and at 6 arcsec to the East (knot B).

2. Narrow-band images The object was observed with the 1-m JKT telescope. We used a CCD with a spatial scale of 0.3 arcsec pixel" 1 . Images through interference filters, FWHM = 50 A, centred at 6563 A for Ha and at 6832 A for the continuum were obtained. Figure l a shows the resulting Ha image after coaddition by using the Richardson-Lucy algorithm. We can see a very low surface brightness emission in the North-East direction. The region is resolved into four main components. These include knot A, resolved into a compact t The data presented in this contribution are part of the GEFE collaboration. GEFE, Grupo de Estudios de Formacion Estelar, is an international collaboration of astronomers from Spain, the U. K., France, Germany, Denmark and Italy, formed to take advantage of the international time granted by the Comite Cienti'fico International at the Observatories in the Canary Islands. 117

118

Gonzalez-Delgado et al.: The Giant HII Region NGC 2363

120

140

ISO

ISO

FIGURE 1. (a) Ha image after coadding using the Richardson-Lucy algorithm for image deconvolution, (b) The internal structure of NGC 2363 is expanded in panel b.

bright core and a fainter component 1.3 arcsec to the South, knot B 6 arcsec to the East with a head-tail structure, and an additional component 3 arcsec East of knot A. The total Ha flux measured after correcting for reddening is 8.7 10~ 12 erg s" 1 cm " 2 which translates to an Ha luminosity of 1.4 1040 erg s" 1 and an ionizing photon luminosity of 1.5 1052 photon s" 1 . We have calculated some physical parameters for the region. The average electron density is 5 cm" 3 , the total ionized gas is 2.4 105 M© and the total mass of the cluster 3.4 105 M Q .

3. Long-slit spectroscopy The spectroscopic observations were obtained with the 4.2-m WHT telescope using the blue and the red arms of the ISIS spectrograph and an EEV CCD in each arm. The dispersion was 1.4 A pixel" 1 , and spatial sampling of 0.33 arcsec pixel" 1 and 0.66 arcsec pixel" 1 . The object was observed at three position angles, in knot A, at p.a. = 187°, and in knot B, at p.a. = 185° and p.a. = 130°. At both positions, we can see more than twenty lines of the hydrogen Paschen series (Figure lb), including the series limit, a prominent nebular Hell A4686 and low-intensity

Gonzalez-Delgado et al.: The Giant HII Region NGC 2363

A B

te([OIII ]) 1.61 1.52

U{[ SIII])1 1.71 1.59

<e([OII]) 1.45

*.(l;sir]) 1.34 1.54

TABLE 1. te is in units of 10 4 K,

C(H/J) 0.20 0.16

O/H

S/H

N/H

7.86 7.89

6.14 6.22

6.35 6.31

119

He/H 0.0806 0.0749

abundances are 12+log(X/H)

broad components in Ha and [OIII]. And at the knot B also WR bumps at 4686 A and 5810 A. Both knots present high-excitation spectra. 3.1. Spatial distribution of the emission lines The distribution of the emission lines Ha, [OIII], and [SIII] are quite similar. The continuum and the emission lines peak at the same position for knot A; however, for knot B, the peak of the continuum is shifted 1.2-1.5 arcsec to the North with respect to the peak of the emission lines (Figure 2a). The excitation [OIII]/H/3, and ratio [OIII]/[OII] show a steep fall in the South direction and a shallower decrease to the North (Figure 2a). 3.2. Physical conditions and chemical composition We have calculated the electron temperature, density, reddening and chemical compositions for 15 different zones. We have not found any significant variations in these parameters. The region is characterized by an electron density of 190 cm" 3 , and electron temperature of about 16100 K for knot A and 15200 K for knot B, a very low amount of reddening, C(K/3) ~ 0.2, and very low abundances (Table 1). 3.3. The broad low-intensity components At both knots, we have detected a broad low-intensity component under Ha and [OIII] A5007,4959. This component was previously detected by Roy et al. (1992) in a direction nearly perpendicular to our positions. This broad component is extended over about 500 pc; its flux is about 3% of the narrow, and the FWHM is about 30 A. The existence of this high-velocity gas can be accounted for following the hydrodynamics of breakout (Tenorio-Tagle and GEFE 1994). 3.4. WR features and the nebular Hell emission In the central part of knot B, we have detected wide bumps in emission at 4660 A and at 5810 A. These features are attributed to WC stars. The presence of these stars gives an age for this knot of between 3 and 4 Myr. Recent evolutionary models indicate that, at low metallicity, stars with mass larger than 60 MQ become WC stars (Fagotto et al. 1994). From the calculated luminosities (10 36 8 and 10 3 6 5 erg s" 1 for the bumps at 4660 A and 5810 A respectively) and using the calibration given by Smith (1991) only one WC would be enough to explain simultaneously the luminosities in these two bumps. At both knots nebular Hell A4686 emission was detected, being more extended and brighter in knot B than in knot A. The ratio of the luminosity of this line to H/3 is 0.005 for knot A and 0.015-0.035 for knot B. These ratios indicate that the ionizing spectrum is very hard, harder in knot B than in knot A. However, the excitation [OIII]/H/? is lower in knot B than in knot A, which means that the presence of WR stars in a cluster increases the number of photons with energy higher than 54 eV considerably but do not change the Lyman continuum photons much.

Gonzalez-Delgado et al.: The Giant HII Region NGC 2363

120

3900

5200

6500 7800 wavelength (A)

25

9100

I

r

Knot B

20 CM

E

o 'w

15

£> a> CO

b

10

X 3

LL

5 ^W*^ '.

I

. . .

3900 FIGURE

I

I

5200

.

.

I

6500 7800 wavelength (A)

9100

2. (a) Merged spectrum for knot A. (b) Merged spectrum for knot B.

Gonzalez-Delgado et al.: The Giant HII Region NGC 2363

1

M eg

"uE

i

•

•

•

i

•

•

•

i

•

•

121

•

continuum

6

400

f

5

d) (0

300

4 X

E 3

Knot B

3

200 <3

2

o

C

100

w

^

0 -8

- 6 - 4

-2 0 2 distance (arcsec)

8

FIGURE 3. WR bumps at 4660 A and at 5810 A for knot B.

3.5. Temperaturefluctuationsinside the region We have observed the Paschen discontinuity in emission at both positions. Using the ratio APa/H/? we have derived the electron temperature in the region. The values obtained are lower (13000 K and 10200 K for knot A and B, respectively) than the value derived using collisional lines. This indicates the presence of large temperature fluctuations within the nebula; these fluctuactions affect the emissivity of the ionic lines and consequently the derived chemical composition. To compute this effect, we have obtained the average temperature, To, and the r.m.s. temperature fluctuation, t%, proposed by Peimbert (1967). We have obtained 14600 K and 12200 K for To, and 0.064 and 0.098, for knot A and B respectively. The effect in the He abundance is quite small, less than 6%, however the O/H abundance is increased in 0.2 dex and 0.43 dex for knot A and B respectively. This has an important influence in the determination of the primordial helium abundance, since the increase of the O/H abundance and the decrease of He/H turn the slope of the y vs. O/H regression to steeper values.

122

Gonzalez-Delgado et al.: The Giant HII Region NGC 2363

4. Conclusion The physical properties of the region indicate that NGC 2363 is undergoing a violent burst of recent star formation, where WR stars, strong nebular Hell A4686 emission and high velocity gas are present. We gratefully acknowledge partial support from DGICYT Grant PB91-0531(GEFE) and the NATO Grant CRG920198 for collaborative research.

REFERENCES FAGOTTO,

F.,

BRESSAN,

A.,

BERTELLI,

G. &

CHIOSI,

C. 1994 Astr. Astrophys. Supp. Ser. In

press. M. 1967 Astrophys. J. 150, 825. ROY, J.-R., AUBE, M., MCCALL, M. &; DUFOUR, R.J. 1992 Astrophys. J. 386, 498. SMITH, L.F. 1991 In IAU Symp. no. 143: Wolf-Rayet stars and interrelations with the massive stars in galaxies (ed. Karel A. van der Hucht & Bambang Hidayat), p.601. Kluwer. TENORIO-TAGLE, G. & G.E.F.E. 1994. In preparation. PEIMBERT,

Photometric Diagrams of NGC 2366 By A. APARICIO1, J. CEPA 1 , H. O. C. CASTANEDA 1 , C. GALLARTSC. MUNOZ-TUNON1, E. TELLES2 AND G. TENORIO-TAGLE1 1

Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain

2

Royal Greenwich Observatory, Madingley Road CB3 OEZ, Cambridge, UK

The irregular magellanic galaxy NGC 2366 is usually assumed to belong to the M 81/ NGC 2403 group. (B — V)-V and (V — R)-V photometric diagrams of its stellar content are presented. Using its brightest blue and red stars, its distance is estimated to be about 3.0 Mpc.

1. Introduction For some time now, a considerable effort has been devoted to the analysis of the stellar content of nearby galaxies through the photometry of their resolved stars. These data, combined with spectroscopic and radio observations, provide interesting information about the history of the star formation and about the star formation processes in galaxies. The understanding of these processes is relevant, since they define the path followed by the galaxy through its evolution. Our project is devoted to the analysis of the stellar content in nearby galaxies and is included in the GEFEf project. We present the first results of our photometric analysis of NGC 2366. This galaxy is usually assumed to belong to the M 81-NGC 2403 group. Table 1, summarizes the global parameters of NGC 2366.

2. Data and results Observations of the stellar content of NGC 2366 were carried out in 1992 using the 1150x1250 EEV5 chip at the prime focus of the 2.5-m Isaac Newton Telescope of the Observatory of Roque de los Muchachos in La Palma (Canary Islands, Spain). The limiting magnitudes, are about B = 22.1, V = 23.2 and R = 22.9 (defined as those for which 50% of the stars are lost). The colour-magnitude diagrams of NGC 2366 are shown in Figure 1. The field stars have been statistically removed from the diagrams. The Figure shows the typical pattern of a galaxy with active star formation. The colour-magnitude diagrams are characterized by the presence of two plumes separated by a scarcely populated Hertzsprung gap. The blue one is predominantly populated by main sequence and blue-loop stars, and the red one is populated by red supergiants. No Cepheid stars have been observed in this galaxy, and therefore only indirect estimates of its distance modulus have been obtained. Using the method of the brightest blue and red supergiants in the galaxy, we estimate its distance modulus to be about (m — M)o = 27.4, or about 3Mpc, in good agreement t GEFE, Grupo de Estudios de Formation Estelar, is an international collaboration of astronomers from Spain, the UK, France, Germany, Denmark and Italy, formed to take advantage of the international time granted by the Comite Cienti'fico International at the observatories in the Canary Islands 123

Aparicio et al.: Photometric Diagrams of NGC 2366

124

18 -

—

20

•' •. . *

22

•

24

•

••

- • » • • ' • » • .

.

.

•

~

• • ^ S i f c x •"••••••..

* •

,

.

-1

. 1 . . . . 1" 1 . . . .

0

1

2-0.5

(B-V)

1. . . . 1. . . . 1. . . . 1

0

0.5 (V-R)

1

1.5

FIGURE 1. The V vs. (B — V) and V vs. {V — R) colour-magnitude diagrams of the stars of NGC 2366 after correction for contamination by foreground stars.

with the determinations of other authors, which range from 27.07 to 27.65 (Tikhonov et al. 1991; Sandage k Tammann 1974; de Vaucouleurs 1978; Tully 1987). The full analysis of the data presented here will be submitted to one of the major journals. We gratefully acknowledge financial support for this project from NATO Grant CRG920198 for collaborative research and DGICYT Grant PB91-0531 (GEFE).

REFERENCES G. 1978 Astrophys. J. 224, 710. A. & TAMMANN, G. A. 1974 Astrophys. J. 191, 603. TIKHONOV, N. A., BILKDMA, B. I., KARACHENTSEV, I. D. & GEORGIEV,

DE VAUCOULEURS, SANDAGE,

TS.

B. 1991 Astron.

Astrophys. Suppl. 89, 1. TULLY,

R. B. 1987 Nearby Galaxies Catalog. Cambridge University Press, Cambridge.

Spectroscopical Imaging of Star-Forming Regions ByJ. M. MAS-HESSE 1 , C. MUNO Z-TUNON 2 , J. M. VILCHEZ 2 , H. O. CASTANEDA 2 AND D. CARTER 3 'Laboratorio de Astrofisica Espacial y Ffsica Fundamental, POB 50727, E-28080 Madrid, Spain 2 Instituto de Astrofisica de Canarias, E-38200 La Laguna, Spain 3 Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, UK As part of the GEFEf collaboration, observations of star-forming regions with high spectral resolution and long-slit sampling are being undertaken. 2D maps of physical parameters like density, excitation, extinction...etc. have been produced with l" spatial resolution and 2" spatial sampling. Some preliminary results on the giant HII Region NGC 5471 and the irregular galaxy NGC 4214 are presented. Very high velocity components have been detected at some particular positions on the nebulae, as well as other peculiar kinematical structures (redshifted secondary emission peaks, line splitting etc.). The whole emitting area of NGC 5471 behaves as a unique entity with respect to excitation, with no correlation with the emitting knots. On the other hand, differentiated star-forming regions can be identified on NGC 4214. Finally, the distribution of dust particles seems to be rather inhomogeneous and anticorrelated with the distribution of emission-line intensities.

1. Introduction: aim and targets The ultimate aim of the GEFE collaboration is to determine which are the physical parameters that control the formation of a violent burst of star formation. Within this framework and in order to fulfil this main objective it is important to know the physical properties of star-forming regions with high enough spatial resolution as to determine variations of the measurable parameters within the emitting nebulae. We aim to use measurements of age, excitation degree, velocity dispersion and chemical composition to know whether we are dealing with single star-forming regions or with well differentiated physical entities within a patch of ionized gas, which cause misinterpretation in our understanding of the "physical object" (Mufioz-Tunon et al. 1993). The observational programme consists of mapping starburst regions with high spectral resolution. With this approach it is possible to map the parameters and also to differentiate the kinematical features that allow us to identify those regions in the nebulae where the effects of massive stars are relevant. Both physical parameters such as density, chemical composition, temperature etc. and kinematics are observationally independent tracers of violent star formation. The idea of mapping ionized regions using classical long-slit spectroscopy at different positions in a nebula is becoming more and more popular nowadays (Arsenault and Hippelein, private communication; Chu & Kennicutt 1994, among others). The procedure we are using to map the regions is explained in the next section. f GEFE, Grupo de Estudios de Formation Estelar, is an international collaboration of astronomers from Spain, the U. K., France, Germany, Denmark and Italy, formed to take advantage of the "International Time" granted by the Comite Cientffico Internacional at the Observatories in the Canary Islands. 125

126

Mas-Hesse et al.: Spectroscopical Imaging of Star-Forming Regions

A better understanding of the physics of the ISM is of course also part of our goal and interesting results can be obtained with our obvervational approach. We want to study the spatial distribution of interstellar dust particles, its relation with the gas density distribution and their effects on both continuum and nebular extinction. Finally, by combining kinematical features such as jets, bubbles, high velocity gas etc. with other physical parameters we will derive interesting information about the mixing process and the chemical enrichment of the ISM. Peculiar features, such as very broad components within giant HII regions, have already been reported in the literature, in particular in NGC 2366 (see Gonzalez-Delgado et al. 1994 and references therein). No physical mechanism that could account for these observed features has been found up to now. In future work we will proceed to characterize the physical properties of this high-velocity gas, as a way to determine plausible accelerating mechanisms. The present data base allows us to follow this approach. The observational sample of GEFE includes Giant HII regions, dwarf galaxies and starbursts in spirals and irregular galaxies. The high-resolution mapping is however restricted to the highest surface brightness and most extended objects. We have already mapped the giant ionized regions in M101, M33, NGC 4214, NGC 2403 and NGC 6822. In this paper we will present some preliminary results concerning two of the targets, NGC 4214 and NGC 5471.

2. Observational technique and data processing Observations were taken with the 4.2-m William Herschel telescope (WHT) at the Observatorio del Roque de los Muchachos (La Palma, Spain) using the double-arm ISIS spectrograph. Two spectral ranges are achievable simultaneously for each slit position. The selected ranges were from 4700 A to 5050 A in the blue arm and from 6400 A to 6800 A in the red one; with a spectral resolution of around 0.33 A pixel"1. The slit, 180" long and 1" wide, was located at several positions (typically 12) along the object sampling the emission at 2" steps. Typical seeing values were about 1" and spatial pixel size along the slit was 0'-'35. The TWEAK procedure at the WHT together with the pointing and tracking accuracy of this telescope proved to be good enough for this observational technique to be successfully employed. In Figure la the observational scheme is sketched and in Figure lb we present one of the possible 2D maps that can be built from data provided by this technique. The Ha intensity map presented in this figure can be compared with that taken for NGC 5471 (see Munoz-Tunon, this volume) with the TAURUS Fabry-Perot. Although the spatial resolution is clearly worse, the produced image is well in accordance with that obtained with TAURUS, proving the validity of the employed technique. The data have been processed using standard IRAF routines, including the STSDAS tables package. The output parameters derived from the different measurements were used to build up the corresponding 2D images and maps. The uniform sampling of each region is important to perform a global interpretation.

3. Preliminary Results We discuss the results that can only be obtained with our observational technique, but not with TAURUS. This means that a detailed discussion about velocity field, velocity dispersion maps and any other feature also achievable using 2D TAURUS mapping will be analyzed in another communication. As mentioned above, we will restrict the discussion

Mas-Hesse et al.: Spectroscopical Imaging of Star-Forming Regions

127

FIGURE 1. a) Sketch of the employed long-slit scanning procedure; b) NGC 5471 Ha emission

line intensity map. North is right and east is up.

to two of the targets: NGC 5471, one of the classical Giant HII regions in M101, and NGC 4214, an irregular galaxy showing numerous star-forming knots.

3.1. NGC 5471 We show in Figure 2 the electron density and excitation maps for NGC 5471. The electron density has been parameterized by the [SII] ratio, so that the isocontour values correspond indeed to line ratio values. The excitation parameter has been estimated, as usual, by the flux ratio between the [OIII] and H/? emission lines. The most remarkable results from these maps are the following: 1. The electron density distribution is not correlated with emitting knots. 2. The excitation does not show any clear sub-structure, such as those found in the distribution of the emission-line intensities. Therefore, although different emitting knots are resolved, as shown in Figure lb, in terms of excitation NGC 5471 behaves as a single entity. This result is consistent with that obtained from a purely kinematical approach (see Mufioz-Tunon, this volume). In Figure 3a we present an extinction map of this region derived from the distribution of the F(RQ)/F(R0) ratio. It is clear that the interstellar extinction is anticorrelated with the Ha and continuum emitting peaks. The observational technique we have used allows not only to produce global maps of the region, but also to analyze the profile of the emission lines with high spectral and spatial resolution. In this sense we have found, for example, high velocity gas in certain, relatively small areas of the emitting regions.

Mas-Hesse et a/.: Spectroscopical Imaging of Star-Forming Regions

128

FIGURE

2. a) Density map F([SII] A6716 A)/F([SII] A6731 A) and b) excitation map F([OIH] A5007 A)/F(H/?) for NGC 5471.

1.50E-15F

1.25E-15 •

1.00E-I5 -

7.5OE-16 -

5.00E-16

2.5K-16 -4000

FIGURE

-2000

0 2000 Velocity Ika/il

4000

3. a) Extinction map of NGC 5471 based on i r (Ho)/F(H/?) and b) spectrum of region A showing the broad emission component.

Figure 3b shows the Ha line profile of region A. An underlying broad component with FWHM of about 2000 km s" 1 is clearly detected, confirming the results of Castaneda et al. (1990). The origin of this broad emission-line component is not yet well understood and will be analyzed in detail in the future. 3.2. NGC 4214

NGC 4214 is an irregular galaxy located at 6.5 Mpc. Several star-forming knots are clearly visible in the Ha emission map shown in Figure 4. An important WR population is known to be present in some of the knots (Sargent k Filippenko 1991; Mas-Hesse k Kunth 1991) and from evolutionary population synthesis it is known that the brightest starburst is around 4-5 Myr old (Mas-Hesse k Kunth 1991). As in the case of NGC 5471,

Mas-Hesse et ai: Spectroscopical Imaging of Star-Forming Regions

129

FIGURE 4. Ho intensity map of NGC 4214.

FIGURE 5. NGC 4214 Ha surface diagram at position A. Dispersion on the X-axis; spatial extension on the Y-axis. Note the emission component redshifted by about 80 km s" 1 and extended over 3".

we have also found significantly higher values of the reddening in NGC 4214 outside the regions where the emission lines are brighter, possibly indicating that the dust particles are not uniformly distributed. On the other hand, the distribution of the excitation parameter shows well differentiated regions, contrary to the evenness found in NGC 5471. Analysis of the line profiles has led to the detection of several interesting features. We summarize here the three most remarkable ones: 1. A redshifted (around 80 km s" 1 ) emission component extended over 3" (Figure 5), just on the side of a region of very high interstellar absorption. 2. A three-peak split line, apparently associated with a bubble with a diameter of around 85 pc and an expansion velocity close to 80 km s" 1 (Figure 6a). It could be related to one of the WR stars which are known to be present in this object.

Mas-Hesse et al.: Spectroscopical Imaging of Star-Forming Regions

130

1

1

1

-

3.00E-16

4.006-16 3.50E-I6

2.50E-16 3.00E-16 2.00E-16 2.50E-16

U\ V \

1.50E-16

1.00E-16

-

5.00E-17

i

-200

i

0 Velocity Ika/il

-

-

2.00E-16 1.50E-16 1.00E-16

i

1

200

400

5.00E-17

6540

6550

6560 6570 6580 6590 Wivalingth (angstroasl

6600

FIGURE 6. a) Ha line splitting detected at position B. Note that the centroid of the split components is clearly redshifted with respect to the central component; b) Ha broad component (FWHM around 550 km s" 1 ) in region C, compared with the emission-line profile in a nearby region.

3. A broad Ha line with FWHM around 550 km s" 1 whose origin is still unknown to us (Figure 6b). We are grateful to the CCI (Comite Cientifico Internacional de los Observatorios de Canarias) for allocating the 5% of the International Time to the GEFE collaboration. It has always been very nice to enjoy lively discussions with the rest of the GEFE team and in particular to feel the enthusiasm and support from Guillermo Tenorio-Tagle as leader of the project. The WHT is operated on the island of La Palma by the RGO at the Observatorio del Roque de los Muchachos of the IAC. The Observatory staff has always provided excellent support. This work has been financed by DGICYT grant No. PB910531 (GEFE), the NATO Grant CRG920198 for collaborative research and the IAC. JMMH acknowledges the hospitality of the IAC where most of this work was carried out.

REFERENCES H. O., VILCHEZ, J. M. & COPETTI, M. V. F. 1990 Astrophys. J. 365, 164. CHU, Y. H. & KENNICUTT, R. C. 1994 Astrophys. J. In press. CASTANEDA,

GONZALEZ-DELGADO, R. M., PEREZ E., TENORIO-TAGLE, G., VILCHEZ, J. M., TERLEVICH, E., TERLEVICH, R., GARCIA-VARGAS, M. L. & DIAZ, A. 1994 Astrophys. J. Submitted.

J. M. & KUNTH, D. 1991 In Wolf-Rayet Stars and interrelations with other Massive Stars in Galaxies (ed. K. A. van der Hucht & B. Hiyadat), p. 613. MUNOZ-TUNON, C. 1994 Supersonic Motions in Giant HII Regions. This volume. MUNOZ-TUNON, C , VILCHEZ, J. M. &; CASTANEDA, H. O. 1993 Astion. & Astiophys. 278, 364. SARGENT, W. L. W. & FILIPPENKO, A. V. 1991 Astron. J. 102, 107. MAS-HESSE,

A Study on the HII Regions of NGC 4449 By O. FUENTES-MASIP, H. O. CASTANEDA AND C. MUNOZ-TUNON Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain Observations in the line of Ha and of [OIII] A5007 of the galaxy NGC 4449 taken at the William Herschel Telescope of the Observatorio del Roque de los Muchachos with TAURUS-2, a FabryPerot interferometer, are being used to study the correlation between the diameter orfluxand the velocity dispersion of its HII regions. Two different catalogues of HII regions are being compared. In the first we consider each flux relative maximum as a differentiated HII region, while in the second sample we use kinematical criteria to identify the different regions.

1. Introduction In 1981, Terlevich and Melnick established a correlation between the diameter or the luminosity of giant HII regions and the turbulent width of their emission lines, which represented a chance to improve present extragalactic distance estimations. To study the possible existence of similar correlations for HII regions in a single galaxy, we have chosen NGC 4449, a giant irregular type I galaxy. It is located 5 Mpc away from the Milky Way and is very rich in HII regions. New criteria to identify HII regions are explored and compared with the samples and measurements already present in the literature. Detailed information will be presented in Fuentes-Masip et al. (1994, in preparation). 2. Data analysis The product of the calibrated Fabry-Perot observations was a three dimensional set where the X-Y axes were the spatial coordinates, and the Z axis was wavelength calibrated (the dispersion direction). The analysis of the data was done with MATADORf. A first photometric analysis of the cubes was done by collapsing the cubes in the spectral direction. This produced flux maps where we applied FOCAS (a program from the IRAF package), characterizing an HII region as an area on the image being a relative maximum of flux above a threshold (three times the sigma of the local background), with an area greater than the seeing disk. Secondly, radial-velocity and velocity-dispersion maps were produced by fitting the emission lines of the data cubes. For each of these maps we obtained its associated gradient map. In this way we identified every HII region on a kinematical basis, as the area with an almost constant radial velocity and velocity dispersion (varying only by about 0.1 km s" 1 pixel"1), differentiated from its neighbours by an abrupt change in these magnitudes (of about 1 km s" 1 pixel"1). Isolated HII regions were limited by high values of these gradients because of the random values of the radial velocity and velocity dispersion in the areas devoid of emission. We extracted and fitted the integrated spectrum of each HII region, correcting the velocity dispersion (
Fuentes-Masip et al.: A Study on the HII Regions of NGC 4449

132

1.2

1.4

1.6

Log a (Km/s) FIGURE 1. Log(radius [pc]) vs. log(corrected velocity dispersion [km s—1]), for all HII regions with single Gaussian profiles found by FOCAS in the Ha flux maps. Bigger circles indicate brighter regions. The dashed line represents the correlation by Terlevich and Melnick (1981) and the dotted line, the one by Melnick et al. (1987).

3. Discussion In Figure 1 our HII region radii are, on average, smaller than expected from the results of Terlevich and Melnick (1981) and Melnick et al. (1987). This could be due to the diffuse luminosity that permeates the main body of NGC 4449, and which constitutes the real background of the galaxy. Because of this, radii are easily underestimated, since FOCAS detects, in our case, only what is above three times the sigma of the local background. It is also possible that other authors overestimated some HII region sizes due to visual merging of independent objects. All the HII regions detected present supersonic velocity dispersions. The diffuse luminosity of NGC 4449 could be responsible again, preventing the detection of the smaller and fainter normal HII regions which have subsonic velocity dispersions. On the other hand, perhaps we are measuring the profiles of some regions as single Gaussians while they are split profiles of stellar wind-driven bubbles in the poorly-defined disk of this irregular galaxy. That would artificially broaden the emission lines. In a similar work, Arsenault et al. (1990) observed in NGC 4321 that only the HII regions with highest surface brightnesses displayed the expected correlations. Following their idea, we have plotted in the example of Figure 1, bigger circles for regions with higher surface brightness. The radii and corrected velocity dispersions of the 6 brightest regions follow the relation log -R[pc] = 3.01ogcr[km s~ x ]-2.6. The slope of this correlation is very close to the value of 3.68 found by Melnick et al. (1987).

REFERENCES ARSENAULT,

R.,

FUENTES-MASIP,

ROY,

O.,

J.-R. &

J. 1990 Astron. Astrophys. 234, 23. H. O. &; MUNOZ-TUNON, C. 1994 In preparation.

BOULESTEIX,

CASTANEDA,

MELNICK, J., MOLES, M., TERLEVICH, R. & GARCIA-PELAYO, J.-M.

849. TERLEVICH, R. & MELNICK, J. 1981 M. N. R. A. S. 195,

839.

1987 M. N. R. A. S.

226,

Metallicity Effects on the Properties of Very Young Star Clusters By M. L. GARCIA-VARGAS 1 A. BRESSAN 2 AND A. I. DIAZ 1 J

Dep. de Fisica Teorica, Universidad Autonoma de Madrid, 28049-Cantoblanco, Spain

2

Osservatorio Astronomico di Padova, Vicolo dell' Osservatorio 5, 35122-Padua, Italy

We present theoretical evolutionary models for star clusters at ages between 1.0 and 5.5 Myr and different metallicities. Clusters at these ages can provide enough ionizing photons to explain the observations of Giant Extragalactic HII Regions (GEHRs). The emergent ionizing continua from the clusters are used as input for the photoionization code CLOUDY to obtain the corresponding emission line spectra.

1. Ionizing star clusters The ionizing clusters have been assumed to form in a single burst with a standard power-law IMF, (j>(m) = m~a, with a — 2.35 (Salpeter), considering the masses of formed stars between the lower limit, rniow = 0.85 M©, and the upper limit, mup = 120 M©. Models are computed for total cluster masses, My = J"m "p m(rn)dm, between 1.2 x 104 M© and 6 x 104 M©. Clusters in this range of masses can provide the necessary ionizing photons, Q(H), to produce the observed H a luminosities (logQ(H) m 51) of GEHRs. The clusters are evolved along the evolutionary tracks of Bressan et al. (1993) for solar metallicity (Z = 0.02) and Fagotto et al. (1994) for the rest of them (Z=0.05, 0.008, 0.004, 0.001). The tracks are computed taking into account the new OPAL opacities (Iglesias, Rogers &; Wilson 1992), mass loss (scaled with the square root of the metallicity) and moderate core overshooting. These tracks are not corrected for the effect of the thickness of the wind, which means that Tejf refers to the hydrostatic core. Isochrones for different ages between 1.0 and 5.5 Myr have been computed. Figures la-d show the isochrones at 2-5 Myr respectively for a cluster of 6 x 104 M©. Different symbols correspond to isochrones of different metallicities, as labelled in the plot. 2. The Wolf-Rayet and red-supergiant phases Special attention is paid to some important evolutionary stages in these young clusters, like the Wolf-Rayet and red supergiant phases, which can provide important observational clues about the age and the metallicity of the star clusters responsible for the ionization of the gas in GEHRs. The presence of Wolf-Rayet stars can be detected through some features in the optical emission line spectrum, since their hard ionizing spectrum can control the excitation of the regions. On the other hand, red supergiants affect the V — K colour and, if they are present, a strong calcium triplet in the infrared (AA8498, 8542, 8662 A), with equivalent widths greater than 9 A should be detected. Cervirio and Mas-Hesse (1994) have studied the metallicity effects on the WR and RSG phases, using the evolutionary tracks of Schaerer et al. (1993) and Charbonnel et al. (1993), finding that there are less WR features at low abundances since the WR phase is shorter when the metallicity decreases, and a lack of RSGs at low metallicity. In the present models, the WR phase has a maximum at solar metallicity although the phase is seen in burts with metallicities between 0.2 Z© and 2.5 Z@. On the other hand, the RSG phase exists 133

134

Garcia-Vargas et al.: Metallicily effects... Figure la. (2 Myr)

Figure lb. (3 Myr)

CO

J

J bO

ao

. Z=0.050 . Z=0.020

• Z=0.050 • Z=0.020 * Z=0.008 o Z=0.004 A Z=0.001

3

* Z=0.008 o Z=0.004

* Z=0.001

10°

10°

5x10* *eff

(K)

'eff

Figure lc. (4 Myr)

CO

5x10* (K)

Figure Id. (5 Myr)

•

\

-

J 60 O

• Z=0.050 . Z=0.020 *Z=0.008 oZ=0.004 AZ=0.001

o -

OS

\

V1

• Z=0.050 . Z=0.020 • z=o.ooe

3

1. \

oZ=0.004 »Z=0.001

1

1.5xlO6

10° T BW (K)

5x10*

2xlO6

1.5xlO5

106

5x10*

T f i f f (K)

FIGURE 1. lsochrones at 2 Myr (la), 3 Myr (lb), 4 Myr (lc) and 5 Myr (Id) for a 6 x 104 M 0 cluster. Different symbols correspond to isochrones of different metallicities, as labelled in the plot.

Garcia-Vargas et al.: Metallicity

effects...

135

Z/Z 0

fWR

MWR M$£ tRSG

2.50 1.00 0.40 0.20 0.02

1.75 1.91 2.60 2.82

TABLE

1. WR and RSG phases at different metallicities.

1.99 2.54 1.55 1.30

40 40 40 40 _

6.52 5.70 6.20 9.90 14.6

7,.6 13.0 14.3 11.0 5.8

20 30 30 20 15

at metallicities as low as 1/50 Z 0 although the phase is shorter at low metallicity. Table 1 gives the times at which the WR phase begins (in Myr), tlWR, the length of the phase in Myr, AtwR , and the minimum initial mass of a star to enter the phase, M^J^; the same parameters are given for the RSG phase: t'RSG, AIRSG and MRgQ, respectively.

3. Spectral energy distributions For each of the clusters in each evolutionary stage, we have synthesized the emergent spectrum by calculating the number of stars in each point of the HR diagram and assigning to it the most adequate stellar atmosphere model (the closest one in effective temperature and surface gravity). The corresponding stellar spectrum has then been scaled to the luminosity of the corresponding theoretical star in the HR diagram. The atmosphere models we have used are those of Clegg & Middlemass (1987) for stars with Teff > 50000 K (corresponding to the last evolutionary stages of massive stars) and those of Kurucz (1992) for stars with 3500 K < Teff < 50000 K. We have chosen these models for two reasons: first, because there is a good agreement between Kurucz's and Clegg & Middlemass' models of a common T e /;, which provides a consistent temperature scale, and second, because Kurucz's grid includes models in a wide range of metallicities, between 10~5 Z 0 and 10 Z 0 , and this fact allows us to build a consistent grid of models, in which the effect of the metallicity is taken into account not only in the evolutionary tracks but also in the atmosphere models. Figures 2a and 2b show the synthesized spectra of a cluster of 6 x 104 M 0 with the "standard" IMF (a=2.35, m,ou,=O.85M0, m up =120M o ) at 2 and 4 Myr respectively. Different lines correspond to SEDs for clusters at different metallicities (Z = 0.001, 0.004, 0.008, 0.02 and 0.05), as labelled in the figure. The most remarkable characteristic of the spectra is the ionizing continuum beyond 1 Ryd, which becomes harder when the extreme WR appear, and this fact depends on both the age and the metallicity of the cluster, as discussed in the previous section. Since the number of Lyman ionizing photons per solar mass of the cluster, Q m (H), is controlled by massive O-B main sequence stars, this number does not change much during the evolution between 1.0 and 5.5 Myr, or with the metallicity. This fact implies that the intensities of recombination emission lines of the gas like those of the Balmer series will be mainly dominated by the mass of the cluster, but not by its age or metallicity. This effect is showed in Figure 3a, in which we have represented the number of ionizing photons per solar mass of the cluster versus the age. Only in the case of Z — 0.05 (2.5 solar) is the behaviour of this parameter clearly different from the rest of models. However, the ratio of the number of helium ionizing photons, Q(He) (y >1.8 Ryd), to total ionizing photons, Q(H) (v > 1 Ryd) changes appreciably, both with the age and the metallicity of the cluster, since <3(He) is dominated by the extreme WR stars, and

Garcia-Vargas ei al: Metallicily effects...

136

Figure 2a .Metallicity effect in the SED at 2 Myr 1

i

' _ IUE range

' ..Optical _

12 fim

C IV

Ug II

1.0 Ryd 1.8 Ryd 2 6 Ryd

'"""•"'

"v Aj

if-

-^

•

s 00

•&?:.. z

o

— —

= z= z= = z= i

-2

•, X1

0.050 0.020 0.008 0.004 0.001 .

.

-1.5

1

I

\

\ .

.

1

.

1. 1 . . . 0.5

.

-0.5

-1

log,, Figure 2b .Metallicity effect in the SED at 4 Myr ' ..Optical ., IjirnH, Hf

12 fim

^ IUE rang^ Ug n

C IV

1

1.0 Ryd 1.8 Ryd 2 6 Ryd

#?•-

_ _ ^'l

\

o

— z — z -•- z z -2

= = = =

'j :. • • • ! S o "

0.050 0.020 0.008 0.004

\

\ ',

= 0.001

-1.5

-0.5

-1

0.5


FIGURE 2. Synthesized spectra of a 6 x 10 M 0 cluster at 2 Myr (Figure 2a) and 4 Myr (Figure 2b). In the upper part of the plot it has been marked the different wavelenghs ranges: 12-/im IRAS band, optical range (1/iin -3200 A), WE range (3200-1012 A), and the points of the ionization potentials for Hydrogen (1 Ryd), Helium (1.8 Ryd), and Oxygen (2.6 Ryd).

the metallicity of the cluster controls the beginning and the duration of this evolutionary stage. In other works (Mas-Hesse k Kunth 1991; Garcia-Vargas k Diaz 1994) it has been defined an equivalent ionizing temperature,Ttq, for each cluster as the T e // of an ionizing star whose spectrum has the same value of Q(He)/Q(U). Figure 3b shows the evolution of Teq with the age of the cluster at different metallicities. For ages less than 3 Myr a trend is observed in the sense of having lower T e? at higher metallicities as expected. For ages greater than 3 Myr, when WR stars form part of the cluster, this trend disappears and the situation becomes more complicated since the value of Teq is very model dependent; in particular it depends on the chosen atmosphere models.

Garcia-Vargas ei al: Metalliciiy effects...

1

1

t'

1

137

'

•

0-

«-,.... ^o 1 ^ -

....

.

.... z ... z , _

o-

Z

= o.ooi

= = Z = z =

°- ...y

0.004 0.008 0.020 0.050

^

\

0 1

•

LJ

Age

FIGURE

3. (a) Number of ionizing photons per mass unit of the cluster versus the age. (b) Evolution of Teq of the ionizing cluster at different metallicities.

4. Photoionization models The cluster ionizing spectra computed as described in the previous section have been used as input for the photionization code CLOUDY (Ferland 1990). The resulting emission-line spectra of the ionized regions have been computed under certain hyphotheses about the geometry and the physical and chemical conditions of the gas. A static spherical geometry is assumed in the models, considering the star cluster in the centre of the sphere, and the gas surrounding it, located in a shell at an inner radius R from the ionizing source. Constant density throughout the nebula has been assumed for simplicity. The detailed results of the nebular emission-line spectrum for models with a wide range of parameters are going to be published in a forthcoming paper. The model results concerning emission lines can be more easily examined by means of diagnostic diagrams. As a first result, we have applied the models to high metallicity GEHRs {Z > ZQ), although the application to low metallicity GEHRs is in preparation. Figure 4 shows an excitation diagram in which the value of log ([OIIIj/H^) is plotted versus log ([OII]/[OIII]). Different types of open symbols represent the evolution of different masses clusters, as labelled in the figure. These points, representative of the ages of 1.5, 2, 2.5, 3, 3.5, 4, 4.5, 5, 5.1, 5.2, 5.3 and 5.4 Myr are joined with a solid line ( Z 0 ) or a dashed line (2Z@). The observational data for extragalactic low-excitation (high-metallicity) HII regions are from the compilation of Garcia-Vargas k Diaz (1994), and new data from Oey &; Kennicutt (1993) and Zaritski, Kennicutt & Huchra (1994). They are plotted as filled symbols as labelled in the figure. The models cover the observations of GEHRs, and in principle it should be possible to determine the physical parameters of the star cluster responsible for the ionization (mass, age and metallicity), by fitting the emission line spectrum of the ionized gas. In a previous paper (Garcia-Vargas & Diaz 1994) models for GEHRs at solar metallicity were presented using Maeder &. Meynet (1989) evolutionary tracks, and the implications for the IMF and the stellar evolution derived from the comparison between models and observations were discussed. We found that models predicted high-metallicity high excitation regions, which have not been observed. This apparent contradiction seems to

Garcia-Vargas et al.: Melallicity

138

effects...

•

-

A A

0.6e5 Mo

A 5t~~~^-S^ ^ o o

o

-

71

-

^ 0 . 2 e 5 Mo

3e5 Mo

\ * \

k Kennnicutt o• Oey Oey & Kennnicutt A .*

^ & ^ •/^ M \^**-

A

. -

-0.5

i

0

.

.

.

.

i

v

N v* v

1993 o 1993. Z > 2 Z o Zaritsky, Kennnicutt Ic Huchra 1994 Other high metallicity HII regions

CM

FIGURE

° - 4 e 5 MQ

b

.

.

^ x

^

^

^

^

it

'A

\ \ h b A i

if .

.

.

0.5 1 log ([OII]3727/[0III]4959+5007)

.

1.5

4. Diagnostic diagram. Comparison between our new grid of evolutionary models for high metallicity CEHRs and observational data.

go away in these new models by using the new evolutionary tracks from Bressan et al. (1993), and the new atmosphere models. We therefore think they will eventually allow us to determine the physical parameters associated with the star clusters responsible for the ionization of the gas: mass, age and metallicity.

BRESSAN,

A.,

FAGOTTO,

REFERENCES F., BERTELLI, G. & CHIOSI, C. 1993 Astr. Astrophys. Suppl 100,

647. M. & MAS-HESSE, J. M. 1994 Astr. Astrophys. Submitted. C , MEYNET, G., MAEDER, A., SCHALLER G. & SCHAERER, D. 1993 Astr. Astrophys. 279, 338.

CERVINO,

CHARBONNEL,

CLEGG, R. E. S. & MIDDLEMASS D. 1987 Mon. Not. R. astr. Soc. 228, 759.

F., BRESSAN, A., , BERTELLI, G. & CHIOSI, C. 1994 Astr. Astrophys. In press. FERLAND, G. J. 1990 HAZY a brief introduction to CLOUDY V.76.03, June, 1990. GARCIA-VARGAS, M. L. & DIAZ, A. I. 1994 Astrophys. J. Suppl. Accepted. IGLESIAS, C. A., ROGERS, F. J.& WILSON, B. G. 1992 Astrophys. J. 397, 717. KURUCZ, R. 1992 Precision Photometry: Astrophysics of the Galaxy, (ed. P. G. Davis Philip, A. R. Upgren & K. A. Janes). L. Davis Press. MAEDER, A. & MEYNET, G. 1989 Astr. Astrophys. 210, 55. MAS-HESSE, J. M. & KUNTH, D. 1991 Astr. Astrophys. Suppl. 88, 399. FAGOTTO,

OEY, M. S. & KENNICUTT, R. C. 1993 Astrophys. J. 411, 137. SCHAERER, D., CHARBONNEL, C , MEYNET, G., MAEDER, A. & SCHALLER, G. 1993 Astr.

Astrophys. 280, 346. D., KENNICUTT, R. C. &

ZARITSKI,

HUCHRA,

J. P. 1994 Astrophys. J. Submitted.

The Prototype Starburst Galaxy NGC 7714: Physical Conditions of the Gas and the Stellar Population! By ROSA GONZALEZ-DELGADO1, ENRIQUE PEREZ1, MARIA LUISA GARCIA-VARGAS2, ELENA TERLEVICH3, ROBERTO J. TERLEVICH AND JOSE M. VICHEZ1 'institute) de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain 2

Depto. Fisica Teorica CIX, Universidad Autonoma, Cantoblanco, 28049 Madrid, Spain 3

Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, UK

We present narrow-band Ho imaging and long-slit optical and near-infrared spectroscopy of the starburst galaxy NGC 7714. We have detected WR stars in the starburst region, which indicate an age for the burst of between 3 and 5 Myr. We have obtained the physical condition of the gas in the starburst region and in three HII regions. These have moderately low abundances, while the nucleus has half solar abundance, with an overabundance of N.

1. Introduction A typical starburst galaxy can be defined as a spiral galaxy with a bright nucleus bluer than expected for its morphological type, which emits strong narrow emission lines similar to low-ionization HII region spectra, as a consecuence of the photoionization by the ultraviolet radiation of hot stars, with typical Ha luminosities ranging from 10 40 to 10 42 erg s" 1 . During this intense recent burst of star formation between 107 and 10 10 MQ of massive stars are formed within a radius of a few hundred pc about its nucleus. NGC 7714, the prototype of the starburst (henceforth SB) galaxies (Weedman et al, 1981) and classified as a SBb peculiar, is in interaction with the irregular galaxy NGC 7715. The X-ray luminosity (6 10 40 erg s" 1 ) is explained with about 104 supernova remmants in a volume of 280 pc radius (Weedman et a/.1981). The 6-cm radio map shows a weak double radio structure separated by about 1 arcsec at p.a = 30°. The object was observed in La Palma in narrow band Ha imaging and in long slit spectrophotometry from 3700 to 9700 A at p.a. = 110° and 216°) across the nucleus. 2. N a r r o w b a n d i m a g e s The object was observed with the 1-m JKT telescope. We used an EEV CCD with a spatial scale of 0.3 arcsec pixel" 1 . Images through interference niters, FWHM = 50 A centred at 6607 A for Ha and at 6925 A for the continuum were obtained. Figure 1 shows the Ha image after the continuum was removed from the on-band image. Several morphological features can been seen in the continuum image. There is a bar at p.a. ~ 143°, from the end of which two spiral arms emerge, and a very distorted disc, showing a loop and two tails, one to the South-West and a longer one to the East, where the t The data presented in this contribution are part of the GEFE collaboration. GEFE, Grupo de Estudios de Formacion Estelar, is an international collaboration of astronomers from Spain, the U.K., France, Germany, Denmark and Italy, formed to take advantage of the international time granted by the Comite Cientifico Internacional at the Observatories in the Canary Islands. 139

140

Gonzalez-Delgado et al.: The starburst galaxy NGC77J4 400

250 -

200

250

FIGURE

300

350

1. Ha image.

companion galaxy is found at 2 arcmin (23 kpc) from the nucleus of NGC 7714. The Ha image reveals extended emission clearly attributable to a nuclear starburst. Giant HII regions appear to be located in the spiral arms, at less than 2.5 kpc from the centre of the galaxy, and very weak HII regions are also found at 69 arcsec from the nucleus in the bridge connecting with the companion galaxy. 2.1. Physical parameters and the luminosity function We have measured the total Ha flux of the galaxy and of the individual 23 detected HII regions. The observed flux was corrected for the extinction deduced from our spectroscopic data, and also for the contribution of the [Nil] A6548, which falls within the filter pass band. The total Ha luminosity is 10 4176 erg s" 1 , which is 1.4 times the luminosity in the circumnuclear SB; the luminosity per unit area is equal 4.8 1034 erg s" 1 pc~ 2 . The SB region comprise 73% of the total mass of the ionizing stars in the galaxy; however, the ionized gas only represents 31% of the total ionized gas. So most of the star formation occurs in the circumnuclear region, making the SFR/area higher in the central zone. The Ha luminosities of the HII regions range from 10 386 to 10 406 erg s" 1 . The faintest regions can be ionized by several O5V stars; however, the brighter regions would need the equivalent of several hundred O5V stars. 74% have luminosities larger than 1039 erg s" 1 . Following Kennicutt et a/.(1989), we have fitted a power law to the cumulative luminosity function and the resulting power is -1.67, which is near to the mean value, -2, obtained by Kennicutt et a/.(1989) for a sample of irregular and spiral galaxies. 3. Long-slit spectroscopy The spectroscopic observations were made with the 4.2-m WHT telescope using the blue and the red arms of the ISIS spectrograph and an EEV CCD in each arm. The dispersion was 1.4 Apixel"1, with a spatial sampling of 0.66 arcsec pixel"1. The object was observed at p.a. = 110° and 216° across the nucleus. Along the slit we have observed the starburst region and three HII regions, located at 6 arcsec to the SE (region A), 12

Gonzalez-Delgado et al.: The starburst galaxy NGC7714

141

600

^ •

^ E

500

J3

400

i

I 300 T o 20° CO

"x uZ

100 0

4000

5000

100

I

6000 7000 8000 wavelength (A) I

I

I

I

I

9000 10000

I

1

E

80

(0

60

Nucleus

o •

'b

40

Flux(

co

20

Hu, f 4000

5000

6000

7000

8000

9000 10000

wavelength (A)

FIGURE

2. Merged spectrum for the nucleus.

arcsec to NW (region B) and 18 arcsec (region C) to the SW. The spectra are typical of intermediate-excitation HII regions. The Balmer lines in the nucleus are clearly affected by absorption features associated with the underlying stellar population. The synthetic spectrum (given by Bica 1988) which best matches the continuum and the absorption spectrum is S7, a combination of an average of spiral galaxies dominated by young stellar populations (22% with age < 50 Myr). We have found the Ca II triplet in absorption, and broad Hell A4686 emission attributed to WR stars. 3.1. Spatial distribution of the emission lines The spatial distribution of Ha, [OIII] and [SIII] are quite similar, the maximum emission being offset 0.5-1 arcsec to the West from the maximum in the continuum distribution. The excitation ratio [OIII]/H/J and the ratio [OIII]/[OII] present their local maxima at

Gonzalez-Delgado et al.: The starburst galaxy NGC7714

142

5

1

'

,

,

|

.

1

.

OIII]/[OII]

4

•

• - O - -

p-a = 110°

LOIII]/HP

P'

•

3

•i

/ .©€

' • • * * • «

1

•©'

•m

0

•

-

CD

-

2

•

0

I

-15

i

, , 1 . .

-10

. 1 . , . . 1

-5

1

1

/r

'to-£ "

/

-^* 1 i

0 5 distance (arcsec)

10

15

20

-* [Nll]/Ha O---[SII]/Ha a - - 2 [Ol]/Ha

-15

FIGURE

-10

-5

3. (a) Spatial distribution of the ratios [OIII]/H/?, and [OIII]/[OII]. (b) Spatial distribution of the ratios [SII]/Ha, [OI]/Ho and [NI]/Ha.

the HII regions and at 1 arcsec NW of the nucleus. The [SII]/Ha and [0I]/Ha show a similar behaviour, with a local minimum at the same positions where [OIII]/H/? and [OIII]/[OII] reach local maxima. This anticorrelation between low and high ionization emission lines is not observed in the case of [Nil]/Ha, since this ratio presents a maximum in the nucleus. This could indicate either the presence of an additional source of excitation or, more likely, an overabundance of nitrogen, since the ratio [NII]/[OII] is almost flat, except in the nucleus where there is a maximum. 3.2. Physical conditions and chemical composition The electron density is determined from the ratio [SII] A6717/A6732. The HII regions (B and C) have low density, ~ 100 cm"3, which remains constant across the regions; however,

Gonzalez-Delgado et al.: The starburst

Nucleus

A B C TABLE 1.

OIII]|) 1.14 1.28 1.11 1.02

<e([SIH]) 0.75 0.80 1.36 1.20

Nil]) 0 .76

*e([SII]) 0.72 0.71 1.11

galaxy NGC7714

Ne 550 164 109 96

O/H 8.5 8.47 8.19 8.27

4

143

S/H

N/O

6.95 6.89 6.40 6.46

-0.68 -0.85 -1.04 -1.18

te is in units of 10 K, the abundances of O/H and S/H are abundances of N/O and Ne/O are log(X/Y)

Ne/O -0.83 -0.72 -0.71 -0.65

12+log(X/H), the

a very steep increase in the density toward the center is observed in the circumnuclear region, reaching a m a x i m u m of 650 c m " 3 at the nucleus. T h e density profile is flatter t h a n the r~2 law, the best fit is obtained with the function n = with n0 - 653 and a = 2.6. T h e electron temperature was obtained from the line ratios of [OIII], [SHI], [SII], and [Nil]. In the circumnuclear region [OIII] A4363 is in the absorption wing of H7, which makes its measured flux very uncertain. We have adopted i e ([SIII]) for the nucleus and region A, and < e ([OIII]) for the regions B and C. T h e HII regions have moderately low metallicity, and the nucleus has a half solar abundance, with an overabundance of N, which can be attributed to pollution by winds from the W R stars detected. 3.3. Stellar population

in the starburst

region

W R features are detected in the nucleus at 4680 A in the form of a double b u m p , with Hell A4686 and NIII A4634,4642 transitions. Since the spectrum does not clearly show NV A4602,4620, the W R population could be composed of W N of intermediate type ( W N 6 - 7 ) . T h e broad Hell has an equivalent width of 1 A, and a F W H M of 23 A, coincident with the typical value (about 1500 km s" 1 ) for many W R stars. T h e luminosity of this broad line is 1 0 3 9 1 erg s " 1 , which implies t h a t ~ 1260 W N stars are present in the cluster. T h e C a l l triplet is observed in absorption in N G C 7714, extending 10 arcsec. T h e equivalent width of the lines A8542 and A8662 in the nucleus, measured according to the m e t h o d described in Diaz, Terlevich & Terlevich (1989) is 4.5 A. This value is significantly lower t h a n the 7 A reference given by Terlevich, Diaz & Terlevich (1990), as indicator of the presence of red supergiant stars, and since we detect the presence of W R stars, which indicate an age for the present burst between 3 and 5 Myr, we deduce t h a t the observed C a l l triplet is produced by red giants a n d / o r supergiants associated with a previous episode of star formation.

4. Conclusion NGC 7714, considered as the prototype starburst galaxy, is clearly experimenting a violent burst of star formation, most of which is taking place in the circumnuclear region, making the SFR per unit area higher in the nucleus of the galaxy than in the rest of the galaxy. The detection of WR features indicates the presence of evolved massive stars with M > 30 MQ in the ionizing cluster. The Call triplet absorption lines indicate that the galaxy has experienced a previous burst of star formation.

Gonzalez-Delgado et al.: The starburst galaxy NGC7714

144

5c

50

[Felll]X4658

I 48

Hell X4686

3 42 "4600

8500 FIGURE

4650

4700 wavelength (A)

8600

8700 wavelength (A)

4750

8800

4800

8900

4. (a) WR bump at A4686. (b) Call triplet lines in absorption at the nucleus.

We gratefully acknowledge partial support from DGICYT Grant PB91-0531(GEFE) and the NATO Grant CRG920198 for collaborative research.

REFERENCES BICA, E. 1988 Astr. Astwphys. 195, 76. DIAZ, A. I., TERLEVICH, E.,& TERLEVICH, R. 1989 M. N. R. A. S. 239, 325. KENNICUTT,

R. C ,

EDGAR,

B. K. & HODGE, P. W. 1989 Astrophys. J. 337, 761.

WEEDMAN, D. W., FELDMAN, F. R., BALZANO, V. A., RAMSEY, L. W., SRAMEK, R. A.

Wu, C.-C. 1981 Astrophys. J. 248, 105.

Tracing Violent Star Formation: HI Observations of Nearby Galaxies By ELIAS BRINKS National Radio Astronomy Observatoryt, P.O. Box "O", Socorro, NM 87801, USA HI observations are an excellent tool to probe the conditions of the ISM, giving information on the distribution and velocity structure of the cool and warm atomic gas in galaxies, and how that is affected by violent star formation. Currently, more distant galaxies have come within reach, allowing HI studies of galaxies of different Hubble types. In this talk I will review some of the work on the nearest galaxies and draw comparisons with our own. I will use the example of Holmberg II, a dwarf galaxy, to show how much can be learnt regarding violent star formation and its effects on the ISM by going to low-mass systems. Lastly, I will introduce some of the new and exciting projects which are under way.

1. Introduction Carl Heiles' pioneering work on the structure of the Interstellar Medium (ISM) of our Galaxy confirmed the picture of it being a violent environment. His HI observations showed a wealth of structure in the form of shells and supershells, many of them expanding (Heiles 1979, 1984). His observations agreed, in general, with the prediction by Cox and Smith (1974) who, argued that supernovae create cavities of hot, coronal gas and play a significant role in the shaping of the ISM. Although some work was done on the Magellanic Clouds, it wasn't until the mid-80's that Heiles' results were dramatically confirmed by high-resolution (aperture synthesis) HI maps of the nearest galaxies, M31 and M33, enabling studies of the distribution and characteristics of HI shells across entire galaxies. There have been several review articles in the recent past dealing with the influence of Violent Star Formation on the ISM in galaxies. An excellent general introduction to the ISM is given by Kulkarni k Heiles (1988). Brinks (1990) deals with the cool, atomic phase of the ISM in nearby galaxies, focussing on the nearest neighbours, i.e., the Magellanic Clouds, M31 and M33. Van der Hulst k Kamphuis (1991) and Kamphuis (1993) include some more distant objects, notably M101 and NGC 6946, and include a discussion of HI at high velocities with respect to the "local" velocities in their galactic disks. Of all the papers which deal with theoretical aspects and modelling Ifindthe following especially useful: McCray k Kafatos (1987), Tenorio-Tagle k Bodenheimer (1988), Norman k Ikeuchi (1989) and Heiles (1990). It is well accepted that massive stars have a major impact on the surrounding ISM, first through photo-ionisation and stellar winds. Moreover, as they explode as supernovae (SNe), they deposit some 1051 erg of energy into the ISM, creating a large cavity filled with coronal gas. Cox and Smith (1974) realised that the coronal gas can take a long time to cool down and that therefore the ISM is likely to befilledwith a tunnelling network filled with hot gas. It took several years after their paper had appeared for people to appreciate that violent star formation involves large numbers of massive stars within a relatively small volume. It is beyond the scope of the present contribution to go into further detail and the reader is referred to the theoretical papers listed above. f The National Radio Astronomy Observatory is operated by Associated Universities, Inc., under contract with the National Science Foundation 145

146

Brinks: Tracing Violent Star Formation

When observing in HI, one can map the distribution and radial velocity of the cool and warm neutral gas with temperatures ranging from a few hundred to a few thousand Kelvin. This allows one to map the direct surroundings of a star-forming region, usually a giant HII region fuelled by an OB association.

2. Past—The Galaxy, Magellanic Clouds and Local Group Heiles (1979) lists 63 HI shells and supershells in the Galaxy, located within |6| < 10° and 10° < / < 250°. The shells range up to 1200 pc in size, contain up to 107 MG and many of them show expansion velocities of order 20 kms" 1 . There was substantial uncertainty about the absolute sizes of the HI shells due to the inherent difficulty of determining distances in the Galaxy. Also, because of our vantage point within the Galaxy and the large optical extinction, especially at such low galactic latitudes, it wasn't at all obvious what could have caused these structures. Heiles calculated that if these HI shells were created by a single explosion, energies up to 6 x 1053 erg were required, or the equivalent of several thousands of Type II supernovae. No correlation with Population I objects was found. In retrospect it is not surprising that Heiles couldn't find any direct evidence for a precursor for the HI shells and supershells. Not only does optical extinction make a search for precursors difficult, if not impossible, but by the time an HI shell has expanded to a few hundred pc in diameter, its age is a few xlO7 yr and all stars more massive than 8 M© will have exploded as SNe. As we will see below, even in other, nearby galaxies, it is not at all obvious that OB associations and giant HII regions are the cause of the HI shells (or HI holes) which have been detected. Any evidence has been largely circumstantial. By going to nearby galaxies, one has a much better vantage point. This comes at a price, though, as the linear resolution decreases. To obtain a similar resolution as Carl Heiles in the Galaxy with a single dish telescope one will have to push aperture synthesis telescopes to their limits, even for the nearest of galaxies. Only in the last couple of years has this been achieved with telescopes such as the Westerbork Synthesis Radio Telescope (WSRT), the NRAO Very Large Array (VLA) and most recently the Australia Telescope (AT). Table 1 summarizes some characteristics of each instrument when observing at a wavelength of 21 cm, i.e. corresponding to the hyperfine transition of atomic neutral hydrogen. Because of their proximity, about 50 kpc, the Magellanic Clouds can usefully be observed with a single dish instrument, such as the Parkes 64-m telescope. Brinks (1990) gives a summary of work published before 1990. Since then, Luks & Rohlfs (1992) have resurveyed the LMC resulting in maps at much improved sensitivity and velocity resolution. All eyes are on the AT to produce an HI survey of the Clouds at much higher spatial resolution, possibly achieving a linear resolution of a few pc! At that resolution, a comparison with optical data (e.g. Kennicutt, this volume) will be exciting. Out to about one Mpc, the WSRT is able to provide a linear resolution comparable to that achieved with single dish instruments in the Galaxy. The surveys of M31 (Brinks 1981; Brinks & Bajaja 1986) and M33 (Deul k den Hartog 1990), both at a distance of approximately 0.7 Mpc, have uncovered a wealth of information, confirming the picture Heiles produced for the Galaxy. The ISM in M31 and M33 is dominated by HI holes (in nearby galaxies the term HI holes was used rather than HI shells, as the linear resolution didn't allow the detection of the thin HI shells surrounding the cavities). Much for the same reason as quoted above, no clear correlation with Population I objects was found,

Brinks: Tracing Violent Star Formation

Instrument Configuration Number of elements Antenna diameter (m) Primary beam (arcmin) System temperature (K) Sensitivity (mJy beam" 1 )

WSRT A 14 25

37.6 55 2.0 7.2 2.8

B

VLA C

27 25 30 30 0.5 1.4 15 3.4 11.4

147

D

AT Compact Array 6 22 35 18 1.1 16 6.0

140 0.16 (K) 1.03 Longest Baseline (km) 36.4 0.036 0.030 Shortest Baseline (km) 0.68 0.21 0.073 0.033 3.9 44 12.5 6.5 x 6.5/ sin 8 Synthesized beam (arcsec) 13 x 13/ sin 8 1.4 2 15 14 7 12 0.6 Largest structure (arcmin) TABLE 1. Characteristics of the WSRT, VLA and AT-Compact Array at 1420 MHz. Note that the sensitivity in mJy beam"1 corresponds to the one-sigma rms noise reached after a 12h period. The frequency resolution of respectively 19.5, 12.2, and 15.6 kHz converts to a velocity resolution of 2.5-4.1 kms" 1 . The sensitivity in terms of surface brightness, expressed in Kelvin, assumes that the emission is distributed uniformly and fills the beam.

although there were suggestions for a correlation with OB associations and HII regions (see Brinks 1990 for a review). 3. Present—Nearby groups out to 10 Mpc 3.1. Holmberg II

From one Mpc out to perhaps ten Mpc, the power of the VLA (mainly using the Barray) allows one to achieve a linear resolution similar to that achieved in M31 and M33 in galaxies beyond the Local Group. A stunning example of the effect of violent star formation on the ISM is the HI image of Holmberg II. Ho II is a gas-rich dwarf galaxy in the M81 group. It was imaged by Puche et al. (1992) with the VLA at a spatial resolution of 4-5 and at a velocity resolution of 2.5 kms" 1 . Figure 1 shows the HI column density image at full resolution. At an assumed distance of 3.2 Mpc the HI extent is about 14 kpc. The entire galaxy is shot with holes, some 51 having been definitely identified. At about the same scale, Figure 2 shows an outline of the location, size and shape of the HI holes which were identified. It was assumed that their shape can be approximated by an ellipse. Using the same definitions, and assumptions, as Brinks & Bajaja (1986), such properties as the size, location within the galaxy, and ellipticity of the HI hole and expansion velocity of the surrounding HI shell were determined. From these observables, the age, amount of displaced HI and an estimate for the required energy were derived. In Ho II we were able, for the first time, to check that the HI holes are indeed expanding HI bubbles. This was happily assumed by all workers in the field, but hadn't been convincingly demonstrated. Figure 3 shows a cross-cut through hole number 35. We plot here the observed HI expansion velocity across this hole as derived from the observed line splitting of the HI profiles. The projected (radial) component follows beautifully the expected cosine dependency confirming that, to first order, this hole is spherically expanding (although we cannot tell an expanding bubble from a contracting one, the former option is far more likely than the latter). The properties of the HI holes in Ho II are similar to those found in other well-studied galaxies such as in M31 and M33, and they are in line with what has been found for

148

Brinks: Tracing Violent Star Formation

50

100

71 00 • <

70 58

& w

-

X v-

56

CD

z

54

Z

52

O UJ Q

50

48 I

08 15 00

14 30 00 13 30 00 RIGHT ASCENSION (B1950) Grey scale flux range= 0.0 100.0 JY/B*M/S

FIGURE

1. HI column density map of Holmberg II by Puche et al. (1992).

the Galaxy. Linear sizes range from about 100 pc (the resolution limit) up to 1500 pc and expansion velocities reach 15 k m s " 1 ; characteristic ages (based on size divided by expansion velocity) are 107 to 108 yr; kinetic energies are 1050 to 10 53 erg and the HI mass which has been displaced ranges from 104 to 107 M©. A more careful comparison between the properties of HI holes in large (M31 and M33) galaxies and a dwarf system (Ho II) shows that the HI holes in Ho II are larger, that they have on average lower expansion velocities, and that the HI holes are rounder. This can be understood as follows. Ho II is a low-mass system. Based on the HI rotation curve we find Mdyn = 2 x 109 M Q - The maximum rotational velocity is a low 40 k m s " 1 which means that differential rotation or shear has a much smaller influence on the shape and longevity of an HI hole than in more massive systems (Palous et al. 1990).

Brinks: Tracing Violent Star Formation 71°00'

—

T

149 I-

58' —

56'

q d

54<

~ O ill

52'

ID O)

Q

50'

48'

46'

70°44'

-

I h

m

j_

8 15 30

s

14 m 30 s

13 m 30 s

12 m 30 s

RA (1950.0) FIGURE 2. Location of the HI holes in Holmberg II.

The smaller gravitational potential has some other interesting effects. We measure in Ho II a velocity dispersion which is about equal (6-7 kms" 1 ) to that found in large spirals. This then implies that the HI layer will have a much larger scale height, 625 pc rather than 120-130 pc. This has two consequences. Firstly, the gas volume density in Ho II is lower than in more massive spiral disks, which makes for larger HI holes for the same amount of SN energy input. Because the HI scale height is larger, the HI holes can grow larger before experiencing "break-out" and subsequent loss of spherical symmetry and internal pressure. 3.2. What fills the holes?

The fact that an HI hole is seen does not necessarily imply that there is a true absence of material, although this is the most likely explanation. In fact, the material inside, or filling, an HI hole could be molecular. Searches for CO as a tracer for molecular gas in some of the most well-defined HI holes in M33 have resulted in upper limits (Deul &

150

Brinks: Tracing Violent Star

Formation

15

•

/

en

E

IK

I \

I/I

Q

5 -

i

•

i; /

\

\

T

-

\ .

.

/ 1

: / " -40

-20

0

20

40

arcsec FIGURE 3. Cross-cut through hole number 35 in Holmberg II. Plotted is the expansion velocity AV along a cut through the HI shell. The solid curve corresponds to a cosine function which is expected in the case of spherical expansion.

den Hartog 1990). Assuming that the CO-to-H 2 conversion factor is similar in M33 and the Galaxy, the most likely explanation is that the HI holes trace a genuine absence of material. This is corroborated by another important result on Ho II which supports the growing evidence that violent star formation events, or rather the after effects, are indeed responsible for creating the HI holes. A comparison with CCD direct images through a narrow-band Ea filter shows that many of the smaller holes are filled with RQ emission. The largest holes show H a emission along their rims. This suggests that, by the time the first supernovae go off to create the beginnings of an HI shell, there are still plenty of ionising stars around. Once all the ionising stars (with masses typically larger than 8 M Q ) have exploded as supernovae, the HI shell has grown much larger and no obvious stellar remnant, other than perhaps a cluster of blue stars, will remain. Gas piling up along the expanding rim experiences shocks and is likely to exceed some threshold for star formation, and sites of secondary star formation are formed. The HI holes are genuinely empty with volume densities of the remaining gas of order 0.001 atom c m " 3 . Moreover, this gas must be hot with temperatures typically 5 x 105 to 10 7 Kelvin. The best evidence to date for this coronal gas is formed by Einstein observations of LMC 2 in the Large Magellanic Cloud presented by Wang k Helfand

Brinks: Tracing Violent Star Formation

151

59 08

58 5 8 00 18 15

00

17 45 30 RIGHT ASCENSION

15

FIGURE 4. The non-thermal superbubble in IC 10 as observed by Yang & Skillman (1993) in grey scale superimposed on the HI column density map of Shostak & Skillman (1989). Contour levels are: 0.5, 1.0, 1.5, 2.0, and 2.5 x 1021 atom cm . The positions of Wolf-Rayet stars are indicated with crosses.

(1991) and through ROSAT observations of LMC 4, a supergiant shell coinciding with an HI minimum centred on Shapley Constellation III. The ROSAT data are presented by Bomans et al. (1994), who derive a temperature of 2.4 x 106 K, an electron density of 0.0056 atom cm" 3 and an X-ray luminosity of 1.4 x 1037 erg s" 1 . They also find that the superbubble still has considerable overpressure as compared to the pressure of the warm ionised medium, which implies that the largest superbubble in the LMC hasn't experienced "break-out" yet and thus confirms that the scale height of the ISM in the LMC is much larger than in the Galaxy, and probably comparable to that found in Ho II. The scenario presented above suggests that there must be precursor objects to the HI holes. Obviously, these objects are rare as the phase during which one has massive O stars embedded in their forming molecular environment before they go off as supernovae is extremely short, a few x 106 yr at most. It seems, though, as if one such object has been found. Yang & Skillman (1993) report on a non-thermal superbubble in the Local Group irregular galaxy IC 10. Figure 4 shows an overlay of their 20-cm VLA radio continuum map (grey-scale) with a contour map of the HI column density from Shostak & Skillman (1989). The radio continuum emission coincides with a giant HII region, about the size of 30 Dor in the LMC. The non-thermal radio emission indicates that at least a few SNe have occurred. However, there is no sign, yet, for an HI shell or hole (at least at the

152

Brinks: Tracing Violent Star Formation

rather modest linear resolution of 200 pc). Rather, the radio emission coincides with a peak in the HI emission. This all argues for this region to be older than 5 x 106 yr, i.e. old enough for some 0 stars to have exploded as SNe, but younger than 107 yr, or too young for an HI supershell to have evolved. Follow-up VLA HI observations (Wilcots, Miller &: Hodge, private communication) may allow detection of a supershell in statu nascendi, possibly showing large expansion velocities. 3.3. Holes, mfall and HVCs Kamphuis (1993) presents extensive new data on three more galaxies, M101 (at 7.2 Mpc), NGC 628 (at 10.4 Mpc) and NGC 6946 (at 10.1 Mpc). His WSRT observations have a linear resolution of about 400 pc. He therefore misses the smaller holes, and some of his HI holes might be blends. Having said that, he finds ample evidence for fine-scale structure in the ISM of these galaxies, similar to that found in the Galaxy, M31 and M33. Interestingly, Kamphuis finds evidence for high velocity gas, i.e. gas with velocities in excess of 30-100 kms" 1 with respect to the "local" disk velocities in some of his objects. A good example is NGC 6946. Many of these high velocity features seem to correlate with holes which have suffered "break-out" which suggests that gas is blown into the halo to a height of up to a few kpc above the disk. Although the vast majority of HI holes can be explained by violent star formation, there are cases where the energy requirement to produce an expanding bubble seem far too large, i.e. mass involved in the expansion of a few xlO8 MQ and a kinetic energy of 1055 erg. A good example is M101, where van der Hulst & Sancisi (1988) find a highvelocity complex (HVC) which they explain as due to the collision of a dwarf galaxy or tidal debris with the disk of the galaxy (see also Kamphuis 1993). Another example is NGC 628, where an Hi-rich companion may have produced two HVCs (Kamphuis & Briggs 1992). Two extreme HI supershells were recently reported by Rand & van der Hulst (1993) in NGC 4631. The authors again postulate that a collision of small companions or massive gas clouds are the origin for these features. Given the highly disturbed and filamentary distribution of HI in and around NGC 4631, this seems quite a plausible explanation.

4. Future—Recent results and ongoing projects This presentation can be summarised as follows: (a) Violent Star Formation has a tremendous influence on the ISM in galaxies. (6) HI observations are eminently suited to map the effects of Violent Star Formation on the general ISM as it traces the morphology as well as the velocity of the cool and warm neutral medium. (c) Typically, 50-150 HI shells larger than 100 pc are found. (d) Energy requirements to create the shells range from several hundred Type II supernovae to several thousand within the volume of the violent star formation region. A typical example of such a region is the nearby giant HII region 30 Doradus in the LMC. (e) Violent Star Formation has a more serious (deleterious) effect on lower-mass, i.e. dwarf galaxies. (/) There is little doubt that the majority of the HI shells is indeed caused by the combined effects of stellar winds and supernova explosions of the most massive stars, creating a cavity filled with coronal gas. (g) Up to one or two of the largest HI shells in some galaxies could be due to infall of tidal debris or a dwarf companion.

Brinks: Tracing Violent Star Formation 200

153 400

68° 481

44' o m

|

o Q

361

68° 32* 10 h 26 m

10 h 25 m

10 h 24 m

RIGHT ASCENSION (B1950) FIGURE

5. HI column density map of IC 2574 by Martimbeau et al. (1994). Grey levels range from 0 to 540 mJy beam"1 kms" 1 .

With current instruments, such as the WSRT, the VLA, and the AT, a whole range of galaxies has come within reach. A linear resolution of 50 to 100 pc, i.e. similar to that of the M31 and M33 studies, can now be achieved out to about 5 Mpc. This has prompted a flurry of activity. Without claiming completeness, I list here some of the projects which are currently being undertaken. As already mentioned, IC 10 is being studied by Wilcots et al. (private communication). VLA C-array (i.e. 15 arcsec spatial resolution) maps show an ISM dominated by HI shells and holes. Focussing on the Sculptor group, Cote et al. (1994) are making an inventory of all its dwarf galaxies. It will only be a matter of time before the larger Sculptor galaxies will be reobserved. Martimbeau et al. (1994) have made a detailed study with the WSRT of IC 2574, a member of the M81 group of galaxies. Figure 5 shows the HI column density map for this object. They note that, as in Ho II, HII regions are predominantly found along the rim of the larger HI shells and

154

Brinks: Tracing Violent Star Formation

that HII regions in general tend to avoid regions of high HI column density. Regarding the M81 group of galaxies, Puche & Westpfahl (1994) continue mapping all gas-rich, dwarf companions at resolutions similar to Holmberg II. One of the fascinating results is that the smaller the dwarf galaxy, the more dramatic the effect of violent star formation becomes. In the smallest objects such as in Holmberg I only one or two HI shells are found which occupy the entire (tiny) HI disk. They speculate that the HI is expelled from the dwarf galaxy by the after effects of the star formation event and that new star formation is put on hold until the gas cools and falls back onto the dwarf galaxy. A long overdue effort to map M81 itself is under way at the VLA (Adler & Westpfahl 1993), and initial maps already show a wealth of detail. A unique opportunity to study violent star formation in ultimate detail is offered by the Magellanic Clouds. Eventually, HI maps made with the AT will become available and will no doubt shed new light on the interaction of supernovae with the surrounding ISM. HI studies of nearby galaxies have now reached a point where they can be usefully compared with optical data. A first attempt to compare H a data, obtained with the TAURUS Fabry-Perot interferometer with HI data is presented by Brinks et al. (1990), who compare several 1 kpc diameter fields in M31. No doubt this kind of work will become more and more important in the future. I am most grateful to Guillermo Tenorio-Tagle for inviting me to "la isla bonita" and to the organisers for their efforts to make this workshop such a resounding success.

REFERENCES D. S. & WESTPFAHL, D. J. 1993 Bull. Am. Astton. Soc. 25, 1393. D.J., DENNERL K. & KURSTER, M. 1994 Astion. Asttophys. In press. BRINKS, E. 1981 Astron. Astiophys. 95, LI. BRINKS, E. 1990. In The Interstellar Medium in Galaxies (ed. H. A. Thronson, Jr. &; J. M. Shull), p. 39. Kluwer. BRINKS, E. &; BAJAJA, E. 1986 Astion. Astiophys. 169, 14. BRINKS, E., BRAUN, R. & UNGER, S. W. 1990 In IAU Coll. No. 120 on Structure and Dynamics of the Interstellar Medium (ed. G. Tenorio-Tagle, M. Moles & J. Melnick). Lecture Notes in Physics, Vol 350, p. 524. Springer. COTE, S., FREEMAN, K. & CARIGNAN, C. 1994. In ESO/OHP Workshop on Dwarf Galaxies (ed. G. Meylan & B. Binggeli). In press. ESO. Cox, D. P. & SMITH, B. W. 1974 Astiophys. J. (Letteis) 189, L105. DEUL, E. R. & DEN HARTOG, R. H. 1990 Astion. Astiophys. 229, 362. HEILES, C. 1979 Astiophys. J. 229, 533. HEILES, C. 1984 Astiophys. J. Suppl. 55, 585. HEILES, C. 1990 Astiophys. J. 354, 483. KAMPHUIS, J. J. 1993 PhD thesis, University of Groningen. KAMPHUIS, J. & BRIGGS, F. 1992 Astion. Astiophys. 253, 335. KULKARNI, S. R. & HEILES, C. 1988 In Galactic and Extragalactic Radio Astronomy, 2nd edn. (ed. G. L. Verschuur & K. I. Kellermann), p. 95. Springer. LUKS, T H . & ROHLFS, K. 1992 Astion. Astiophys. 263, 41. MARTIMBEAU, N., CARIGNAN, C. & ROY J.-R. 1994 Astron. J. In press. MCCRAY, R. & KAFATOS, M. 1987 Astiophys. J. 317, 190. NORMAN, C. A. & IKEUCHI, S. 1989 Astiophys. J. 345, 372. PALOUS, J., FRANCO, J. & TENORIO-TAGLE, G. 1990 Astion. Astiophys. 227, 175.

ADLER,

BOMANS,

Brinks: Tracing Violent Star Formation PUCHE, D. & WESTPFAHL, D. 1994 In ESO/OHP & B. Binggeli). In press. ESO.

155

Workshop on Dwarf Galaxies (ed. G. Meylan

P U C H E , D., WESTPFAHL, D., BRINKS, E. & R O Y , J.-R. 1992 Astron.

J. 103, 1841.

RAND, R. J. & VAN DER HULST, J. M. 1993 Astron. J. 105, 2098. SHOSTAK, G. S. & SKILLMAN, E. D. 1989 Astron. Astrophys.

214, 33.

TENORIO-TAGLE, G. &; BODENHEIMER, P . 1988 Ann. Rev. Astion. Astrophys. 26, 145. VAN DER HULST, T. &C KAMPHUIS, J. 1991 In IAU Symp. No. 144 on The Interstellar Disk-Halo Connection in Galaxies (ed. H. Bloemen), p. 201. Kluwer. VAN DER HULST, T. & SANCISI, R. 1988 Astron.

WANG, Q. & HELFAND, D. J. 1991 Astrophys. YANG, H. & SKILLMAN, E. D. 1993 Astron.

J. 95, 1354.

J. 379, 327.

J. 106, 1448.

Massive Star Formation and Supergiant Shells in the Irregular Galaxy NGC 55 ByDOMINIK J. BOMANS1 AND EVA K. GREBEL 21 'Sternwarte, Univ. Bonn, Auf dem Hiigel 71, D-53121 Bonn, Germany 2

European Southern Observatory, Casilla 19001, Santiago 19, Chile

The highly inclined very late type galaxy NGC 55 is a perfect target to study the effects of massive star formation on the interstellar medium, especially in the disk-halo interface. We present deep broad- and narrow-band images of a field centered on the two largest HII regions of NGC 55. The color-magnitude diagrams produced from broad band V, R, i, and z images show the effects of a strong burst of recent star formation and imply the presence of a population of blue clusters which show up as clearly exdended sources on our best seeing images. In the Ha image a spectacular arrangement of filaments, shells and supergiant shells is visible. This ionized gas at distances of up to 1 kpc from the nearest star-forming region makes NGC 55 a new example of a galaxy in 'chimney mode'.

1. Supergiant shells in irregular galaxies HII supergiant shells (SGSs) in irregular galaxies were first detected in the Magellanic Clouds on deep Ha plates taken with the UK Schmidt telescope (Goudis & Meaburn 1978). Thereafter, more and more such structures were detected in nearby irregular galaxies using modern detectors (e.g. Hunter k Gallagher 1990; Bomans & Hopp 1992; Hunter, Hawley & Gallagher 1993). The objects seem to be related to active star formation and HI holes, but the pattern is far from being understood. According to a widely accepted theory massive stellar associations create supergiant shells by excavating a large hole in the ambient HI by the wind of their massive stars and later supernova explosions (Leitherer, Robert & Drissen 1992). This cavity is filled with hot gas (Bomans et al. 1994). The HII filaments defining supergiant shells are the shell walls ionised by the Lyman continuum photons from the hot stars inside the cavity and maybe by thermal conduction with the hot interior via turbulent mixing layers (Slavin, Shull & Begelmen 1993). The outward-propagating shock wave may induce secondary star formation (e.g. Elmegreen 1993), while star formation is effectively suppressed in the interior. This way supergiant shells can control the large-scale star formation activity of irregular galaxies. During the expansion phase of a supergiant shell, this also grows perpendicular to the galactic plane (the HI layer appears to be flattened also in irregular galaxies (Skillman 1994). The lowering of the density with z speeds up the shock and leads finally to a breakthrough or even breakout of the supergiant shell though the galactic gas layer (see Heiles 1990; Tenorio-Tagle & Bodenheimer 1988). As a consequence hot, metal-enriched gas escapes into to the halo, as observed, for example, in NGC 253 (e.g. Pietsch k. Triimper 1993) where not only the central starburst seems to contribute to the X-ray halo. Venting up enriched gas into the halo of a galaxy can act as an effective mechanism for mixing the ISM and may control the temporal evolution of global bursts via star formation caused by cooled clouds falling back to the galaxy. This would lead to a self-regulating circle for the whole galaxy (Tenorio-Tagle & Bodenheimer 1988). 156

Bomans & Grebel: Massive Star Formation in NGC 55

-

*

• ' .

FIGURE

•

.

.

•

.

.

.

,

.

.

;

,

•

157

••

•

•

1. Scan of a B image of NGC 55 taken from the ESO LV database. The size and location of the analysed field is marked.

The superwinds observed in some starburst galaxies (e.g. Heckman, Armus & Miley 1990) may be similar phenomena at much higher energies. The conditions for the venting of large amounts of gas are not so clear in less extreme galaxies. Low-density HI layers with large z extent may effectively prevent supergiant shells from breakout. Also, the large scale height of diffuse ionized gas observed in NGC 891 by Dettmar (1990) and Rand, Kulkarni & Hester (1990) is a linked phenomenon. Lyman continuum photons from large HII regions in the galactic plane seem to be able to produce these halos (Domgorgen & Mathis 1994) if there are channels of low HI density, which could be produced by supergiant shells. NGC 55 with its proximity, large inclination angle, and active star formation is a nice target to study the effects of massive star formation on the ISM and especially the interactions with the lower halo. 2. The data We observed NGC 55 with the ESO/MPIA 2.2-m telescope at La Silla equipped with the EFOSC2 focal reducer/imager. The field of view is 5-8 x 5-8 on a low-noise Thompson 10242 chip. The resulting pixel size is 0"338. We took frames in the Bessell V and R and Thuan-Gunn i and z filters, and through an Ha filter. The Ha filter has a central wavelength of 655 nm and a width of 6 nm, well fitted to the radial velocity of NGC 55. The seeing was between 0-9 and l'-'3. Figure 1 shows a B plate of NGC 55 with the location and size of our CCD field indicated. All data were reduced using MIDAS on the workstation cluster of the Astronomical Institutes of Bonn University. The point-source photometry was performed with DAOPHOT II. 3. Resolved stellar photometry NGC 55 is nicely resolved into stars on our CCD frames. Also, the background sheet of red stars as decribed by Graham (1982) is clearly visible (Figure 2), implying resolution down to the tip of the asymptotic giant branch or even red giant branch. The obvious major problem in performing photometry of the stellar images is the

Bomans & Grebel: Massive Star Formation in NGC 55

158

*•*

FIGURE 2. V frame of the NGC 55 field.

highly variable background due to the edge-on orientation of NGC 55 and the bright HII regions. We tried to compensate for this by subtracting a low-pass filtered image from the original frame. The low-pass was effected by first subtracting all bright stars with DAOPHOT II and then running a 15 by 15 pixel median filter on the subtracted image, similar to the method of Pierce, McClure & Racine (1992). The resulting image contains only the high spatial frequencies, i.e. the stars. Some dust lanes are clearly visible, which implies a quite large amount of dust especially in the large eastern star-forming region. We should expect differential reddening to affect the stars in this region. Our final photometry contained 5000 well-fitted stars with both V and R magnitudes and still 2700 stars measured in all four broad-band colors. Calibration of the V, R, and i images was done using standard fields of Landolt (1992). The resulting color-magnitude diagrams (Figure 3) are quite typical of those of irregular galaxies, see, for example, NGC 3109 (Greggio et al. 1993). The typical blue plume is present as well as the red supergiant branch. Comparing these features with those of NGC 3109 clear differences in the structure of the CMDs are visible. They imply a strong recent burst of star formation in the NGC 55 field with

Bomans & Grebel: Massive Star Formation in NGC 55

159

i

16M,,

16 '-

18

' .

18

_

> 20 20 22 22

• • • •

• - .

- |

, . . i , . . . i , . , ,:

24

-1

[' "*$$$ ': |,

- 0.5

0

0.5

1

1.5

2

-1

-0.5

0

0.5 R-l

1

1.5

2

FIGURE 3. V - R, V and R-i, R color-magnitude diagrams of our NGC 55 field.

an age less than 4 x 106 yr, as one might expect from the presence of the giant HII regions. The large population of stars of intermediate color around V = 20.5 mag may be a combined effect of reddening and unresolved blends. The group around 19.5 mag may be real and would then be an important feature to account for in an analysis using synthetic CMDs. The red supergiant branch is consistent with stars as old as 1 x 107 yr, which implies a high star formation rate over at least this time interval for this region in NGC 55. A more detailed analysis of the resolved stellar content of our field using synthetic color-magnitude diagrams is in preparation. The large group of very luminous stars somewhat redward of the blue supergiant location may in part be an effect of the differential reddening in the observed field, but in this case the stars would be extremely luminous. This finding and the clearly extended shape of some of these objects on our best seeing images implies that we see here a population of blue clusters similar to the ones found in the Magellanic Clouds. The objects are not compact HII regions because they do not appear in the continuumsubtracted Ha image. We derived a distance for NGC 55 using the i-band brightness of the tip of the red giant branch (Lee, Freedman & 1993). We arrived at a distance modulus of 25.9±0.2, in good agreement with earlier estimates. Thus NGC 55 is at the front side of the Sculptor galaxy group. 4. HII shells and filaments in NGC 55 The continuum-subtracted Ha image of our field in NGC 55 reveals a spectacular collection of shells, filaments and supergiant shells (Figure 4). Two regions are of special interest: the large, coherent loop with a diameter of about 800 pc north of the main body of NGC 55 (SGS 1) and the large web of loops and filaments of comparable size in the south-east (SGS 2). South of the western giant HII region a large shell (SGS 3) is located, and east of this giant HII region a bright arc is present as well as numerous smaller shells. Diffuse emission with filaments up to 150 pc long is present around the two HII region complexes.

160

Bomans & Grebel: Massive Star Formation in NGC 55

FIGURE

4. Continuum-subtracted Ha image of the centralfieldof NGC 55.

The SGS 1 was already seen in the OIII image of Graham and Lawrie (1982). The loop is relatively bright in OIII compared to Ha, which is somewhat unusual compared to the (small number of) supergiant shells with published line ratios. Comparing the location of the bluest and brightest stars with the Ha emission, we see no such stars inside the supergiant shell. The only locations in which we can detect these stars are the giant HII regions south of this supergiant shell. The same holds for the large web of Ha filaments and SGS 2. Here a 400 pc long horizontal filament is particularly noteworthy. If it is part of a loop (which is indicated judging from some other parts of the web), it rivals SGS 1 in size. Some very faint surface brightness filaments are visible south of the web. Their distance from sites of recent star formation is more than a kpc, which makes these filaments similar to those detected in NGC 4449 by Hunter & Gallagher (1992) and Bomans k Hopp (1993). Interestingly, there are some bright blue stellar objects north of the eastern giant HII region, where no, or only very weak, diffuse emission is present. 5. Discussion The supergiant shells 1 and 2 in NGC 55 are comparable in size and appearance of the HII emission to the supergiant shells in the Magellanic Clouds. The apparent lack

Bomans & Grebel: Massive Star Formation in NGC 55

161

of a large association inside is a strong distinction though. The supergiant shells LMC 4 and 3 both contain very large associations. LMC 2 is somewhat less clear, and only LMC 1 contains a small number of hot stars, which may be an age effect. When comparing NGC 55 and the LMC one has to keep in mind that we see the two galaxies approximately edge on and face on, which means that the LMC supergiant shells look like holes in the gas structure, while the loops in NGC 55 can be interpreted as the expanding caps of supergiant shells. If this interpretation is right, we see a breakthrough into the lower halo region taking place in NGC 55. The size of the shells pose an interesting question: if they expand with similar speed as LMC 4 or LMC 2 (less than 50 km s" 1 ), the supergiant shells in NGC 55 are older than the giant HII regions in the main body of NGC 55. Two solutions are possible: the supergiant shells (or more precise the caps) expand much more faster into the lower halo, or these filaments are remnants of an older burst of star formation in the same region which triggered the recent burst. In the first case it is difficult to understand how the filaments are so well ordered and not affected by plasma instabilities. The very faint filaments even farther away from the HII regions are difficult to explain by either possibility.

REFERENCES BOMANS, D. J. & HOPP, U. 1992, In Evolution of interstellar matter and dynamics of galaxies (ed. J. Palous, W. B. Burton & P. O. Lindblad), p. 63. Cambridge Univ. Press, Cambridge. BoMANS, D. J. fe HOPP, U. 1993 In Star forming galaxies and their interstellar medium (ed. J. Franco, F. Ferrini & G. Tenorio-Tagle), p. 159. Cambridge Univ. Press, Cambridge. BOMANS, D. J., DENNERL, K. & KURSTER, M. 1993 Astron. Astrophys. Accepted. Dettmar, R.-J. 1990 Astton. Astrophys. 232, L15. DOMGORGEN, H. & MATHIS, J. S. 1994 Astrophys. J. Accepted. ELMEGREEN, B. G. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto & F. Sanchez), p. 381. Cambirdge Univ. Press, Cambridge. Goudis, G. & Meaburn. J. 1978 Astron. Astrophys. 68, 189. GRAHAM, J. A. 1982 Astrophys. J. 252, 474. GRAHAM, J. A. &; LAWRIE, D. G. 1982 Astrophys. J. Lett. 253, L73. GREGGIO, L., MARCONI, G., TOSI, M. & FOCARDI, P. 1993 Astron. J. 105, 894. HECKMAN, T. M., ARMUS, L. & MILEY, G. K. 1990 Astrophys. J. Suppl. 74, 833. HEILES, C 1990 Astrophys. J. 354, 483. HUNTER, D. A. & GALLAGHER, J. S. 1990 Astrophys. J. 362, 480. HUNTER, D. A., HAWLEY, W. N. & GALLAGHER, J. S. 1993, Astron. J. 106,

1797.

A. U. 1992 Astron. J. 104, 340. M. G., FREEDMAN, W. L. & MADORE, B. F. 1993 Astrophys. J. 417, 553. LEITHERER, C , ROBERT, C. & DRISSEN, L. 1992 Astrophys. J. 401, 596. PIERCE, M. J., MCCLURE, R. D. & RACINE, R. 1992, Astrophys. J. 393, 523. PlETSCH, W. & TRUMPER, J. 1993, In Space Astronomy (ed. J. Triimper, C. Cesarsky, G. G. C. Palumbo & G. F. Bignami), Advances in Space Research 13(12), 171. RAND, R. J., KULKARNI, S. R. & HESTER, J. J. 1990 Astrophys. J. Lett, bf 352, Ll (Erratum Astrophys. J. Lett. 362, L35). SLAVIN, J. D., SHULL, J. M. & BEGELMAN, M. C. 1993 Astrophys. J. 407, 83. SKILLMAN, E. 1994 This conference. TENORIO-TAGLE, G. & BoDENHEIMER, P. 1988, Ann. Rev. Astron. Astrophys. 26, 145. LANDOLT, LEE,

Galactic Supershells By S. A. SILICH 1 ,! J. FRANCO 2 , J. PALOUS 3 AND G. TENORIO-TAGLE 4 1

Main Astronomical Observatory of the Ukrainian Academy of Sciences, 252127, Kiev-127, Goloseevo, Ukraina 2

3

Instituto de Astronomia UNAM, Apartado Postal 70-264, 04510 Mexico D.F., Mexico

Astronomical Institute, Academy of Sciences, Bocni II 1401, 141 31 Prague 4, Czech Republic 4

Instituto de Astrofi'sica de Canarias, 38200-La Laguna, Tenerife, Spain

The results of 3-D numerical simulations for superbubbles expanding in a sheared and cloudy gaseous disk are presented. Assuming a disk in rotational equilibrium, the effects due to the gravitational force perpedicular to the Galactic plane, differential rotation, cloud evaporation, and radiative cooling of the gas inside the hot cavity have been included. We consider a set of different stratifications for the cloud component and discuss their influence on the final bubble properties.

1. Introduction Massive stars are born in groups and their collective energy input (via UV radiation, stellar winds, and supernova explosions) causes the agglomeration of large masses of gas in cold expanding supershells (see Tenorio-Tagle fc Bodenheimer 1988 and references therein). These expanding supershells are common interstellar features in our Galaxy (Heiles 1979; Lozinskaya & Sitnik 1988) and in other nearby galaxies (see Brinks 1994 and references therein). In this contribution we continue the study of expanding multisupernova shells with 3-D numerical simulations (Palous 1990; Bisnovatyi-Kogan & Silich 1991; Palous 1992; Silich 1992; Silich et al. 1994a) as an attempt to understand their link with the general structure of the ISM and with the star formation process. 2. Equations and model parameters The expansion of large multi-supernova shells is described using the thin layer approximation. This approximation corresponds to the late stages of supernova remnant evolution and has been applied to a wide variety of astrophysical problems with both analytical and numerical methods (Bisnovatyi-Kogan & Blinnikov 1982; Tenorio-Tagle & Palous 1987; Mac Low k, McCray 1988; see recent reviews by Ostriker & McKee 1988; Franco et al. 1992; Tomisaka 1993). After the mass ejected in a supernova explosion is thermalized, the remnant enters the quasi-adiabatic stage,when, due to the low densities and high temperatures within the structure, radiative cooling is unimportant. Later on, when the remnant evolutionary time becomes larger than the cooling time in the shocked gas, radiative losses reach a maximum value, and a thin and cold shell structure forms (see Franco et al. 1994). The expansion velocity remains supersonic after thin shell formation and the shell continues to collect the ambient gas. The thin layer approximation method assumes that the swept-up mass is accumulated in the thin shell, and that the hot medium inside the cavity can be described with the average values of the hydrodynamical variables. The shell is divided into N Lagrangian elements and the motion of each of these t Present address: Astronomical Institute, Bocni II 1401, 141 31 Prague 4, Czech Republic 162

Silich et al: Galactic Supershells

163

is followed using t h e equations of mass and m o m e n t u m conservation. T o include t h e influence of small and dense interstellar clouds, this set of equations is coupled with three additional equations for the cloud component: t h e total number of clouds inside the r e m n a n t , the total evaporated mass Mev, and the equation of energy conservation. We define nci(z) as the mean space density of clouds with radius Rci and mass m c /, a n d assume t h e classical or saturated thermal evaporative mass loss rate mev for t h e clouds engulfed by t h e shock wave (Cowie et al. 1981). We do not explicitly take into account the distortion generated during the interaction between t h e clouds and t h e expanding shell (Klein et al. 1994; Mac Low et al. 1994), nor the reduction of cloud radius due t o cloud ablation (Hartquist et al 1986). We only assume t h a t t h e number of clouds inside the remnant, JVci, is reduced due to the evaporated mass rate dMev/dt :

^

= mm N. cl. evN evcl at The total number of clouds inside the remnant is described with the equation: dNc,

A

,

,

,„

(2.1)

1 dMev

— = g M u - vel)n,dE, -—-#-,

(2-2)

where u and vcj are the shell and cloud velocities, n* is the unit vector perpendicular to the surface of the Lagrangian element, and d£,- is the surface area of the Lagrangian element. Three effects leading to a continuous change in the total thermal energy of the remnant are included: • The mechanical energy rate L(t) from supernova explosions, • the P — V work done by the shell expansion, and • ionization of the evaporating cloud material in the hot bubble interior:

Here q is the energy required to ionize a unit mass of cloud material, Nj is the number of particles, \j ls * n e ionization potential, and rrij the mass per particle. The sub-index j denotes each of the different components in the gas mixture. We assume a "normal" chemical composition by number, with 90% hydrogen and 10% helium. The equation of energy conservation is written as:

^

= L(t) - Pindtt - nlA(T)Q' - Mevq

(2.4)

where P, n and n, n are the mean pressure and number density of ions inside the remnant, A is the gas cooling function, $1' = (1 — /„)V is the volume of the remnant which is filled by the intercloud material, and fv is the cloud volume rilling factor (Cowie et al. 1981). The average structure of the gaseous disk in the ^-direction is approximated with the density distribution discussed by Dickey & Lockman (1990): H

/ \ *

"

l

"

2

+

"

3

+

H2y

+

exp(| z \ /H3)'

/

„

,., ^

^

where nx are the midplane densities of the 3 Gaussian components, and Hx are the corresponding scale-heights. This average density distribution is in turn redistributed into two components: a warm intercloud medium, and a cloudy component with small spherical clouds with a radius of 5 pc and a density of 10 cm" 3 The fraction of mass in clouds is: fm{z) = (P:f] pg[z)+Pcl{Z)

,y

(2-6)

164

Silich et al.: Galactic Supershells

f

*max

Zmax

Zmin

[Myr]

[PC]

[PC]

0.1 0.3

30.3 30.1

1090 1162

-1090 -1162

0.1 0.3

30.0 30.0

1390 1488

-567 -600

Nd [103 MQ] R = 8.5 kpc 1730 5220 R = 8.5 kpc 1640 4860

z 0 = 0 pc 36.5 55.6 z0 = 50 pc 35.1 53.0

M!h

Ltot

[106 M 0 ]

[10 36 erg s" 1 ]

2.94 2.61

1.58 4.23

2.60 2.30

1.38 3.57

T A B L E 1. Superbubble parameters at the final state

where pg is the intercloud gas mass density, pc\ is the contribution of the clouds to the mean ISM density, and pg + pc\ is given in equation (2.5). We assume that the number of clouds per unit volume nc\ follows the same Gaussian distribution as the first term of equation (2.5): nci

, ,„ , , , exp(z/#i)2

(^-7)

where nc;(0) = po(z = 0)/m c | is the cloud number density at midplane, z — 0. The cloud distribution then is specified by / m (0) = / , which is the fraction of mass in clouds at midplane. 3. Results and discussion

The parameters for the gaseous disk of the Galaxy vary as a function of the galactocentric radius and here we focus on the evolution at the solar circle. A more detailed discussion and other sets of different models are discussed by Silich et al. (1994a, 1994b). The 3-D shapes for superbubbles generated with a mean energy rate L = 1.05 x 1038 erg s" 1 and two different initial positions relative to the Galactic plane (z0 — 0 and 50pc) are shown in Figure 1. This figure illustrates the effects caused by the original source's relative location on the superbubble's final appearance. If the energy source is in the symmetry plane, the bubble remains confined within the disk. A displacement of about 50 pc in the z-direction ganerates a strongly asymmetrical structure, and the bubble breaks out of the Galactic disk. The most remarkable features are the peaks in the luminosity profiles. These peaks appear also in the X-ray luminosities (Silich et al. 1994a) and are a consequence of the inhomogeneity in the cloud distribution. To illustrate this conclusion we also show in Figures 2-4 the results of numerical simulations with the same input parameters but homogeneous gas and cloud density distributions (dashed lines). Thefinalparameters of the model are summarized in Table 1. Figure 2 shows the evolution of the total luminosity, Ltot, as a function of time. The absolute value of Ltot is approximate because we only use the average values of the thermodynamical quantities inside the cavity (see Silich et al. 1994a).

Silich et al.: Galactic Supershells

1500 -

1200 900

1200

600

900

300 600

-,

165

0

1

.— .

—

—

•

—

300

-300

r 1

—I

nn

7//f

0

-600

I

-900 -1200 7.8

4m — -—

_J

8.1

8.5 R(kpc)

.

8.1

%

-600

I

7.8

9.2

8.1

8.5 R(kpc)

8.8

9.2

FIGURE 1. Superbubble morphologies for different locations of the parent OB-association relative to the Galactic plane. Left: The OB-association is at midplane. Right: The OB-association is 50 pc above the plane. The simulations were done with £38 = 1.05 and at R - 8.5 kpc.

5

10

15

20

25

30

1 (Myr) FIGURE 2. The total superbubble luminosity as a function of time. The OB-association is at midplane. The dashed line represents the homogeneous (unstratified) gas and cloud space number density distributions. The rest of the parameters are the same as in Figure 1.

166

Silich et al: Galactic Supershells 5000

4000 -

3000 -

2000 -

1000-

0

5

10

15

20

25

30

t (Myr) FIGURE

3. The total number of clouds inside the remnant as a function of time for the cases described in Figure 2. 50.0

0.0 10

t(Myr)

FIGURE

4. The total evaporated mass inside the remnant as a function of time for the cases described in Figure 2.

4. Conclusions We have described the expansion of superbubbles in a cloudy medium for different positions of the energy source relative to the Galactic plane. Sources located off the plane generate highly distorted structures and the bubbles can break out of the Galactic disk. The total and X-ray luminosities display a different behavior with respect to cases with a homogeneous distribution. In the case of a stratified interstellar mass distribution the luminosity reaches a maximum at about 10 Myr and then becomes nearly constant. This is a consequence of the cloud concentration into a thin slab near the Galactic plane. The clouds engulfed by the expanding shell transfer some (3 - 6) • 104MQ to the bubble interior during the 30 Myr of evolution. This amounts to a few percent of the total mass of swept-up intercloud matter. SAS was partially supported by a grant from the International Science Foundation N1005/3. JF was partially supported by DGAPA-UNAM through the grant IN103991, a CRAY R&D grant, and the EEC grant CI1*-CT91-0935. JP was supported by the

Silich et al.: Galactic Supershells

167

grant No. 205/93/0090 of the Grant Agency of the Czech Republic. GTT was partially supported by the EEC grant CI1*-CT91-O935, DGICYT Grant PB91-0531(GEFE) and NATO Grant CRG920198 for collaborative research. Some of the simulations for this work were performed with the CRAY/YMP of the Supercomputing Center-UN AM.

REFERENCES G. S. & BLINNIKOV, S. I. 1982 Astion. Zh. 59, 876. BISNOVATYI-KOGAN, G. S. & SILICH, S. A. 1991 Astron. Zh. 68, 749. BRINKS, E. 1994 This volume. Cowm, L. L., MCKEE, C. F. & OSTRIKER, J. P. 1981 Astrophys. J. 247, 908. DICKEY, J. & LOCKMAN, F. 1990 Ann. Rev. Astron. Astrophys. 28, 215. FRANCO, J., BODENHEIMER, P., TENORIO-TAGLE, G. & ROZYCZKA, M. 1992 In Evolution of Interstellar Matter and Dynamics of Galaxies (ed. J. Palous, W. B. Buiton & P. O. Lindblad), p. 83. Cambridge University Press. BISNOVATYI-KOGAN,

FRANCO, J., MILLER, W., ARTHUR, S. J., TENORIO-TAGLE, G. &; TERLEVICH, R. 1994

As-

trophys. J. Submitted. HARTQUIST, T. W., DYSON, J. E. & SMITH, L.J. 1986

Mon. Not. R. astron. Soc. 221,

715.

HEILES, C. 1979 Astrophys. J. 229, 533. KLEIN, R. I., MCKEE, F. & COLELLA, P. 1994 Astrophys. J. In press. LOZINSKAYA, T. A. &; SITNIK, T. G. 1988 Pis'ma. Astron. Zh. 14, 240. MAC LOW, M.-M. & MCCRAY, R. 1988 Astrophys. J. 324, 776. MAC LOW, M.-M., MCKEE, C. F., KLEIN, R. I., STONE, J. &; NORMAN, M. 1994 Astrophys J. In press. OSTRIKER, J. P. & MCKEE, C. F. 1988 Rev. Mod. Phys. 60, 1.

J. 1990 IAU Symposium No. 144: The Interstellar Disk Halo Connection in Galaxies (ed. H. Bloemen), p. 101. Kluwer. PALOUS, J. 1992 In Evolution of Interstellar Matter and Dynamics of Galaxies (ed. J. Palous W. B. Burton & P. O. Lindblad), p. 65. Cambridge University Press. SILICH, S. 1992 Astrophys. Space. Sci. 195, 317. SILICH, S., FRANCO, J., TENORIO-TAGLE, G. & PALOUS, J. 1994a Numerical Simulations in Astrophysics (ed. J. Franco et al.). Cambridge Univ. Press. In press. SILICH, S., FRANCO, J., PALOUS, J. & TENORIO-TAGLE, G. 1994b In preparation. TENORIO-TAGLE, G. & BODENHEIMER, P. 1988 Ann. Rev. Astron. Astrophys. 26, 145. TENORIO-TAGLE, G. & PALOUS, J. 1987 Astron. Astrophys. 186, 287. ToMlSAKA, K. 1993 Numerical Simulations in Astrophysics (ed. J. Franco et al.) Cambridge Univ. Press. In press. PALOUS,

Violent Star Formation in Dwarf Irregular Galaxies By EVAN D. SKILLMAN Department of Astronomy, University of Minnesota, Minneapolis, MN 55455, USA Star formation in general and violent star formation in particular as observed in dwarf irregular galaxies are discussed. Emphasis is placed on those qualities of dwarf irregular galaxies that may be regarded as controversial. In particular, the conditions leading to star formation and the effects of star formation on the chemical and dynamical evolution of a dwarf irregular galaxy are discussed.

1. Introduction The title of this talk reflects a relatively broad reach, and I have no aspirations of achieving an all-encompassing review. Tofindsuch material, I recommend the reviews of Elmegreen (1992), Franco (1992), Hunter (1992), Kennicutt (1992), and Melnick (1992) in the proceedings of the Third Canary Islands Winter School of Astrophysics on Star Formation in Stellar Systems. Altogether, I believe that these lectures will provide an excellent background from which to discuss the problem of star formation in dwarf galaxies. Instead, I would like to take this opportunity to discuss some of what I consider to be the points of contention one might encounter in the more lengthy reviews. To get right to the point, my discussion will take the form of a presentation of my prejudices. In preparing this talk, I was able to assemble my prejudices and review the observations and theories that led to their development. To admit that these are prejudices, I think, allows them to be openly confronted by both myself and others.

2. Star formation in dwarf galaxies 2.1. A brief definition of terms One stumbling block encountered by everyone who works on dwarf galaxies is the lack of a coherent terminology and pervasive abuse of the existing terminology (Tammann 1994). For this talk, I will divide star forming dwarf galaxies into two cases, normal dls and dls with dominant starbursts. Seen in evolutionary snapshots, the normal dls experience star formation, as indicated by the presence of HII regions, which wanders over the face of the galaxy. My image of a prototype comes from the studies of NGC 6822 and IC 1613 by Hodge (1980) where star formation histories reconstructed from the clusters and associations reveal a migratory pattern. This is also seen in the smallest of dls like Sextans A (Aparicio et al. 1987) and GR 8 (Wyatt k Dufour 1993). In the smallest galaxies, the "cell size" for the star formation events is comparable to the size of the galaxy, so the galaxy will experience episodic star formation as envisioned by Gerola, Seiden, & Schulman (1980) in their picture of self propagating star formation. In larger galaxies, there are a sufficient number of cells for the star formation rate to be nearly constant with time, and Gallagher, Hunter, & Tutukov (1984) propose some observational evidence supporting this picture. Note, however, that for the LMC, which 168

Skillman: Violent Star Formation in Dwarf Irregulars

169

is a relatively large galaxy, there is very good evidence that the star formation history has not been constant in any sense (Bertelli et al. 1992). My picture of starburst dls requires a large, central burst of star formation which dominates the optical view of the galaxy (e.g. NGC 1569). The evolutionary history of these types of galaxies is not so clear, both in what were the prerequisites for the burst and in what will happen to the galaxy in the future (with possibilities ranging from returning to normal dl status to a complete disintegration of the galaxy). If we might compare the evolutionary histories of dls to the geological history of the island of La Palma, at one point the island would have been classified as a starburst during the eruption which left behind the Caldera de Taburiente but today the island would be classified as normal with perhaps a single giant HII region in the vicinity of Fuencaliente. Note that this discussion has concentrated on the optical appearance of the galaxy. Before continuing, I would like to introduce my first prejudice: Prejudice #1: Know Thy HI Actually this is the weak form of my prejudice; the strong form states that unless you have a synthesis HI image of the neutral hydrogen in a dwarf irregular galaxy, you don't know anything at all about that galaxy. The case behind this prejudice lies partly IN the rationale that you cannot have a complete picture of a galaxy unless you know the distribution of the ISM of the galaxy and partly from the fact that total masses and mass distributions of irregulars require HI imaging (Skillman 1992). Note that the second is not necessarily true for spiral galaxies. The optical isophotes of a spiral galaxy give a good indication of its inclination, so that it is possible to correct a single dish HI spectrum for inclination effects in order to divine the true height of the rotation curve. (The success of the Tully-Fisher method of distance measurement depends on this.) For the typical dl, only very deep surface photometry will reveal the ellipticity of the underlying stellar disk. The HI distribution and velocity field is a much more reliable guide to understanding the orientation of the galaxy. 2.2. Star formation thresholds By comparing their Ha imaging with the HI maps of the LMC, Davies et al. (1976) followed up on the work of McGee & Milton (1966) and confirmed a strong correlation of the high surface brightness HII regions with the regions of large HI column density. This observation led them to propose a surface threshold density for massive star formation in the LMC. Later, Gallagher & Hunter (1984), based on the very few HI observations of dls that were available at the time, proposed that a surface threshold density for star formation may be universal for dls and guessed that the threshold was about 5 x 1020 atom cm" 2 . Note that any HI column density measurement is resolution dependent. It is important that a linear resolution be given before comparing thresholds between galaxies. This is particularly important for distant galaxies where resolutions exceeding several kiloparsecs result in artificially low column densities. Note also, that resolution is much less important when discussing radially averaged HI distributions (van der Hulst et al. 1993). Skillman (1987) compared the HI distributions obtained from synthesis observations at a linear resolution of 500 pc with Ha imaging for seven dls and found a very strong correlation between the presence of HI over a column density of 10 21 atom cm" 2 and the presence of adjacent HII regions. It was also discovered that the HI sizes (at column densities of a few x 1019) could be close to the optical size of the galaxy or several times

170

Skillman: Violent Star Formation in Dwarf Irregulars 0

1.5

01 02 45

30 15 00 RIGHT ASCENSION (B1950)

3.0

01 45

30

1. The neutral hydrogen column density in the Local Group dwarf irregular galaxy IC 1613 compared with the distribution of HII regions. The HI is shown in greyscale with contours at 0.25, 0.5, 1, and 2 x 1021 atom cm" 2 . The image was constructed from VLA observations with a circular beamsize of 1 arcmin (FWHM) which corresponds to 210 pc at the distance of IC 1613. The HII region positions and sizes, which are represented as circles, were taken from Hodge e al. (1990). FIGURE

larger, but that this characteristic was unimportant with regard to the star formation properties of the galaxy. The only gas that was important to the massive star formation within the galaxy was the gas that exceeded the threshold. An excellent example of this is seen in IC 1613 (Figure 1) where the HI distribution (Lake & Skillman 1989) shows a single strong peak above 1021 atom cm" 2 and almost all of the bright HII regions are clustered around that HI peak while the rest of the HII regions are associated with HI above 5 x 1020 atom cm"2 (Hodge, Lee & Gurwell 1990). 2.2.1. Thresholds to molecular cloud formation? What is the physical cause underlying the observed star formation threshold? I know of two credible possibilities. It is possible that the HI column density is simply indicative of a sufficient dust column density to shield the neutral gas from the ambient radiation field allowing molecule formation to proceed as envisioned by Federman, Glassgold &; Kwan (1979). Under this scenario one would expect a metallicity dependence, since the dust/gas ratio should be dependent on metallicity (Franco & Cox 1986; Elmegreen 1989).

Skillman: Violent Star Formation in Dwarf Irregulars

171

This possibility assumes that the problem of forming stars is one of forming molecular clouds, and here we encounter a second prejudice: Prejudice #2: Star formation always occurs in molecular clouds There have been some controversial results concerning the presence of molecular clouds in actively star forming dls. Since the pioneering work of Elmegreen, Elmegreen & Morris (1980), it has been clear that the CO molecule is difficult to detect in dls and therefore that the CO surface brightnesses are much lower in dls than in spiral galaxies. One school of thought has held that the CO/H2 ratio is constant everywhere, and thus the molecular gas content of dls must be very low (Young, Gallagher & Hunter 1984; Tacconi & Young 1987). However, Israel et al. (1986) hypothesized that due to the lower metallicities of dls the dust/gas ratios could be lower and the ambient radiation field could be stronger in dls relative to the Galaxy. This would result in smaller CO core sizes for the same size H2 clouds. Why is the CO/H2 ratio dependent on metallicity and the local radiation field? Van Dishoeck & Black (1988) have described the photodissociation of CO and point to the fact that while both CO and H2 dissociation are dominantly line processes, self-shielding is much more important for H2 due to the factor of at least 104 abundance difference. This large difference in abundance also means that dust shielding (which is unimportant for H2) competes effectively with self-shielding for CO. In addition, the dissociation line widths of CO are broad because the dominant process is pre-dissociation as opposed to spontaneous radiative dissociation which dominates for H2. Thus, it is not surprising that the relative sizes of the CO cores of H2 clouds can vary, depending on local physical conditions. Indeed, Maloney & Black (1988) show dramatic differences in the fraction of C in the form of CO due to a small change (factor of 4) in C abundance. This picture is borne out by observations of the LMC and SMC where the CO/H2 ratio is shown to be l/6th and l/20th of the Galactic ratio (Cohen et al. 1988; Rubio et al. 1991). On the other hand, Sage et al. (1992) found values ranging from the Galactic value to l/3rd of the Galactic value for blue compact galaxies and Wilson <Si Reid (1991) found the ratio to be only 1/2 for IC 10. My prejudice was strengthened when I made a direct detection of strong H2 emission in the near infrared in the vicinity of Hubble V, a giant HII region in NGC 6822 (Israel 1988a,b) which previously showed no evidence of associated CO emission and only recently has been detected in very sensitive CO observations (Israel 1991, Wilson 1992). In addition to smaller CO cores, the H2 clouds may also be smaller. Elmegreen (1989) has argued this on theoretical grounds, and Rubio et al. (1991) note that the H2/HI ratio is 15 times smaller in the SMC compared to our Galaxy. The one dl that has not fitted into this picture has been LGS-3. Despite its very small luminosity, which would imply a very low metallicity (Skillman, Kennicutt &; Hodge 1989), LGS-3 was detected in CO by Tacconi & Young (1987). However, Lo et al. (1993) report that they have been unable to confirm the CO detection, and the original detection was reported with a radial velocity of 310 km s" 1 while the HI radial velocity is ~ 1 -285 km s" . 2.2.2. Thresholds to cloud collapse? The disk instability criterion for star formation which has been shown to be very convincing for spiral galaxies (Kennicutt 1989) may also be valid for dls. Since dls are dominated by solid-body rotation over their optical disks, the disk instability criterion produces a constant value of the star formation threshold across the entire disk of the

172

Skillman: Violent Star Formation in Dwarf Irregulars

galaxy. Using the solid-body rotation approximation and accounting for the presence of He, calculations of the critical surface density threshold typically result in values within a factor of two of 1 x 10 21 atom cm" 2 . For example, I find values of 2.0 x 10 21 for Sextans A (Skillman et al. 1987), 1.2 x 10 21 for IC 10 (Shostak k Skillman 1989), 5.4 x 1O20 for IC 1613 (Lake k Skillman 1989), 1.2 x 10 21 for DDO 154 (Carignan k Freeman 1988), and 1.0 x 10 21 and 4.4 x 10 20 for IC 3522 and 6°18 respectively (Skillman et al. 1987). There does not appear to be any trend in this theoretical surface-density threshold with mass, implying that one might be able to distinguish between the two hypotheses of cloud collapse and molecular cloud formation. One might ask if it is appropriate to treat dls as thin-disk systems, and here we find another prejudice: Prejudice #3: dls have HI in flattened, rotating disks Again we find some controversy on this point in the literature. Perhaps some early confusion arose when it was discovered that the single-dish HI profiles of most dls were Gaussian in shape and did not show the double-horned profile common in spiral galaxies. Of course, the profiles of spirals reflect the extended fiat part of the rotation curve that most dls lack. This, combined with the irregular optical appearance of dls led some to model dls as chaotic systems. Staveley-Smith, Davies k Kinman (1992) found the average intrinsic axial ratio of the starlight of dwarf galaxies to be 0.57, indicating some flattening. It seems reasonable to assume that the gaseous component will be more flattened than the stellar component. While it is clear that the gaseous disks of dls are not nearly as thin as those of the more massive spiral galaxies, modeling the gaseous disk of a dl as a pressure-supported sphere is certainly inappropriate. All dls, when observed in HI, show velocity fields of rotating systems. Again, this point is a bit controversial for the lowest-mass systems, where the rotation velocity is of order of the velocity dispersion of the gas. Recently, Lo et al. (1993) have presented VLA HI observations of 9 very low mass dls and find that only 2 of these 9 show a velocity gradient indicative of rotation. They state that these very low mass systems are dominated by chaotic motions rather than rotation. When I look at their velocity fields, I find evidence of rotation (with rotation curve amplitudes in excess of 5 km s" 1 ) in all but 2 of the systems. Simple statistics of inclinations would appear to explain the low observed rotation velocities in the other two. 2.3. Giant HII regions in dls Given the above picture of star formation in dls, one interesting problem is the formation of giant HII regions, regions comparable in their production of new stars to the 30 Dor region as reviewed by Kennicutt (this volume). It would appear that the problem of forming a giant HII region is that of forming a giant gas cloud capable of sustaining star formation on this scale. Again we might consider two options, that the giant HII regions represent the upper end of a normal distribution of sizes of star-forming regions, or that the production of giant HII regions requires special conditions. I think that the former case has been effectively ruled out by the studies of Hodge and collaborators (cf. Strobel, Kennicutt & Hodge 1991). Histograms of HII region sizes of dls can be characterized as exponential in nature, but the giant HII regions which are frequently found in dls fall well outside of these normal distributions (Hodge 1983). The locations of giant HII regions often suggest special conditions. Many, like 30 Dor, are associated with bar structure (Elmegreen k Elmegreen 1980). It seems likely that

Skillman: Violent Star Formation

in Dwarf Irregulars

173

bars are establishing flows which pile up the gas required for these m o n u m e n t a l starforming events.

3. Prerequisites for violent star formation Here I will discuss the starburst dls which often fall under the label of Blue Compact Galaxies and HII galaxies. 3.1. HII galaxies with companions

I need to digress a bit from the main theme of this talk in order to give some background for some observations I will present. This digression consists of a very brief description of the Ph.D. thesis work of Chris Taylor done in collaboration with Eli Brinks and me. Based on a study of a small sample of HII galaxies, Brinks (1990) hypothesised that interactions with other dwarf galaxies could be responsible for triggering the current burst of star formation. Using this hypothesis as the basis for a search for extragalactic HI clouds, Taylor, Brinks & Skillman (1993) made VLA HI observations to search the vicinity of nine HII galaxies in right ascension, declination, and redshift for HI companions. Four companion objects were found in this pilot sample. A follow-up study of a statistically well-defined sample from the lists of Salzer, MacAlpine & Boroson (1989a,b) resulted in detections of companions in all but 4 of 19 targets. Some of the companion clouds have readily detectable optical counterparts, but a few have yet to be detected at optical wavelengths and might be the closest things to "Intergalactic HI" that we currently know of (Brinks 1994; Brinks & Taylor 1994). The goal offindingextragalactic HI clouds may have been met, but the hypothesis that the starbursts in HII galaxies are triggered by encounters is another matter altogether. Of course this topic already has a long history, although most of the literature concerns larger galaxies. Larson & Tinsley (1978) were the first to point out that interactions and mergers were accompanied by larger current star formation rates, but even the most extensive studies (e.g., Kennicutt et al. 1987) have very few dwarf galaxies in their samples. Models of interacting galaxies with gaseous disks have shown that gas can be driven to the center (Negroponte & White 1983) which presumably is followed by increased star formation rates (Noguchi & Ishibashi 1986; Noguchi 1987; Hernquist 1989; Barnes & Hernquist 1991). Again, I would note that these models are based on spiral-like galaxies. I know of no simulations of interacting dwarf galaxies. The high detection rate of companions to HII galaxies may be validating the hypothesis of bursts of star formation triggered by interactions. The fact that we don't see companions to all of the HII galaxies could point to triggering by gas-poor dwarfs in a few cases. On the other hand, a suitable "control" sample has yet to be established so that we might know the statistical significance of our high companion detection rate. I think that it is also worth noting that our detection rate may not be the correct observable to focus on; from modeling calculations, Mihos, Richstone & Bothun (1991, 1992) find that only a small fraction of encounters result in a large increase in the star formation rate. If the results of these models (again carried out on spiral-type galaxies) are relevant to dwarf galaxies, then to test the hypothesis requires understanding the nature of the interaction (e.g. prograde versus retrograde). 3.2. Star formation in HII galaxies The original HI observations of Taylor et al. (1993) were combined with deeper, higher resolution HI observations and Ha imaging in order to study the star formation properties

174

Skillman: Violent Star Formation in Dwarf Irregulars

of the galaxies (Taylor et al. 1994). At first, the HI observations revealed no surprises and fit in well with the picture of normal dls as given in the first section of this talk. All of the galaxies shared three characteristics: (1) The radial HI profiles were strongly centrally peaked. Note that this is in comparison to other low surface brightness galaxies where the radial HI profiles are relatively flat (van der Hulst et al. 1993). (2) The central HI maxima exceeded the star formation threshold criteria as derived by Skillman (1987) and Kennicutt (1989). (3) Using the disk instability criterion of Kennicutt, a star formation threshold contour was identified for each of the galaxies and this contour encompassed the high surface brightness patch of recent star formation in all cases. All three of these observations are consistent with the picture as sketched in the previous section. However, there was one surprise in the data: (4) All HII galaxies showed large HI sizes when compared with their optical size. This was not necessarily expected, since the large envelopes of HI are not participating in the current burst of star formation. When combined with the small sample of other HII galaxy prototypes with existing synthesis HI observations, a pattern appeared. NGC 1705, NGC 4214, NGC 4449, and NGC 5253 are all HII galaxies with large HI envelopes. This prompts the following extended hypothesis for starburst irregulars, that they not only require an interaction but that they also require a large HI halo. Since the large central HI overdensities are not found in typical dls, it would appear that the role of the triggering interaction is to build that central overdensity (as opposed to a picture wherein galaxies with pre-existing central HI concentrations are waiting for interactions to trigger the star formation). In sum:

Prejudice #4-' Starburst dls require both pre-existing conditions and a trigger A comparison of the luminosities and HI masses of low surface brightness galaxies and HII galaxies indicates that some low surface brightness galaxies could serve as the hosts to HII galaxies (Taylor et al. 1994). I think that Figure 2 bears some inspection. Note that the very low mass systems studied by Lo et al. (1993) fall off to the left of the diagram with HI masses less than 108 M 0 . Although the linear sizes of the optical disks of the lowest luminosity systems and the HII galaxies may be comparable (~ a few kpc), their total HI masses are very different. I interpret this to mean that the lowestmass dwarf galaxies cannot serve as hosts to the very luminous bursts of star formation as observed in HII galaxies. This could be a misconception arising from a preference to conduct HI sythesis observations of galaxies with relatively large HI fluxes. Salzer (private communication) has found many Blue Compact Galaxies with relatively small HI masses. 4. The effects of violent star formation on dls Brinks (this volume) has shown us a number of cases in which it is clear that star formation is having a tremendous impact on the ISM in galaxies. Since dls define the low-mass end of star forming galaxies, it is natural to consider that star formation could have devastating effects on a galaxy. In the past it has been very popular to discuss these effects in terms of a wind. I would like to ask, "Are winds important for dls?" This discussion is split into two sections, the first concentrating on chemical evolution concerns and the second concentrating on dynamical concerns. Very importantly, I would

175

Skillman: Violent Star Formation in Dwarf Irregulars 11

i

i

i

i—I

i

i

i

i—I

I

i

I

I

I

I

i

i

•

10

r

VII Zw 8

* Mrk314 A Haro 26

rP A

Mrk 600 D

o

D

o

D

CD

3

. O A W794

_o I

I

I

I

I

8.5

I

I

I

I

I

I

I

I

I

II

9 9.5 Log M(HI) solar masses

I

I

I

A a

HII galaxy LSB spiral

• O

HSB spiral — LSB dwarf -

1

i

10

i

i

-

i

10.5

FIGURE 2. LogL(B) versus logM(HI) for HII galaxies (filled triangles; Taylor et al. 1994), LSB galaxies (open squares; van der Hulst et al. 1993), HSB galaxies (filled pentagons; sample from Kennicutt 1989), and LSB dwarfs (open circles; Skillman et al. 1987; Carignan & Freeman 1988; Carignan et al. 1988; Lake et al. 1990). Note that the HII galaxies have masses more comparable to the spiral galaxies than the LSB dwarfs, despite having optical diameters comparable to the LSB dwarfs.

like to note up front that I am discussing present-day dls. The very early evolution of dwarf galaxies is an important topic, but one that has few observational constraints. 4.1. Galactic winds and chemical evolution: normal dls The idea of metal-enhanced winds playing a dominant role in the evolution of dwarf galaxies is very appealing. Larson (1974) noted that metal-enhanced winds could explain the lack of prominent nuclei in low-mass systems, the lower average stellar abundances in low mass systems, and their smaller radial abundance gradients. Metal-enhanced winds could easily explain the luminosity-metallicity relationship for dls first noted by Lequeux et al. (1979). The selective loss of metals, particularly oxygen, would also resolve the discrepancy between the yields calculated from stellar evolution theory and those measured from observations of dls (Pagel et al. 1992; Maeder 1992; Peimbert, Colin & Sarmiento, this volume). What is the form of these galactic winds? Given the strong motivation for their presence, it is not surprising that one can find several models for winds in dwarfs. Larson (1974) was interested in the possibility that the collective powers of supernovae could sweep out the ISM of an elliptical galaxy, and pointed out the relevance of this mechanism for dwarf galaxies. In much the same way, Dekel & Silk (1986) developed a picture of dwarf galaxies whose earlier histories are dominated by a wind phase in which up to 90% of the mass of the galaxy is lost. In this picture, dls are different from dEs only in their abilities to retain a small fraction of their gas. Vader (1986) emphasizes that the winds must be metal enhanced in order to explain the low metallicities of the dwarf

176

Skillman: Violent Star Formation in Dwarf Irregulars

galaxies. Pantelaki k Clayton (1987) and, later, Clayton k Pantelaki (1993) give a more detailed model in which the evolution of chemically homogeneous, cold, neutral inclusions within a hot, metal enhanced medium are considered. Burkert k Hensler (1989) present a similar model. Such large scale winds may be very important in the early history of dwarf galaxies, especially following a first burst of star formation. However, is there sufficient evidence to support the hypothesis that winds dominate the chemical evolution of present-day dwarf irregular galaxies? I think one pertinent question is what fraction of the gas actually escapes following an episode of star formation? The answer to this lies in understanding the evolution of the bubbles formed around star forming regions (see review by Tenorio-Tagle k Bodenheimer 1988). The evolution of an individual bubble is governed by the energy input rate and the mean density of the surrounding gas. Obviously, insufficient energy input (too few stars and/or supernovae) will not result in break-out from the cold disk. Additionally, if the energy input rate is too low, a larger fraction of the energy will be lost radiatively, also resulting in failure to break-out. The case of break-out is described by Tomisaka k Ikeuchi (1986) and MacLow, McCray k Norman (1989). De Young k Gallagher (1990) presented models of such bubbles forming in low-mass galaxies. They note that the low metallicities of low-mass galaxies will affect the cooling rates, resulting in larger bubbles for identical input energies. In their models, single SNRs do not break out, but regions of multiple SNe do. Until break-out, the metal-rich SNR debris resides near the inner boundary of the bubble. Upon break-out, approximately one sixth of this material actually escapes the galaxy. However, the most important effect for chemical evolution follows break-out. The metal-enhanced production of all future SNe is funneled out so that the final ratio of captured SN production to lost SN production is one third. These results are, of course, very model dependent, and therefore quite uncertain. Most importantly, De Young k Gallagher note that a dark matter halo was not included in their models, and the presence of one would decrease the loss. Additionally, a gas scale height of 150 pc was chosen, which may be on the low side for dwarfs (a larger gas scale height will delay break-out). It would be very interesting to see the results of models of this type for a larger range of galaxy parameters. What can we learn from observations? Two parameters must be identified, the fraction of massive stars that are born in regions which will break out and the fraction of the newly synthesized elements that will be lost to the galaxy as a result of the break-out. Estimating the first quantity will be very difficult, but perhaps studying the distribution of HII region sizes will yield a valid estimate. For the second parameter, we can consider the laboratory for star formation, Ho II (Puche et al. 1992). Here we see that the neutral gas has been shaped by the star formation events, leaving a legacy of holes and/or bubbles. Brinks (private communication) estimates that roughly one sixth of the hole structures have broken out of the disk. This might be taken as a lower limit to the fraction of bubbles that break out since certainly some of the present hole structures will grow sufficiently to reach break-out. To estimate the fraction of newly synthesized elements that are lost during break-out, two extremes might be considered. If all of the SNe have exploded before break-out from the disk, then only the caps of the spheres will actually be propelled away from the disk of the galaxy. These caps may represent less than 20% of the surface area of the shells. If a significant fraction of stars are born in clusters that do not produce bubbles that proceed through to break-out, it could be that only a very small fraction of the gas actually leaves the galaxy (even in a galaxy in which the HI distribution is dominated

Skillman: Violent Star Formation

in Dwarf Irregulars

177

by holes and bubbles). T h e other extreme is given by the theoretical models indicating close to unity efficiency of loss of metals is possible. I think t h a t the nonthermal superbubble in IC 10 (Yang & Skillman 1993) presents another observational case supporting a minor role for winds with regard to chemical evolution. Here we see a giant supershell structure which shows no sign of breaking through the disk and hasn't even disrupted the adjacent HI cloud (Shostak & Skillman 1989). Since much of the motivation for strong winds comes from chemical evolution concerns, I would like to note two points against galactic winds playing a dominant role in the chemical evolution of dls. First, galactic winds as pictured by Larson (1974) and Dekel & Silk (1986) have an upper mass limit, above which they are no longer i m p o r t a n t . T h u s , the luminosity-metallicity relationship which is quite strong in irregular galaxies (Skillman, Kennicutt k Hodge 1989) should be different from t h a t seen in the high mass spiral systems (Garnett & Shields 1987). In fact, Zaritsky, Kennicutt & Huchra (1994) have shown t h a t the low-mass dls and the high mass spirals all fit onto the same luminosity-metallicity relationship. Secondly, relative abundance patterns of dls may argue against dominant galactic winds. If galactic winds dominate the chemical evolution of dls, then C and N will be less affected (compared to 0 ) . Naively, low-mass galaxies, which have the most trouble retaining 0 , should have higher N / O and C / O ratios. In fact, N / O is fairly constant for the lowest mass systems (Garnett 1990) and C / O increases with increasing metallicity ( G a r n e t t et al. 1994). This may reflect a balance between the diminishing effect of galactic winds with increasing mass and the C, N, and 0 yields changing as a function of metallicity. To test this, abundance observations of a larger sample will be needed. In particular, HI synthesis observations of these systems are required in order to compare the abundances with the gas mass fractions (properly accounting for the effects of extended HI haloes and dark m a t t e r haloes). These observations have led me to the following: Prejudice

#5:

Winds are not important

for the chemical evolution

4.2. Galactic winds and dynamical

evolution:

starburst

of dls

dls

In some starburst dls, there is good evidence for a wind. N G C 1569 is the best example t h a t I know of (Israel k, van Driel 1990). Here the superposition of blue, high surface brightness stellar clusters with a hole in the systemic velocity HI and associated high velocity HI clearly points to a star formation driven wind. Here the wind is probably very i m p o r t a n t for the chemical evolution of the galaxy. T h e nucleosynthetic products of the recent starburst are being carried away by the wind. Note, however, t h a t the region of disruption is local and confined. Most of N G C 1569 is not aware t h a t something special is taking place in its center. T h e dynamical evolution will not be affected dramatically by the star formation event. A large HI halo will not be blown away by this wind. T h e galaxy is in no danger of disintegrating or transforming into a dwarf elliptical. Similar arguments can be made for N G C 1705 (Meurer et al. 1992; Meurer, private communication). Which means:

Prejudice #6: HII Galaxies do not become dEs via galactic winds

If all HII galaxies have large HI haloes, then the present violent star formation, which

178

Skillman: Violent Star Formation

in Dwarf

Irregulars

will have d r a m a t i c affects in the vicinity of the optical disk, will not be able to drive away even a large fraction of the associated neutral gas. Present day HII galaxies will not evolve into dEs via galactic winds. Binggeli, Tarenghi & Sandage (1990) have enumerated several more objections to t h e proposition t h a t dEs could be formed by removing gas from dls, and Binggeli & Cameron (1991) have collected arguments in favor of gas removal a t early epochs as t h e fundamental process required to produce a d E .

5. Conclusions There are a number of still unsolved problems concerning the evolution of dl galaxies which can be approached through observational programs. I have listed a number of my own prejudices on several important questions, but I think that the very act of recognizing them as prejudices reveals their tenuous nature. On the optimistic side, since giving this talk, many researchers have let me know that they share several of my prejudices. Perhaps for some of these items of contention the weight of observational evidence is driving a consensus. The debate concerning the importance of galactic winds for dl evolution is likely to remain lively. In the future, I think that obtaining HI synthesis observations of otherwise well studied dwarf irregulars will help to constrain the many possible evolutionary scenarios. A perusal of the VLA observing schedules for the last few years indicates that we can expect an abundance of new HI observations of dls to be available. I also have the impression that arguments related to relative chemical abundances have not played as strong role in developing evolution scenarios as they might. I would like to thank the conference organizers for producing such a pleasant conference. I am grateful to Elias Brinks, Jay Gallagher, Don Garnett, Paul Hodge, Bernard Pagel, John Salzer, and Chris Taylor for comments on a draft of this paper and for many useful discussions with conference participants. I would also like to acknowledge the use of the Peridier Library at the University of Texas and thank all of those who have worked to maintain this wonderful facility. Partial support for this work has been provided by NASA LTSARP grant No. NAGW-3189.

REFERENCES APARICIO,

A.,

GARCIA-PELAYO,

J. M., MOLES, M. & MELNICK, J. 1987 A. & A. Suppl. 71,

297. BARNES, J. E. & HERNQUIST, L. E. 1991 Ap. J. Lett. 370, L65. BERTELLI, G., MATEO, M., CHIOSI, C. &: BRESSAN, A. 1992 Ap. J. 388, 400. BINGGELI, B., TARENGHI, M. & SANDAGE, A. 1990 A. & A. 228, 42. BINGGELI, B. &; CAMERON, L.M. 1991 A. & A., 252, 27.

BRINKS, E. 1990 In Dynamics and Interactions of Galaxies (ed. R. Wielen), P. 146. Springer. BRINKS, E. 1994 In The Cold Universe (ed. T. Montmerle, C. J. Lada & J. Tran Tranh Van), Editions Frontieres. In press. BRINKS, E. & TAYLOR, C. L. 1994 In Dwarf Galaxies (ed. G. Meylan). ESO. In press. BURKERT, A. &; HENSLER, G. 1989 In Evolutionary Phenomena in Galaxies (ed. J. Beckman & B. E. J. Pagel), P. 230. Cambridge University Press. CARIGNAN, C. & FREEMAN, K. C. 1988 Ap. J. Lett. 332, L33. CARIGNAN, C , SANCISI, R. & VAN ALBADA, T. S. 1988 A. J. 95, 37.

CLAYTON, D. D. &; PANTELAKI, I. 1993 Phys. Rep. 227, 293.

Skillman: Violent Star Formation in Dwarf Irregulars

179

COHEN, R. S., DAME, T. M., GARAY, G., MONTANI, J., RUBIO, M. &; THADDEUS, P. 1988

Ap. J. Lett. 331, L95. DAVIES, R. D., ELLIOT, K. H. k. MEABURN, J. 1976 Mem. R. A. S. 81, 89. DEKEL, A. & SILK, J. 1986 Ap. J. 303, 39.

DE YOUNG, D. S. &; GALLAGHER, J. S. 1990 Ap. J. Lett. 356, L15. ELMEGREEN, B. G. 1989 Ap. J. 338, 178.

B. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto &; F. Sanchez), p. 381. Cambridge University Press.

ELMEGREEN,

ELMEGREEN, B. G., ELMEGREEN, D. M. &; MORRIS, M. 1980 Ap. J. 240, 455. ELMEGREEN, D. M. &; ELMEGREEN, B. G. 1980 A. J. 85, 1325. FEDERMAN, S. R., GLASSGOLD, A. E. & KWAN, J. 1979 Ap. J. 227, 466.

J. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto &; F. Sanchez), p. 515. Cambridge University Press.

FRANCO,

FRANCO, J. & Cox, D. P. 1986 P. A. S. P. 98, 1076. GALLAGHER, J. S. & HUNTER, D. A. 1984 A. R. A. A. 22, 37. GALLAGHER, J. S., HUNTER, D. A. & TUTUKOV, A. V. 1984 Ap. J. 284, 544. GARNETT, D. R. 1990 Ap. J. 363, 142. GARNETT, D. R. & SHIELDS, G. A. 1987 Ap. J. 317, 82. GARNETT, D. R., SKILLMAN, E. D., DUFOUR, R. J., PEIMBERT, M., TORRES-PEIMBERT, S.,

SHIELDS, G. A., TERLEVICH, E. & TERLEVICH, R. J. 1994 In Dwarf Galaxies (ed. G.

Meylan). ESO. In press. GEROLA, H., SEIDEN, P. E. & SCHULMAN, L. S. 1980 Ap. J. 242, 517.

HERNQUIST, L. 1989 Nature 340, 687. HODGE, P. W. 1980 Ap. J. 241, 125. HODGE, P. W. 1983 A. J. 88, 1323. HODGE, P. W., LEE, M. G. & GURWELL, M. 1990 P. A. S. P. 102, 1245.

D. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto & F. Sanchez), p. 67. Cambridge University Press. ISRAEL, F. P. 1988a In Millimetre and Submillimetre Astronomy (ed. R. D. Wolstencroft & W. B. Burton), p. 281. Kluwer. ISRAEL, F. P. 1988b In Molecular Clouds in the Milky Way and External Galaxies (ed. R. L. Dickman, R. L. Snell & J. S. Young), p. 428 . Springer. ISRAEL, F.P. 1991 In IAU Symp. No. 146: Dynamics of Galaxies and their Molecular Clouds Distributions (ed. F. Combes & F. Casoli), p. 13. Kluwer, HUNTER,

ISRAEL, F. P., D E GRAAUW, T H . , VAN DE STADT, H. &; DE VRIES, C. P. 1986 Ap. J. 303,

186. ISRAEL, F. P. & VAN DRIEL, W. 1990 A. & A. 236, 323.

KENNICUTT, R. C. 1989 Ap. J. 344, 685. KENNICUTT, R. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto & F. Sanchez), p. 191. Cambridge University Press. KENNICUTT, R. C , KEEL, W. C , VAN DER HULST, J. M., HUMMEL, E. &; ROETTIGER, K.

A. 1987 A. J. 93, 1011. LAKE, G., SCHOMMER, R. A. &: VAN GORKUM, J. H. 1990 A. . 99, 547. LAKE, G. & SKILLMAN, E. D. 1989 A. J. 98, 1274. LARSON, R. B. 1974 M. N. R. A. S. 169, 229. LARSON, R. B. & TINSLEY, B. M. 1978 Ap. J. 219, 46. LEQUEUX, J., PEIMBERT, M., RAYO, J. M., SERRANO, A. &; TORRES-PEIMBERT, S. 1979 A.

& A. 80, 155. Lo, K. Y., SARGENT, W. L. W. & YOUNG, K. 1993 A. J. 106, 507.

180

Skillman: Violent Star Formation in Dwarf Irregulars

MAEDER, A. 1992 A. & A. 264, 105. MACLOW,

M.-M.,

MCCRAY,

R. & NORMAN, M. L. 1989 Ap. J. 337, 141.

M C G E E , R. X. & MILTON, J. A. 1966 Aust. J. Phys. 19, 343. MALONEY, P. & BLACK, J. H. 1988 Ap. J. 325, 389.

MELNICK, J. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto & F. Sanchez), p. 253. Cambridge University Press. MEURER, G. R., FREEMAN, K. C., DOPITA, M. A. & CACCIARI, C. 1992 A. J., 103, 60.

Mraos, J. C ,

RICHSTONE,

D. O. & BOTHUN, G. D. 1991 Ap. J. 377, 72.

MIHOS, J. C , RICHSTONE, D. O. & BOTHUN, G. D. 1992 Ap. J. 400, 153. NEGROPONTE, J. & WHITE, S. D. M. 1983 M. N. R. A. S. 205, 1009. NOGUCHI, M. 1987 M. N. R. A. S. 228, 635. NOGUCHI, M. 1991 M. N. R. A. S. 251, 360. NOGUCHI, M. & ISHIBASHI, S. 1986 M. N. R. A. S. 219, 305. PAGEL, B. E. J., SIMONSON, E. A., TERLEVICH, R. J. & EDMUNDS, M. G. 1992 M. N. R. A.

S.. 255, 325. I. & CLAYTON, D. D. 1987 In Starbursts and Galaxy Evolution (ed. T. X. Thuan, T. Montmerle & J. Tran Thanh Van), p. 145. Editions Frontieres.

PANTELAKI,

PUCHE, D., WESTPFAHL, D., BRINKS, E. & ROY, J.-R. 1992 A. J. 103, 1841. RUBIO, M., GARAY, G., MONTANI, J. & THADDEUS, P. 1991 Ap. J. 368, 173. SAGE, L. J., SALZER, J. J., LOOSE, H.-H. & HENKEL, C. 1992 A. & A. 265, 19. SALZER, SALZER,

J. J., J. J.,

MACALPINE, MACALPINE,

G. M. & BOROSON, T. A. 1989a Ap. J. Suppl. 70, 447. G. M. & BOROSON, T. A. 1989b Ap. J. Suppl. 70, 479.

SHOSTAK, G. S. & SKILLMAN, E. D. 1989 A. & A. 214, 33.

E. D. 1987 In Star Formation in Galaxies (ed. C. J. Lonsdale Persson), p. 263. NASA. SKILLMAN, E. D. 1992 In Elements and the Cosmos (ed. M. G. Edmunds and R. J. Terlevich), p. 21. University of Cambridge Press.

SKILLMAN,

SKILLMAN, E. D., BOTHUN, G. D., MURRAY, M. A. & WARMELS, R. H. 1987 A. & A. 185,

61. SKILLMAN, E. D., KENNICUTT, R. C. & HODGE, P. W. 1989 Ap. J. 347, 875. SKILLMAN, E. D., TERLEVICH, R. J., TEUBEN, P. J. & VAN WOERDEN, H. 1988 A. & A. 198,

33. STROBEL, N. V., HODGE, P. W. & KENNICUTT, R. C. 1991 Ap. J. 383, 148. STAVELEY-SMITH, L., DAVIES, R. D. & KINMAN, T. D. 1992 M. N. R. A. S. 258, 334. TACCONI, L. J. & YOUNG, J. S. 1987 Ap. J. 322, 681. TAMMANN,

G. A. 1994 In Dwarf Galaxies (ed. G. Meylan). ESO. In press.

TAYLOR, C. L., BRINKS, E. & SKILLMAN, E. D. 1993 A. J. 105, 128. TAYLOR,

C. L.,

BRINKS,

E.,

POGGE,

R. W. & SKILLMAN, E. D. 1994 A. J. 107. In press.

TENORIO-TAGLE, G. & BODENHEIMER, P. 1988 A. R. A. A. 26, 145. TOMISAKA, K. & IKEUCHI, S. 1986 P. A. S. J. 38, 697. VADER, J. P. 1986 Ap. J. 305, 669.

VAN DER HULST, J. M., SKILLMAN, E. D., SMITH, T. R., BOTHUN, G. D., McGAUGH, S. S. & DE BLOK, W. J. G 1993 A. J. 106, 548. VAN DISHOECK, E. F. & BLACK, J. H. 1988 Ap. J. 334, 771. WILSON, C. D. 1992 Ap. J. 391, 144.

WILSON, C. D. & REID, I. N. 1991 Ap. J. Lett. 366, L l l . WYATT, R. J. & DUFOUR, R. J. 1993 R. M. A. A. 27, 213. YANG, H. & SKILLMAN, E. D. 1993 A. J. 106, 1448.

Skillman: Violent Star Formation in Dwarf Irregulars YOUNG, J. S., GALLAGHER, J. S. & HUNTER, D. A. 1984 Ap. J. 276,

476.

ZARITSKY, D., KENNICUTT, R. C. & HUCHRA, J. P. 1994 Ap. J. 420,

87.

181

Very Metal-Poor Galaxies and the Primordial Helium Abundance By ELENA TERLEVICH1!, EVAN SKILLMAN2 AND ROBERTO TERLEVICH1 'Royal Greenwich Observatory, Madingley Road, Cambridge, CB3 OEZ, UK 2

Astronomy Department, University of Minnesota, 116 Church Street SE, Minneapolis, MN 55455, USA

Critical to the understanding of several fundamental problems in astronomy (among which the determination of the primordial helium is of foremost importance), extremely metal-poor galaxies have been almost impossible to find. In the past few years we have been successful in discovering them. We are embarked on a programme for obtaining with linear detectors, very high S/N spectra of these objects, in order to derive He abundances to better than the 5% per object needed to constrain the Big Bang model of the origin of the universe. We will discuss some results and problems encountered in this quest.

1. I n t r o d u c t i o n Three observational findings sustain the Big Bang model of the origin of the Universe: 1) the relic 3-K microwave background radiation; 2) the expansion of the Universe; and 3) the relative abundances of the light elements (H, D, 3 He, 4 He, Li); for a review, see e.g. Walker ei al. (1991). Accurate measurements of these observables provide better constraints on the details of the model. Even though an accurate measurement of the primordial light-element abundance and in particular of He, is critical to our understanding of the origin of the Universe, it has not been the target of an observational effort comparable to that for the other cosmological observables: the microwave background and the expansion of the universe. Experimental progress has been made on the number of neutrino species and the halflife of the neutron, which together with the nucleon mass density, relate directly to the relative abundances of the light elements just after the Big Bang. Therefore, a determination of the primordial helium abundance, under the assumption of the standard hot Big Bang model of nucleosynthesis, will eventually provide a fundamental parameter in cosmology: the nucleon density in the early universe. The process for measuring helium abundance and hence inferring the primordial value is discussed in §2; uncertainties in the method are shown in §3, general results in §4 and conclusions are drawn in §5.

2. Measuring the primordial helium abundance 2.1. The method Empirical evidence of the existence of primordial He, that is the He that was present before the heavy metals were synthesized in the interior of stars, comes from the early discovery of a plateau in the relation of the abundance by mass of He versus metals (Y vs. Z) in extragalactic HII regions and HII galaxies (Figure 1). Y tends to a finite value as Z tends to zero. Therefore one can gauge the relative contributions of primordial and t Present address: Institute of Astronomy, Madingley Road, Cambridge CB3 OHA, U. K. 182

Terlevich el al.: Primordial Helium Abundance

183

500

250

200

°

150 100

50

0^ 0

20

40

60

SO

100 120 140 160 180 200 220 240 0 / H (106)

FIGURE 1. Abundance by mass of He (Y) vs. O/H for a sample of low-metallicity giant HII regions and HII galaxies taken from the literature. Typical error bars are indicated. Notice the plateau in the relationship.

post-primordial helium by measuring the abundance of a heavier element (e.g. oxygen) that was not produced in the Big Bang. The most accurate method of measuring helium abundance consists of observing diffuse gas ionized by the ultraviolet photons of an inbedded cluster of hot stars in the so-called giant HII regions. The goal is to measure the fraction of helium to hydrogen present in ionized gas of low heavy-element abundance (in practice, oxygen). Early studies (Peimbert & Torres-Peimbert 1974, 1976; Lequeux et al. 1979) pioneered the method of extrapolating the He/H vs. O/H relation to a value of the helium abundance where the oxygen abundance is zero. The smaller the oxygen abundance, the smaller the extrapolation, the more accurate the measurement of the primordial helium abundance, and, finally, the more severe the constraint on the Big Bang model. This method has been refined by Pagel and his collaborators in a series of papers culminating with Pagel et al. (1992), for which they only used best quality data of V vs. O/H and also Y vs. N/H for high excitation, high surface brightness HII regions disregarding those showing any sign of contamination by Wolf-Rayet stars. Figure 2 serves as an illustration. Clearly, it is critical to find objects of very low oxygen abundance. 2.2. The search for very low abundance objects

Unfortunately, several attempts to find such regions have failed. Candidates have been obtained amongst the highest excitation objects from surveys capable of detecting intense

Terlevich et al.: Primordial Helium Abundance

184

270

260 -

250 -

240 -

230 -

220 20

40

270

80 100 120 140 160 180 200 220 240 106 0/H

60

i

•

i

r

^

t

260

250

240

230 (b) 220

i

0

20

40

i

i

60

i

i

I

»

I

,

I

SO 100 120 140 160 180 200 220 240

1 0 7 N/H FIGURE

2. Abundance by mass of (top) Y vs.O/H and (bottom) Y vs. N/H, for a sample of low-metallicity giant HII regions and HII galaxies from Pagel et al. (1992).

Terlevich et al.: Primordial Helium Abundance

185

bursts of star formation. Until 1985, only one object of suitably low abundance was known (I Zw 18, an object with an oxygen abundance of only 2 percent of that of the sun), and no one seemed to know how to find more (Shields 1986). It has been known for more than a decade that low-mass irregular galaxies are much less evolved chemically than our own galaxy (Lequeux et al. 1979; Talent 1980); in the last few years, searches for very low metal-abundance regions in extremely low-mass galaxies have been very successful (Skillman et al. 1988; Skillman, Kennicutt & Hodge 1989; Izotov et al. 1992). Skillman et al. (1989) compare the histogram of number of objects versus their oxygen abundance for three different studies: a determination of the primordial helium abundance (Kunth k. Sargent 1983); the best candidates (Campbell, Terlevich & Melnick 1986), chosen from a survey of over 800 objects (Terlevich et al. 1991) and the objects discovered by Skillman, Terlevich & Melnick (1989) restricting the search to extremely low-mass galaxies, proving that the latter clearly represents a superior way of finding low-abundance objects. Searching for these metal-poor galaxies Stepanian & Terlevich (private communication) have compiled a list of low-luminosity candidates from the Second Byurakan Survey (SBS) of ultraviolet excess galaxies, which have been studied spectroscopically by Lipovetskii and Stepanian. We obtained blue spectra of these galaxies with the 2.5-m Isaac Newton Telescope (INT) in order to select a subsample of the objects that are the brightest, lowest oxygen abundance and of highest excitation. This was followed up by high precision, high spatial and spectral resolution, high signal-to-noise spectroscopy with linear detectors with the 4.2-m William Herschel Telescope (WHT) in La Palma (Skillman et al. 1994; Skillman, Terlevich fc Terlevich, in preparation). ISIS, the double spectrograph at the WHT, allows simultaneous observations in the blue (3600-5100 A) and red (6300-6800 A) regions of the spectrum. Figure 3 shows a typical spectrum. While we have discovered about ten objects with oxygen abundances in the range 2-5% of the solar value, none has an oxygen abundance lower than that of I Zw 18. 2.3. Abundance derivation Ionized gaseous nebulae are fairly well understood, and their physical parameters (temperature, density, chemical abundance) can be deduced by analysing relative intensities of strong emission lines (Aller 1984; Osterbrock 1989). Electron densities (ne) for our sample were obtained from the ratio of the [SII] AA6717, 6731 A or [Oil] AA3726, 3729 A emission lines, corresponding in all cases to the lowdensity limit (ne ~ 100 cm" 3 ). Electron temperatures for the O + + zone, T e (O + + ), have been estimated following standard procedures, from all three [OI11] AA5007,4959,4363 A lines observed simultaneously. Temperatures for the low-ionization zone have been found to be significantly different from the T e (O ++ ) for low abundance HII regions (Stasiriska 1980, 1982). We followed the Pagel et al. (1992) method for estimating the low-ionization zone temperatures, based on Stasinska's (1990) photoionization models. Total oxygen abundances were then derived from the ionic abundances by the expression O/H=O + + /H + + O + / H + , using the temperatures obtained as explained above, and emissivities calculated using thefive-levelatom program of De Robertis, Dufour <5i Hunt (1987) with updated atomic data as described in Garnett (1990). We can assume the O 3+ ionization fraction to be negligible, as the Hell A4686 A line has not been detected. We have also adopted a constant temperature throughout the nebula. Assuming that hydrogen and helium lines are produced from recombination, the He+

Terlevich ei al.: Primordial Helium Abundance

186

4x10 - 1 5

1 «

;

•

/ -

•

'

•

•

•

'

•

1

1

•

ISIS blue arm

[OIII]

1

'

^

M

O

' [OIII]

-

2xl0" 1 6 He I

•

u X

10 - 1 6

0-

•

~ 1

4400

i

1

.

4600

.

.

1

1

i

4800 5000 Wavelength (A)

i

,

-

5200

2xlO"16

^-.1.5x10*

E o

2?

10"

5x10"

6000

6200

6400

Wavelength

6600

6800

(A)

FIGURE 3. ISIS blue and red spectra obtained with the WHT, of one of the newly-discovered very low abundance (O/H < 5 % solar) HII galaxies. Notice the temperature sensitive [OIII] A4363 A line and the density sensitive [SII] AA6717, 31 A doublet, as well as the Hel lines AA4471 and 6678 A.

Terlevich et al.: Primordial Helium Abundance

187

abundance can be determined using the strength of the Hel A A 4471,5876,6678 A lines relative to the reference Balmer lines. From the ratio of ( O + / O + + ) , one obtains the correction for neutral He. Knowing Te and n e , if necessary, one can correct for collisional excitation (Clegg 1987; Peimbert & Torres-Peimbert 1987; Clegg, private communication). 3. Uncertainties in the He abundance determination. Constraints on cosmological models require the abundance of He to be determined to better than 5%. Element abundances are not observed, but derived through tortuous means. Uncertainties in the derivation have therefore to be well understood and corrected for or, better still, avoided whenever possible. One of the problems that beset the determination of the He abundance is neutral helium that, if present, remains undetected leading to an underestimation of He; the socalled ionization correction factor (ICF) is very difficult to estimate. The best solution for this problem is to select objects for which ICF is expected to be negligible; excellent candidates are high-excitation giant HII regions. Stellar populations from the galaxy or from the HII region itself may produce underlying absorptions in the H and He nebular lines. This problem can be particularly severe in the Hel A4471 A line and can lead to an underestimate of the value of He. The strength of helium lines might be affected by fluorescent and collisional excitation effects which need to be estimated and corrected for (Clegg 1987; Peimbert & Torres-Peimbert 1987). Stratification of the nebula constitutes a potential source of error because it could lead to wrong electron-density determinations and hence to a miscalculation of the collisional excitation effects, disregarding which leads to an overestimate of the helium abundance. Data obtained with high spatial resolution is hence of paramount importance. The Hel A 5876 A line intensity might be reduced by galactic Nal absorption in A 5890 A in objects with redshifts between 0.002 and 0.004. This was for many years the case for II Zw 40, a typical object used for these studies. Large differences of intensities exist between He and the H reference lines (4471/H7 is typically about 0.1 while 6678/Ha is about 0.01). For this reason, spectra of such high-precision requirements have to be obtained with linear detectors. Also, due to the weakness of the lines involved, and to the need to disentangle emission from underlying absorption, the observations need high spectral resolution. The question of electron temperature fluctuations that shocks, for instance, can produce in the nebulae (e.g. Peimbert 1969), needs to be addressed. If present and ignored, these fluctuations cause abundances to be systematically underestimated. Objects that show clear evidence for shocks - e.g. W-R and supernova signatures - are better not included in the quest for primordial helium (Pagel et al. 1992). Incidentally, both phenomena will also produce local pollution effects. The uncertainty in recombination coefficients for He has been discussed recently (Smits 1991; Pagel et al. 1992); and its effect on He abundance determinations (through discrepancies for the A 5876 and A 6678 A lines of up to 5% for Te ~ 20000 K) was analysed in Skillman k. Kennicutt (1993), where low O/H abundance data from Pagel et al. (1992) and our own preliminary data are plotted using old Brocklehurst (1972) and new Smits (1991) coefficients. Since then Smits (private communication) has discovered an error in his 1991 paper and his new results agree with Brocklehurst's (Smits 1994, in preparation). The whole subject of the recombination coefficients for helium needs revising.

188

Terlevich el al.: Primordial

Helium

Abundance

Problems inherent to the observation and reduction processes have to be considered as well when seeking line ratios to better than 2%. At this level all aspects are suspect; from the standard star calibrations to the reduction procedures - flat fields, response curve fitting, wavelength calibration, etc - we have found them to be potentially able to introduce errors at levels much higher than 2% r.m.s. (Skillman et al. 1994 present a detailed discussion). Is in this context that we have developed a method that involves the combination of repeated observations with different systems (telescope, spectrograph, standard stars, reduction procedures, etc.) and the final result is achieved only when the different methods agree to better than 2% in the final line ratios.

4. Discussion We have obtained new He abundance measurements for very low-metallicity giant HII regions (oxygen abundances in the range 2-5% of the solar value). These are all highexcitation regions with negligible ionization correction factors. Underlying absorption and collisional excitation effects have been carefully corrected for. The He abundance has been calculated from the Hel A6678 A line, measured from high S/N spectra obtained with linear detectors. Figures 4 and 5 represent Y vs. 0/H and vs. N/H (a) from Pagel et al. (1992), to which our new data (6) has been added with darker trace. A maximum likelihood linear fit to these new data (not shown in the figure) combined with that of Pagel et al. yields a value of the intercept of 0.238 with a very small uncertainty of 0.003. Both regressions - with O/H and N/O - give identical values for Yp. Note that the dispersion in He at a given metal abundance is much larger than the uncertainty in the intercept. The dispersion in He/H may be due to systematic effects and therefore it is probably unwise to take the uncertainty in the maximum likelihood fit as equal to the uncertainty in the primordial helium abundance. The net result is a slightly higher value for Yp, bridging the small gap between the theoretically predicted value (Walker el al. 1991) and the best observationally derived one (Pagel el al. 1991). 5. Conclusions We conclude that we are now in the position of determining observationally He abundances to better than 5%, by obtaining high-quality, high-dispersion, high spatial resolution repeated spectrophotometric observations of low metal abundance, high-excitation compact HII galaxies, and avoiding the cases showing clear signs of local contamination by Wolf-Rayet stars and supernovae. We need to use linear detectors and keep in mind that probably the best line to use is Hel A 6678A, because it has a very small dependence on electron temperature and almost none on reddening correction (thanks to its proximity to Ha) and is only slightly affected by underlying absorption and collisional excitation effects. Crucial to deriving primordial helium abundance now, as well as increasing the number of objects at the low-metallicity end of the sample, is the adoption of the best He emissivity coefficients (as discrepancies in their theoretical derivations are larger than observational errors) and the improvement of the models to include the effects of dust, winds and shock heating.

189

Terlevich et al.: Primordial Helium Abundance

-r ] - r • i |

• • • | • • i | i • • |

.28

.26

.24

.22

I I 1 1 1 I 1 | I 1 1 | 1 I 1 1 I 1 I1 1 1 I| 1 I 1 1 I1I .28

Our New Data Added

.26

.24

fcL .22 I i i i I i i i I i i i I i i i I i i i I i i i I i i i I i i i 1

4

6

8

O/HX10

10

12

14

16

5

FIGURE 4. Abundance by mass of Yvs.O/E (top) from the Pagel et al. (1992) sample of low-metallicity giant HII regions and HII galaxies and (bottom) with our new objects added with darker trace. The regression line is from Pagel et al. (1992).

Terlevich et at.: Primordial Helium Abundance

190

i

i

i

|

i

i

i

|

i

i

.28

.26 TTMTH

.24

.22

I I 1 1 1 II1 I1 I1 I 1 I | I I 1 1 1 II 1 I II1 I1 I | .28

Our New Data Added

.26

.24

• t .22 I I I I I II

20

t i l l

40

I I I I I I I I I I I I I I I I I I I I I

60

80

N/Hx10

100

120

140

I

160

7

5. Abundance by mass of Yvs.N/E (top) from Pagel et al. (1992) and (bottom) with our new objects added with darker trace. The regression line is from Pagel et al. (1992).

FIGURE

Terlevich et al.: Primordial Helium Abundance

191

The INT and WHT are operated on the island of La Palma by the RGO at the Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. We wish to thank Lipovetskii and Stepanian for allowing us to select the objects from their unpublished survey, and PATT for awarding observing time. We acknowledge productive discussions with Bernard Pagel, Don Garnett, Max Pettini and Derek Smits, and financial support from a NATO grant CRG 910269 for collaborative research and from the organizers of what turned out to be a very fruitful and enjoyable conference.

REFERENCES ALLER, L. H. 1984 Physics of Thermal Gaseous Nebulae. Reidel, Dordrecht. BROCKLEHURST, M. 1972 Mon. Not. R. Astron. Soc. 157, 211. CAMPBELL, A., TERLEVICH, R. & MELNICK, J. 1986 Mon. Not. R. Astron. Soc. 223, 811. CLEGG, R. E. S. 1987 Mon. Not. R. Astron. Soc. 229, 31p. DE ROBERTIS, M. M., DUFOUR, R. J. & HUNT, R. W. 1987 J. R. Astron. Soc. Canada 81, 195. GARNETT, D. R. 1990 Astrophys. J. 363, 142. IZOTOV, Yu. I., LIPOVETSKY, V. A., GUSEVA, N. G. & KNIAZEV, A. YU. 1992 in The Feedback

of Chemical Evolution on the Stellar Content of Galaxies (ed. D. Alloin & G. Stasinska), p. 134. Observatoire de Paris. KUNTH, D. & SARGENT, W. L. W. 1983 Astrophys. J. 273, 81. LEQUEUX, J.,

PEIMBERT, M.,

RAYO, J. F.,

SERRANO, A. & TORRES-PEIMBERT, S.

1979

Astron. Astrophys. 80, 155. OSTERBROCK, D. E. 1989 Astrophysics of Gaseous Nebulae and Active Galactic Nuclei. University Science Books, California. PAGEL, B. E. J., SIMONSON, E. A., TERLEVICH, R. J. & EDMUNDS, M. G. 1992 Mon.

Not.

R. Astron. Soc. 255, 325. PEIMBERT, M. & COSTERO, R. 1969 Bol. Obs. Ton. y Tac. 5, 3.

M. &; TORRES-PEIMBERT, S. 1974 Astrophys. J. 193, 327. M. & TORRES-PEIMBERT, S. 1976 Astrophys. J. 203, 581. PEIMBERT, M. & TORRES-PEIMBERT, S. 1987 Rev. Mex. Astron. Astrof. 15, 117. SHIELDS, G. A. 1986 Publ. Astron. Soc. Pac. 98, 1072. SKILLMAN, E. D. &; KENNICUTT, R. C. 1993 Astrophys. J. 411, 655. SKILLMAN, E. D., KENNICUTT, R. C. & HODGE, P. W. 1989 Astrophys. J. 347, 875. SKILLMAN, E. D., MELNICK, J., TERLEVICH, R. & MOLES, M. 1988 Astron. Astrophys. 196, 31. PEIMBERT, PEIMBERT,

SKILLMAN, E. D., TERLEVICH, R. J., KENNICUTT, R. C , GARNETT, D. R. & TERLEVICH, E.

1994 Astrophys. J. In press. E. D., TERLEVICH, R. & MELNICK, J. 1989 Mon. Not. R. Astron. Soc. 240, 563. SMITS, D. P. 1991 Mon. Not. R. Astron. Soc. 251, 316. STASINSKA, G. 1980 Astron. Astrophys. 84, 320. STASINSKA, G. 1982 Astron. Astrophys. Suppl. 48, 299. STASINSKA, G. 1990 Astron. Astrophys. Suppl. 83, 501. TALENT, D. L. 1980 PhD Thesis, Rice University. TERLEVICH, R., MELNICK, J., MASEGOSA, J., MOLES, M. & COPETTI, M. V. F. 1991 Astron. Astrophys. Suppl 91, 285. WALKER, T. P., STEIGMAN, G., SCHRAMM, D. N., OLIVE, K. A. & KANG, H.-S. 1991 Astrophys. J. 376, 51. SKILLMAN,

Implications from HI Composition and Lya Emission of HII Galaxies By D. KUNTH 1 , J. LEQUEUX 2 , W. L. W. SARGENT 3 AND F. VIALLEFOND 2 ^nstitut d'Astrophysique, 98 bis bid Arago, 75014 Paris, France 2

DEMIRM, Observatoire de Meudon, 92195 Meudon Principal Cedex, France 3

California Institute of Technology, Pasadena CA 91125, USA

From HST observations we have detected a damped Lya absorption and metal lines that originate from HI gas in front of a star cluster in the HII galaxy I Zw 18. We have found that this neutral gas is 1000 times more metal-deficient than solar and 30 times less than the O/H abundance of the HII region itself. We discuss the implications of these results for the scenario of star formation, the mixing of synthetised elements and the Population III hypothesis. Moreover no Lyman emission has been detected from the ionized gas. Such a surprising result is discussed in the context of the search for young HII galaxies at high redshift.

1. Introduction We know the importance of blue compact dwarf galaxies in the context of galaxy evolution. These objects were first identified as a class by Searle & Sargent (1972) who recognized their unusual spectroscopical nature making them indistinguishable from giant extragalactic HII regions. Also named HII galaxies they are the closest objects we can find to being pure starburst galaxies (Melnick 1987). Their optical and ultraviolet continuum is dominated by the O and B star population. In one HII galaxy a broad emission typical for WR stars has been first detected at about 4650 A by Allen, Wright &; Goss (1976). Kunth & Sargent (1981) later discovered such a feature in the dwarf galaxy NGC 3115, estimated the number of WR stars (found to be large) and concluded that their presence was further strong evidence that star formation in HII galaxies did occur in bursts rather than continuously. Support that this is indeed the case also comes from color evidence (their integrated blue colors indicate that present star formation rates are larger than the past star formation rate), the presence of large amounts of HI gas (showing that these objects are un-evolved) and their underabundance in heavy elements. Their unusual stellar distribution and neutral hydrogen content combined with the very low heavy-element abundances provided the alternative view that some of these HII galaxies could be genuine "young" galaxies forming their first generation of stars. The ongoing burst would be the first one ever initiated in the HI complex. Whether this idea is correct or whether we are simply looking now at one particular burst event among a series of a few is still a matter of debate. Nethertheless starburst phenomena taking place in chemically unevolved gas can be very well studied in such nearby objects. This has been one of the motivations for the numerous surveys that led to the discovery of hundreds of known HII galaxies. The largest compilation with close to 1000 spectra is that of Terlevich et al. (1991), where 400 galaxies are classified with absolute blue stellar continua ranging from -14 to -24. This compilation as well as other surveys were also aimed at searching for extreme low-metallicity systems but the evidence is that no object with abundance below or equal to that of I Zw 18 has been found (see Figure 1). From recent spectroscopic follow-up observations of a list of galaxies drawn from the Second Byurakan Survey Isotov et al. (1992) revealed a few more objects with O/H close to but 192

Kunth et al.: HI Composition and Lya Emission in HII Galaxies 1

,

I

1

1

1

1

1

i

i

'

i

i

,

i

.

193

i

10 -

£

I

i

5

h n 1

7

1

1

1

1

1

1

7.5

1

i 8

, ,

i

i

i

8.5

i

i

• i

i

9

12 + Log 0 / H FIGURE 1. Distribution of oxygen abundance for HII galaxies taken from Terlevich et al. (1991). not lower than in I Zw 18. To explain the lack of extremely metal-deficient HII dwarf galaxies, Sargent and I have proposed that star-forming galaxies and giant HII regions in general are self-contaminated in heavy elements on time scales shorter than the lifetime of the HII regions from which abundances are measured (Kunth k and Sargent 1986). This means that even one burst raises the oxygen abundance to a level measured in I Zw 18.

2. The HI metal content in I Zw 18 I Zw 18 is a very blue dwarf galaxy with the lowest 0 / H ever measured. Two knots are conspicuous, both are sites of vigorous star formation episodes. Kunth & Sargent hence proposed that I Zw 18 could be a young galaxy undergoing its first burst and with observed metals undiluted in the HI halo. We conjectured that most of the HI envelope should be formed of pristine matter. For this reason Kunth et al. (1994) decided to test this hypothesis by observing the main HII region of I Zw 18 with the GHRS on board the Hubble Space Telescope with the aim of detecting heavy elements - if any - from the HI through their absorption lines against the continuum of the ionizing stars. This project was thought to be feasible since we knew that this galaxy contained a large amount of neutral gas (unpublished VLA observations) while the HII core had a very blue stellar complex with virtually no reddening as exemplified from its WE spectrum (Mas-Hesse 1990). Spectra were obtained in 2 spectral regions with a resolution of 20000. The results are displayed in Figure 2. Three outcomes have strong astrophysical bearings: 1. metal lines were indeed detected in absorption from the HI gas of I Zw 18, 2. no Lyman a emission was detected from the I Zw 18 ionized gas, 3. O and Si lines originate from a galatic high velocity cloud (HVC) at -160 km s" 1 . Its O/H abundance is about 0.2 the solar abundance, ruling out the possibilty that fresh pristine gas is infalling onto the Galactic disc. Furthermore, O and Si absorption lines from a Galactic cloud at zero velocity have been detected.

194

K u n t h ei al.\ HI Composition

Line

and Lya Emission

in HII

Galaxies

Column density (atom cm" 2 )

Abundance (by number)

I Zw 18/Galaxy

0 1 1302A

lower limit best guess upper limit

7.5 10 14 2.8 10 15 7.2 1015

2.1 10~ 7 7.9 10" 7 2.1 10" 6

2.7 10~ 4 1.0 10" 3 2.6 10~ 3

Sill 1304A

lower limit best guess upper limit

1.2 1014 2.6 1014 7.6 1014

3.4 10" 8 7.5 10" 8 2.2 10~ 7

8.0 10" 4 1.9 10" 3 5.5 10~ 3

TABLE 1.

We refer to the original Kunth et al. paper for a full discussion of point 3 and will specifically focus on the I Zw 18 measurements: Lyman a absorption: the spectrum around 1216 A shows a strong damped Lyman a absorption line at the redshift of 740 km s" 1 (that of I Zw 18) that is much larger and deeper than that of the foreground Galactic HI. The restricted spectral range of the GHRS spectrum did not allow for a determination of the continuum on the sides of the broad line. We thus used the WE co-added spectra obtained from the WE archives. The line, fitted with a series of Lyman a profiles gave a column density N(E) - (3.5 ± 0.5) 1021 atom cm" 2

(2.1)

quite in agreement with the column density of HI derived from VLA maps over the same region size. Heavy-element abundances: the other spectrum taken around 1300 A shows strong absorptions from OI 1302.2 A and Sill 1304.4 A redshifted at the velocity of I Zw 18. We obtain the following equivalent widths: W{OI) = 390;^ 5 mA

(2.2)

W(SiII) = 280±^ mA

(2.3)

These lines are saturated. Assuming that their Doppler width is equal to that of the HI cloud as measured in our VLA observations, we get a Doppler parameter 6 = 27(±3) km s" 1 . Corrections for saturation, given W and 6, were obtained from Stromgren (1948). In Table 1 we give three values for the column densities: a strict lower limit obtained by using the lower limit for the equivalent width combined with the upper limit to 6, a best guess corresponding to the most probable values, and an upper limit obtained by using the upper limit for the equivalent width combined with the lower limit for b. We adopted O/H = 8 10~4 for the Galactic value (as measured in Orion) and Si/H = 4 10~5 for the solar value. Using the grid of models of Stasiriska (1990) we find that most of the neutral oxygen must be inside the surrounding HI cloud. Moreover from our measurement of the neutral hydrogen column density in front of I Zw 18 and using the slope of the continuum in the WE spectrum we have determined a lower limit for the gas-to-dust ratio in the HI cloud. We find a maximum E(B - V) - 0.013 mag in I Zw 18 leading to a gas-to-dust ratio N(K)/E(B -V)> 2.7 1023 cm"2 mag" 1 , 50 times the Galactic value.

Kunth et al.: HI Composition and Lya Emission in HII Galaxies

195

u a>

x 3

0 -

1200.0

1290.0

1230.0

1210.0 1220.0 Wavelength (angstrom)

1295.0

1300.0

1305.0

1310.0

Wavelength (angstrom) FIGURE 2. GHRS spectra of the dwarf galaxy I Zw 18. Top: Lyman a region. Bottom: 1300A region showing the raw spectrum and a smoothed version. Dotted marks show the I Zw 18 lines.

196

Kunth et al: HI Composition and Lya Emission in HII Galaxies 2.1. The origin of the metals in the cold gas

There is oxygen in the neutral gas of I Zw 18! However the most probable value for O/H is about 1/1000 solar. This finding is in itself most important and opens ways for study of ISM enrichment processes near primordial abundances. The oxygen abundance in this medium is at least 2.1 10~ 7 and the best guess is 7.9 10~ 7 , corresponding respectively to 2.7 10~ 4 and 1.0 10~ 3 of the abundance in the Orion nebula and in M17 (Peimbert et al. 1992). This is considerably smaller than the abundance in the HII region (1.5 10~ 5 ) of I Zw 18 itself. The large discrepancy between the O/H abundances of the ionized and cold gas points to the young age of the starburst and to the limited efficiency of ISM mixing in I Zw 18. These measurements raise several tantalizing questions: 1. were metals in HI produced in a previous burst? 2. alternatively could the matter in the HI be pristine? 3. if so, why is primordial galactic matter with non zero metals?

2.2. The burst hypothesis and the mixing mechanism It can be shown that a single previous burst with a strength equal to the present one could have raised the metallicity at the level observed in the HI: the amount of HII is about 3 106 MQ whereas the total HI mass reaches 7 10 7 MQ, i.e. 23 times larger, and this is about the dilution factor needed. What would be the timescale for such a dilution? The answer needs a review of the ways in which metals get ejected and are mixed through the galaxy; a full account of the discussion is given in Roy k, Kunth (1994). We know that triggered star formation by expanding shells can ensure radial transport at the kiloparsec scale. Since small galaxies have small binding energies it turns out that shells can expand more freely than in massive galaxies (Puche et al. 1992). A parcel of gas in a typical Magellanic irregular would take at least 4 109 years to disperse by turbulent processes (see also Bateman & Larson, eq. 7, 1993). Hovewer this parcel of gas is from time to time involved as well in a star bursting process. Such a time interval between bursts has been quantified by Elmegreen (1992), who calculated the collapse time of swept-up matter along the perimeters of expanding and decelerating shells. A time as short as 20 Myr is required, corresponding to a dormant phase in between bursts. This is clearly too short to account for the appearance of a galaxy such as I Zw 18 and its metallicity enrichment, nor does it appear valid for extreme gas-rich dwarf galaxies and yet very stable against collapse as recently discussed by Lo et al. (1993). If therefore triggered processes are not efficient enough to maintain star formation, quiescent phases allow turbulent diffusion to act in ionized, neutral and molecular clouds. For a characteristic length scale L with a turbulent velocity u the time scale for diffusion is L/u. For neutral and molecular gas with typical u of 2 and 0.5 km s~l respectively we find time scales of 4 10 8 and 1.6 109 yr, much longer than the lifetime of the HII regions alone. In I Zw 18 metals produced in one single burst, more than 109 years ago, can have fully mixed into the HI core. Obviously, the metals produced over a time scale of few 107 yr are merely mixed into the HII region only. We thus conclude that the O/H observed at present could really be due to self-enrichment as suggested by Kunth and Sargent. Note that inhomogeneous chemical evolution models with self-enriched regions have been advocated by Pilyugin (1992) and by Malinie et al. (1993) to explain the G-dwarf problem in our Galaxy. Although I Zw 18 presents the most obvious case for a large discontinuity in metal abundance, fluctuations are observed in our Galaxy as well and in other gas-rich galaxies such as the LMC, indicating that mixing in the ISM is far from being complete (Roy k Kunth 1994).

Kunth et al: HI Composition and Lya Emission in HII Galaxies

197

2.3. Is the HI gas primordial? We formally can answer this question in the negative since metals (O and Si) have been detected, although in minute quantities. The value derived for the O/H abundance is unprecedentedly low for any local gaseous object. For comparison, the lowest oxygen abundances found in Galactic halo stars are probably 3 10~ 4 times lower than in the Sun (Spite & Spite 1991; Barbuy 1988) and are strikingly similar to our derivation. The upper limit we find for the oxygen abundance in I Zw 18 is higher, approaching in abundance that of the HII region. However assuming with Bouchet et al. (1985) that the gas-todust ratio is inversely proportional to the abundance of heavy elements, consideration of extinction gives an independent indication that the heavy-element abundance in the HI cloud of I Zw 18 is indeed very low. On the other hand the only other gaseous objects with heavy element abundances as low as those in the neutral gas in I Zw 18 are the QSO absorbers at high redshifts. In these objects the most secure abundances are those deduced for the "Damped Lyman a" absorbers which have very large HI column densities larger than 1020 atom cm" 2 . Pettini et al. (1992) have derived abundances not much lower than 1/10 that in the sun but with a large spread about the mean. These systems are thought to be primitive galactic disks that have reached different stages of evolution as early as z = 3. For systems with 10 17 < VV(HI) < 10 20 atom cm" 2 that are revealed from their Lyman limit discontinuities, Steidel (1990) found - 3 < [M/H] < - 1 . 5 . The lowest column densities, i.e. with N(HI) < 10 17 atom cm" 2 are found from the HI lines of the "Lyman a forest" (which are thought to be produced in tenuous, highly ionized, intergalactic clouds with masses similar to those of dwarf galaxies). No heavy elements have been detected with certainty (see Lu 1991 for a possible detection giving [C/H] ~ —3.2). In summary, the lowest heavy-element abundances deduced from QSO absorption lines are encountered in primitive galaxy halos and in intergalactic clouds and are comparable to the value that we have deduced for the neutral gas in I Zw 18.

3. The Lyman a emission in I Zw 18 and the search for primordial galaxies Lyman a emission has not been detected in the GHRS spectrum to an upper limit of 10~14 erg sec"1 cm" 2 . The lack of Lyman a emission is a striking result. Indeed, previous observers have reported that young or unevolved galaxies exhibit weak or absent Lyman a emission combined with a strong UV continuum (Meier k Terlevich 1981). However, it was also noted that the Lya/H/? ratio, although much lower than that expected from recombination theory (case B), correlated with the observed O/H abundance of the ionized gas. This trend has been recently reconfirmed by Terlevich et al. (1993) by adding new IUE observations of two extremely metal-poor HII galaxies. From their Figure 3, we expect I Zw 18 to exhibit a Lya/H/? ratio of about 10. Using an aperture comparable to the one we adopted with the HST, French (1980) finds an observed F(H/?) = 2.4 10~14 erg sec"1 cm" 2 . We thus expect F(Lya) = 2.4 10" 13 erg sec"1 cm"2, 24 times more than our quoted upper limit. This result shows that even when the metallicity is low, Lyman a can be suppressed. The destruction of the Lyman a photons cannot be attributed to the foreground Galactic clouds because of the velocity difference between them and I Zw 18. The most natural explanation for both the weakness of Lyman a and the observed correlation is that dust absorption of Lyman a is combined with multiple scattering in neutral hydrogen. But even in the absence of dust absorption, multiple scattering redistributes Lyman a emission over the whole extent of the HI cloud. As we have observed through a very narrow aperture, we might have lost most of the re-emitted

198

Kunth et al.: HI Composition and Lya Emission in HII Galaxies

photons. Many factors can affect the escape of the Lyman a photons, in particular geometrical ones. In the present I Zw 18 case the HI must be very homogeneous over the slit aperture so as to allow an efficient screening to the escape of the Lyman a photons, and most photons, if not destroyed by dust, must escape from other lines of sight (through HI holes or perpendicular to the observed region). In any case the absence of the Lyman a emission cannot be explained as if I Zw 18 was observed nearly after a burst of star formation as suggested by Valls-Gabaud (1993), since, on the contrary, its colors and the shape of its WE continuum suggest a very young burst (Kunth & Mas-Hesse, 1994). Several attempts have been made to search for primordial galaxies at high redshift under the assumption that their Lyman a line would be easy to detect. It is clear from this work alone and from previous WE observations of emission-line galaxies (see Terlevich et al. 1993 for the updated sample) that the result might often be inconclusive. Recent model calculations from Chariot & Fall (1991) come to the conclusion that strong Lyman a emission requires either a low abundance of dust or an AGN. Moreover the results might be much harder to interpret because of the dependence of the Lyman a photons on orientation. Several cases have been reported of strong Lyman a emission at high redshifts but their equivalent widths, their large velocity widths and the presence of other strong lines of highly ionized species suggest that most of the ionizing radiation originates from AGNs (see Chariot & Fall for references herein). Other observers report spatially resolved emission close to QSOs at high-redshift. The most recent case is due to M0ller &; Warren (1993) who found a gas cloud at less than 10 kpc projected distance. Although the authors strongly believe that star formation is responsible for the Lyman a emission we suspect that at such a distance from the QSO the ionizing field is much larger than the backgroung UV light: it is well known that "proximity effects" observed from Lyman a absorption lines studies extend over more than 300 kpc from the QSO. 4. Conclusions • Metals have been detected in the cold HI gas of the nearby galaxy I Zw 18. • For the first time metals have been quantified through one of the most abundant species, i.e. the oxygen that is mostly neutral in the HI cloud. It was found that the O/H is 1/1000 solar. This abundance compares with that of the oldest Galactic halo stars and with high redshift absorption systems seen in QSO spectra. • Metals in the HI could have originated from previous bursts and have mixed in a time scale of about 109 years. The present HII region is self-enriched and must be at least 107 years old. • Alternatively, metals could originate from a population III star generation. We are observing more dwarf HII galaxies with different metallicities by using the GHRS on the HST to clarify this point. • No Lyman a emission originates from the HIII region of I Zw 18, adding more confusion to the the scheme by which the most metal-deficient dust-free galaxies should emit more Lyman a photons. This result could be due to geometrical effects but indicates that in the line of sight the HI medium must be very homogeneous. For the other reported cases (WE data) dust is shown to be the most efficient way to destroy the ionizing photons.

REFERENCES BARBUY, B. 1988 A. Si A. 191,

121.

BATEMAN, N. P. T. & LARSON, R. B. 1993 Ap. J. 407

634.

Kunth et al: HI Composition and Lya Emission in HII Galaxies

199

BOUCHET, P., LEQUEUX, J., MAURICE, E., PREVOT, L. & PREVOT-BURNICHON, M.-L. 1985

A. &: A. 149, 330. CHARLOT, S. & FALL, S. M. 1991 Ap. J. 378, 471.

ELMEGREEN, B. G. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto &; F. Sanchez), p.381. Cambridge Univ. Press. FRENCH, H. B. 1980 Ap. J. 240, 41. HARTMANN,

D. & BURTON, W. B. 1993 In preparation.

ISOTOV, Y. I., LIPOVETSKY, V. A., GUSEVA, N. G., STEPANIAN, J. A., ERASTOVA, L. K. & KNIAZEV, A. Y. 1992 In The feedback of chemical evolution on the stellar content of

galaxies 3rd DAEC Meeting Observatoire de Paris (ed. D. Alloin & G. Stasinska), p. 127. Observatoire de Paris. KUNTH, D. & MAS-HESSE, J. M. 1994 In preparation. KUNTH, D. & SARGENT, W. L. W. 1981 A. & A. 101, L5. KUNTH,

D., LEQUEUX, J.,

SARGENT,

W. L. W. & VIALLEFOND, F. 1994 A. & A. In press.

KUNTH, D. & SARGENT, W. L. W. 1986 Ap. J. 300, 496. LEQUEUX, J. & VIALLEFOND, F. 1980 A.
Lu, L. 1991 Ap. J. 379, 99. MALINIE, G., HARTMANN, D., CLAYTON, D. D. & MATHEWS, G. J. 1993 Ap. J. 413, 633.

MAS-HESSE, J. M. 1990 Doctoral Thesis. Universidad Complutense: Madrid. MEIER, D. & TERLEVICH, R. 1981 Ap. J. Lett. 246, L109.

J. 1987 In Starbursts and Galaxy Evolution, Proceedings of the 22nd Moriond Astrophysics Meeting (ed. Trinh Xuan Thuan, T. Montmerle & J. Tran Thanh Van), p. 215. Editions Frontieres: Gif sur Yvette.

MELNICK,

M0LLER, P. & WARREN, S. J. 1993 A. & A. 270, 43. PEIMBERT,

M.,

TORRES-PEIMBERT,

S. & Ruiz, M. T. 1992 Rev. Mex. Astron. Astrophys. 24,

155. PETTINI, M., HUNSTEAD, R. W., SMITH, L. J. & KING, D. L. 1992 In First Light in the

Universe, Proceedings of the 8th IAP Meeting (ed. B. Rocca, B. Guiderdoni, M. Dennefeld &; J. Tran Thanh Van), p. 97. Editions Frontieres: Gif sur Yvette. PlLYUGlN, L. S. 1992 A. & A. 260, 58. PUCHE, D., WESTPFAHL, D., BRINKS, E. & ROY, J.-R. 1992 A. J. 103, 1841.

ROY, J. R. & KUNTH, D. 1994 In preparation. SEARLE, L. & SARGENT, W. L. W. 1972 Ap. J. 173, 25. SPITE, M. & SPITE, F. 1991 A. & A. 252, 689.

STASINSKA, G. 1990 A. & A. SuppJ. 83, 501.

STEIDEL, C. C. 1990 Ap. J. Suppl. 74, 37. STROMGREN, B. 1948 Ap. J. 108, 242. TERLEVICH, E., DIAZ, A. I., TERLEVICH, R. & GARCIA VARGAS, M. L. 1993 M. N. R. A. S.

260, 3. TERLEVICH, R., MELNICK, J., MASEGOSA, J., MOLES, M. & COPETTI, M. V. F. 1991 A. &

A. Suppl. 91, 285.

The IR and X-Radiation of the Starburst Dwarf Galaxy UGCA 86 ByGOTTHARD M. RICHTER1, M. BRAUN 1 AND R. ASSENDORP1-2 1

Astrophysical Institute Potsdam, An der Sternwarte 16, 14482 Pottsdam, Germany

'Laboratory for Space Research, Groningen, Landleven 12, P. 0 . Box 800, 9700 AV Groningen, The Netherlands UGCA 86 ( = VII Zw 009) is a companion of IC 342 and is one of the nearest starburst galaxies. It contains at least two starbursts, of which the central one is heavily obscured by dust. The IR radiation (IRAS has a relatively steep (cool) spectrum. The X-radiation (ROSAT pointed observation) seems to come from supernovae.

1. Introduction In our optical observations of UGCA 86 (Richter et al. 1991) we found it to be a low surface brightness dwarf galaxy which contains two star formation regions, a central one and one near the southern border. Whereas the southern burst seems relatively normal, the central one is heavily reddened and has also a "softer" appearance (cf. Saha & Hoessel 1991). Together with the finding of an IR source in the IRAS catalog these facts fit very well the assumption that the central burst contains a large amount of dust. Because the catalogued IRAS position did not fit the optical position well, we reprocessed the IRAS observation and observed the galaxy with a ROSAT pointing.

2. Observations We have reprocessed the IRAS data using the GIPSY-IRAS system developed by the Laboratory for Space Research at Groningen (see Wesselius et al. 1992). High-resolution IRAS images were processed in Groningen using a maximum entropy method. While the IRAS Point Source Catalog shows only one source at the area of UGCA 86, the high-resolution images clearly reveal four sources. These four sources are correlated with different starburst regions of the galaxy: with the central region, with two regions in the southern part, and with a region in the tail. The non-detection of UGCA 86 in the 12-/jm band suggests the lack of hot dust, as it might be located (Natta k Panagia 1976) in or near HII regions, whereas the high emission at 60 and 100 /zm confirms the existence of a large quantity of cooler dust. This quite steep spectrum does not continue to radio wavelengths, because we find no source in the 5-GHz Green Bank catalogues (Gregory k Condon 1991; Becker et al. 1991), corresponding to an upper flux limit of 0.02 Jy. On a pointed ROSAT observation using the PSPC, we found an X-ray source with a count rate of (1.1 ± 0.1) • 10~2 ct s" 1 at a2Ooo.o = 3h59m52s, <52ooo.o = 67°8'38". To calculate the source flux we tried tofita combination of a power law and an absorption model to the X-ray count rate spectrum. We have obtained the model parameters: absorption NH = (7.0 ± 2.0) • 1021 cm- 1 , photon index T = +1.0 ± 1.0, and the source flux in the energy range 0.5-2.4 keV: Fx — (4.8±2.4) • 10~13 erg cm" 2 s - 1 . If we assume 200

Richter et al.: The Dwarf Galaxy UGCA 86

201

a distance of 2 Mpc (the estimates of different authors are between 1 and 4.7 Mpc), we get an X-ray luminosity of Lx = (2.3 ± 1.2) • 10 38 erg s" 1 . The derived power-law index of 1.0 ± 1.0 corresponds to the value of 1.0 ± 0.3 as given by Rephaeli et al. (1991) for a combined X-ray spectrum of 53 JRAS-selected starburst galaxies. All the X-ray reduction was done using the EXSAS package. The galactic HI column density from the Bell survey (Stark et al. 1992) is about 3 10 21 cm" 2 . Hence, the remaining 4-10 21 cm" 2 of the total HI column density is inside UGCA 86. This value fits the measurement of Rots (1979), who obtained an amount of 1 • 10 21 cm" 2 for UGCA 86 at an angular resolution of 10'.

3. Discussion We have used a model by Boller et al. (1993) to determine the most probable source of the X-ray emission. Boller et al. have made an estimate of the contributions of different stellar components - supernova remnants, star-driven superwinds, X-ray binaries, and O stars - to the X-ray emission. They have derived the maximum values for the relation of the X-ray luminosity produced by a given stellar component and the star-forming rate, which can be determined using the far-infrared fluxes: M(A = 2 Mpc) = 0.003M© yr" 1 and an X-ray emission of: Lx/M = 2/cdotl048 erg MQ1 . That means that only supernova remnants may explain the relatively high X-ray emission w. r. t. the model of Boiler et al., according to which only about 0.07 supernova remnants should form per solar mass consumed by star formation. Using the derived star formation rate of UGCA 86, we have obtained a supernova rate of one bright object every 5000 years. The X-ray luminosity obtained is quite high for a dwarf galaxy. Only a few very young supernovae, e. g. of Crab type, could reach it; but in this case one would expect higher radio continuum radiation. On the other hand, the IR radiation dominates the electromagnetic spectrum impressively. Meurs et al. (1992) found a new class of spirals which are IR and X-ray bright. Even if UGCA 86 is below the lower boundary of the respective luminosities of these spirals, maybe it marks the continuation of this type at the faint end of the luminosity function. REFERENCES R. L., EDWARD, A. L. 1991 Astrophys. J. SuppL 75, 1. BOLLER, TH., DENNEFELD, M., FINK, H., MEURS, E. J. A. & MOLENDI, S. 1993 Astron. Astrophys. Submitted. GREGORY, P. C. & CONDON, J. J. 1991 Astrophys. J. SuppL 75, 1011. MEURS, E., BOLLER, TH. & DENNEFELD, M. 1992 In Proc. ESO/EIPC Workshop. Elba. NATTA, A. & PANAGIA, N. 1976 Astron. Astrophys. 50, 191. REPHAELI, Y., GRUBER, D., PERSIC, M. &. MACDONALD, D. 1991 Astrophys. J. Lett. 380, L59. RICHTER, G. M., SCHMIDT, K.-H., THANERT, W., STAVREV, K. & PANOV, K. 1991 Astron. Nachr. 312, 309. ROTS, A. H. 1979 Astron. Astrophys. 80, 255. SAHA, A. & HOESSEL, J. G. 1991 Astron. J. 101, 465. BECKER,

R. H.,

WHITE,

STARK, A. A., GAMMIE, C. F., WILSON, R. W., BALLY, J., LINKE, R. A., HEn.ES, C. & HURWITZ, M. 1992 Astrophys. J. SuppL 79, 77. WESSELIUS, P. R., DE JONGE, A. R. W.,

KESTER, D. J. M. & ROELFSEMA, P. R. 1992

Infrared Astronomy with ISO (ed. T. Encrenaz & M. F. Kessler). Nova Science, NY.

In

High-Resolution CCD Surface Photometry of HII Galaxies By EDUARDO TELLES1'2 AND ROBERTO TERLEVICH 2 institute of Astronomy, Madingley Road, Cambridge CB3 OHA, UK 2

Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, UK

1. Introduction HII galaxies are dwarf emission-line galaxies undergoing violent star formation. They are characterized by having giant HII regions which dominate their observable properties at optical wavelengths. Most HII galaxies are contained as a subsample of Blue Compact Galaxies (BCGs), but due to the different selection criteria only a small percentage of BCGs are HII galaxies. We will stick to the name "HII galaxies" to refer to the systems selected by objective prism surveys and having strong emission lines. Various studies of their spectroscopic properties in optical wavelengths have revealed systems of very low heavy-element abundance and high rates of star formation. Earlier morphological studies have suggested that a large proportion of the sample of HII galaxies observed are starlike and isolated (Melnick 1987). For these, no clear indication of the mechanisms which may have triggered the burst is apparent. This together with the spectroscopic properties have made workers in the field since their discovery pose the question of whether these systems may be truly young galaxies or periodic bursts followed by long quiescent periods in the lifetime of the galaxy. Reviews of the general statistical properties of HII galaxies can be found in the Spectrophotometric Catalogue of HII Galaxies (hereafter SCHG, Terlevich et al. 1991; Melnick 1992). We have pursued the line of work of surface photometry to address the questions of morphology and dynamics (Telles & Terlevich 1993; Telles, Melnick & Terlevich 1994) of HII galaxies as well as the questions of the age of the underlying systems and possible importance of the immediate environment in the onset of the present burst of star formation in HII galaxies. In the present paper we present a brief summary of the latter topics which are discussed at length in a forthcoming paper (Telles & Terlevich 1994). In § 2 we briefly describe the origin of the data and method used to derive the colours. In § 3 we discuss our preliminary results.

2. The sample and observations We have used the SCHG as a database for our study of this class of emission line dwarf galaxies (Telles 1994). The observations were made using the Nordic Optical Telescope (NOT) at the Observatorio del Roque de Los Muchachos, La Palma, Canary Islands on 1991 January 13-15 and the 1-m JKT telescope in March 1992 during international time for the GEFEj project. The nights were photometric and had typically sub-arcsecond seeing conditions. Data reduction was performed using standard techniques. We used the t GEFE, Grupo de Estudios de Formation Estelar, is an international collaboration of astronomers from Spain, the U.K., France, Germany, Denmark and Italy, formed to take advantage of the international time granted by the Comite Cientifico Internacional at the Observatories in the Canary Islands for the study of star formation in young stellar systems. 202

Telles k Terlevich: HII Galaxies

203

narrow-band [OIH] frames to discriminate spatially the region where gaseous emission dominates (the burst region) from where line emission is non-existent (the extensions). Thus, the latter is more representative of the underlying system in these galaxies. We have also attempted to account for the contribution of line emission in the integrated broad-band fluxes in the burst region to allow us to derive the integrated colours of the stellar ionizing cluster (see also Salzer, MacAlpine k Boroson 1989a). All methods are described in detail in Telles k Terlevich (1994). We use H o =100 km s" 1 Mpc" 1 .

3. Results and discussion We defer a thorough description of the morphological and structural properties derived from the analysis of the luminosity profiles to the paper in preparation. We restrict the discussion in this contribution to the presentation of what we consider two important points addressed in the cited paper. 3.1. HII galaxy types One important topic in the study of starburst phenomena is the mechanisms which trigger star formation. This section presents morphological evidence that different triggering mechanisms may be responsible for starburst at different scales and suggests that we may be dealing with two different sub-classes of HII galaxies. Figure 1 presents the dereddened total V absolute magnitude vs. total V — I colour diagram of our sample of HII galaxies. The first point in this diagram is the apparent correlation in the sense that brighter galaxies tend to be redder in V — I, although all colours are still very blue for this type of objects. The same is true for the R — I and V — R diagrams. Moreover, one could also claim that the brighter objects do not belong to same colour-absolute magnitude sequence as the objects of luminosity lower than that of II Zw 40 (My « -16.9). The question that then arises is whether there is indeed a real break in physical properties of these dwarf galaxies at that luminosity, and if so what the underlying causes are. It is interesting to note that Binggeli k Cameron (1991) have recently pointed out a break at M B SS —16 in the systematic photometric properties in their sample of ~ 200 early-type dwarfs in the Virgo cluster. They suggested that this break may be caused by a transition in the internal kinematics of the systems from disky or rotationally-supported more luminous systems to spheroidal, anisotropic less luminous systems. Although we do not claim a clear-cut relation between the break at the photometric properties for the dwarfs in Virgo and the apparent break in the photometric properties of HII galaxies presented here it is inevitable that the break occurs at about the same total absolute magnitudes. Moreover, the distribution of H/? luminosity observed in SCHG also appears to be double peaked with a gap at about L(H/?)~ 3 x 1040 erg s" 1 (Terlevich et al. 1991; Melnick 1992). The nature and reality of this gap is not clear, but if we work out the absolute blue magnitude of the continuum (Terlevich 1981), we find this gap to lie between MB values of about —16.5 and —17.5 mag. In order to understand the origin of the "gap" in the luminosity distribution in these low-metallicity dwarf systems, we split our sample in two subsets or types. We will call t y p e I those HII galaxies b r i g h t e r than My w —16.9 and t y p e I I HII galaxies those fainter than My « —16.9. Typically type I HII galaxies have brighter starbursts, have higher metal abundances, smaller W(Rf3) and larger emission-line widths. Furthermore, type I systems tend to show more extreme signs of tidal effects or interactions like extended tail or galactic fans, whereas type II HII galaxies, although they may have more than one burst, tend to be compact and symmetric and do not show clear evidence of

Telles & Terlevich: HII Galaxies

204

1

:'

1

1

1

0.8h 0.6

1

1

1

™ ($aml212+n48

1

1

1

CamOB-Z8A. UM448* * -

_CamO84O+104? (5JM455

iJM167 "

(JJM461

0.4 -

iIam0B40+1201

0.2

1-

~

'_ <^lark36 i:aml409+1200

0 —1

,

-14

,

,

J,IZw40 1 " 1

-16

-18

1

1

1

1

1—

-20

M v

FIGURE 1. Total dereddened colour-absolute magnitude diagram. Circles are NORDIC data and triangles are JKT data. Filled symbols are the type I HII galaxies whereas type II galaxies are shown as open symbols.

external influence. Figure 2 illustrates this general characteristic by showing the contour maps of the objects in the present sample. A more detailed description of individual objects can be found in Telles & Terlevich (1994). One clearly sees that in general type I objects, shown in the upper panel, provide signifcant evidence for extensions and show signs of complex distorted morphologies, whereas the type II objects shown in the lower panel are more compact, have more regular outer isophotes despite of being contoured at the same surface brightness limit and being lower redshift objects; that is, they are intrinsically much smaller galaxies (c.f. physical scales plotted in the figure). Salzer et al. (1989b) in their study of emission-line galaxies, many of which are contained in SCHG, divide their sample in 10 morphological classes. In a histogram of absolute magnitudes for their sample, all compact objects fall at the low luminosity end, whereas at the high luminosity end one finds objects with more complex morphologies such as giant irregulars, interacting pairs, starburst nuclei and seyfert galaxies. It is clear from this evidence that real distinctions may exist between the systems at the two ends of the observed luminosity range. Although we are not able to draw definite conclusion regarding theories of galaxy formation with this small sample of objects, we are tempted to suggest the intriguing possibility that two different mechanisms of triggering star formation may be acting. More massive (luminous) bursts may be formed preferentially from merging or direct collision of two large gas rich systems while low-luminosity bursts may not have been triggered by a strong external interaction, but rather induced by internal processes. 3.2. Their age In order to put constraints on the ages and metallicities of the stellar population in HII galaxies we have compared our broad-band colours with those derived from detailed chemico-spectro-photometric models of population synthesis by Bressan et al. (1994). Figure 3a-c show our colour-colour diagrams for the total colours, for the starburst colours corrected for the presence of emission lines in the filters and for the colours of

205

Telles & Terlevich: HII Galaxies II Zw 40

UM 167

10 Kpc

1 Kpc

0 Kpc

CamO8-28

TYPE I UM 448 .

C.am0840+1200

MRK 36

1 Kpc

10 Kpc

10 Kpc

10 K; ;

UM 455

1 Kpc

Cam 1409+1200

Cam0840+1044

1 Kpc

UM 238

1 Kpc

TYPE II UM 461

1 Kpc

UM 483

1 Kpc

Caml212+1148

1 Kpc

UM 408

1 Kpc

2. Contour maps. Outer isophotes correspond to V ~ 24 mag arcsec 2 . Type I HII galaxies are plotted in the upper panel. Type II HII galaxies are in the lower panel. Physical scale are shown at the bottom of each object. North is top and East is left. FIGURE

Telles k Terlevich: HII Galaxies

206

1.5

(a) 1 1 1 1 1 1 1 11 1 1 | l l l - Circles - NORDIC data Triangles - JK1 data

ii'd/

Solid symbols -- type I gals - Open symbols -- type II gals /

/ -

1 --

1.5

9/ 1

1 1 1

1i

i

7 6\

7

9/

-

UM408 ^ d ' « 4 6 1

i

3

i

-

1

0.5-

i

0 . 5 -~

i4 /

Mrk36 o •C0840+1044

•

M455O QjM23 3 fUM483 Zw40~

7 -

3 840+1201

-

iP1409+l£00

i

0-

11 r 0 0.2 0.4 0.6 V-R Mil

i i i i i i

1-

/IAS

UM13

0-

(b)

1

11

11

i i i

I I I

I

I

I

1 1 l"

0 0.2 0.4 0.6 V-Rburst (d)

1.5

0.5

> 0.5

10' yr burst + gal.

0 I I

0.2 0.4 0.6 V-R"ext

_ I

1.2xlO» yr burst + gal_ I I I I I I I I I I I I I I I I I

0 0.2 0.4 0.6 V-R

FIGURE 3. Colour-colour diagrams, a) Total observed colours. In this panel the dotted lines drawn from each point illustrate the effect of the emission lines on the total colours, b) Colours of the emission line region alone corrected for the emission line contribution, c) Colours of the extensions. The solid lines in a, b and c are the stellar evolutionary model from Bressan et al. (1994) for Z = 0.008. Ticks are labeled in /og(age). d) Composite galaxy. Dashed lines are models of a burst of the two different ages shown superimposed on a 1O10 yr old underlying galaxy for different burst to galaxy mass fractions as labelled.

Telles & Terlevich: HII Galaxies

207

the extensions, respectively. Figure 3d shows two models for a composite systems as described in the figure caption. We have attempted to put constraints on the age of the stellar population both of the stellar system within the starburst and the underlying stellar system by comparing these derived colours with an adopted evolutionary model. We have adopted the model from table 3 of Bressan et al. with metallicity near the observed heavy-element abundance derived from the nebula emission; that is, for Z = 0.008. Once corrected for the emission line contribution, the integrated colours of the stellar population within the starburst (Figure 3b) seem to be redder than one would expect for a very young single burst taking into account the constraints on age and duration of the burst imposed by the spectroscopic observations (Copetti, Pastoriza & Dottori 1986). This may be indicative of the presence of an intermediate-age stellar system underlying the ionizing stellar cluster which would rule out the hypothesis that these systems are experiencing the very first burst of star formation. More interestingly, the colours of the underlying galaxy, if compared with the single stellar population models with standard IMF (Figure 3c), indicate that these systems have ages ranging from 108 to 109 yr. We are unable to detect conclusively a population older than ~ 109 yr in these galaxies. Moreover, if we assume the underlying galaxy to be a composite system of a recent burst superposed on an old stellar system (Figure 3d) then we inevitably conclude that most galaxies must have undergone this recent event of star formation not earlier than 108 yr ago even assuming that 50% of the total stellar mass has been produced in the recent burst. For the extreme cases, a large fraction of the total mass of the galaxy has been involved in a possibly even more recent event, in this sense constituting a genuine young galaxy. We plan to obtain prime-focus INT data to investigate further the indications raised here. The method applied here will be a useful tool specially with deeper observations. We acknowledge support from DGICYT Grant No. PB91-0531(GEFE) and NATO Grant CRG920198 for collaborative research. ET acknowledges his grant from the Brazilian funding agency CNPq.

REFERENCES B. & CAMERON, L. M. 1991 Astion. Astrophys., 252, 27. A., CHIOSI, C. & FAGOTTO, F. 1994 Astrophys. J., submitted COPETTI, M. V. F., PASTORIZA, M. G.& DOTTORI, H. A. 1986 Astion. Astiophys. 156, 111. MELNICK, J. 1987 In Starburst and Galaxy Evolution (ed. T. X. Thuan, T. Montmerle & J. Tran Than Van), p. 215. MELNICK, J. 1992 In Star Formation in Stellar systems (ed. G. Tenorio-Tagle, M. Prieto k. F. Sanchez), p. 253. SALZER, J. J., MACALPINE, G. M., BOROSON, T. A. 1989a Astrophys. J. Suppl. 70, 447. SALZER, J. J., MACALPINE, G. M. & BOROSON, T. A. 1989b Astrophys. J. Suppl. 70, 479. TELLES, E. 1994 PhD thesis, University of Cambridge. TELLES, E. k. TERLEVICH R. 1993 Astrophys. Sp. Sc. 205, 49. TELLES, E. & TERLEVICH, R. 1994 In preparation. TELLES, E., MELNICK, J. & TERLEVICH, R. 1994 In preparation. TERLEVICH, R. 1981 PhD thesis, Univ. of Cambridge. TERLEVICH, R., MELNICK, J., MASEGOSA, J., MOLES, M. & COPETTI, M. V. F. 1991 Astron. Astrophys. Suppl. 91, 285. (SCHG) BINGGELI, BRESSAN,

Formation of Narrow Hell A4686 Emission in HII Galaxies: Link with X-ray Emission By C. MOTCH 1 - 2 , M. W. PAKULL 1 AND W. PIETSCH 2 1

Observatoire Astronomique, UA 1280 CNRS, 11 rue de PUniversite, 67000 Strasbourg, France 2

Max-Planck-Institut fur Extraterrestrische Physik, W-85470, Garching, Germany

ROSAT and optical observations of NGC 4861 and NGC 5408, two HII galaxies exhibiting both narrow Hell A4686 emission line and luminous soft X-ray emission reveal that X-ray and line emitting regions are spatially distinct. This demonstrates that the X-ray ionization mechanism which accounts for the Helll regions around luminous X-ray sources in the LMC is not at work on a large scale in these two HII galaxies. We Briefly review alternative explanations for the formation of the narrow Hell A4686 line. The X-ray spectrum of NGC 4861 is soft and probably not dominated by young accreting neutron stars. Accreting black holes of stellar masses, hot gas and SNRs may account for the integrated X-ray energy distribution.

1. Introduction Although the strength of most emission lines observed in HII galaxies are in excellent agreement with those expected from the ionization by clusters of massive O stars, the presence of narrow, presumably nebular Hell A4684 emission with an intensity offt*1 % of that of H/? in some HII galaxies is not yet well understood (see e.g. Campbell 1988; Conti, 1991). There are strong indications that the gas is photoionized and that shock excitation, if present at all, is not responsible for the emission of this high-ionization line. The origin of the luminous He+ Lyman continuum is currently a matter of debate since normal O stars have generally been thought to emit only negligible amounts of He+ ionizing photons.

2. Possible sources of E > 54 eV photons O stars: Recent "unified" model atmospheres show enhanced extreme UV flux (Kudritzki et al. 1991). This led Gabler et al. (1992) to suggest that massive 0 stars might excite narrow A4686 emission in HII galaxies. However, stringent observational upper limits to the UV flux of O stars are at variance with the predictions of these new models and agree with the "old-fashioned" hydrostatic NLTE models (Pakull, this conference). X-ray sources: Another proposed mechanism is ionization by the very luminous X-ray sources known to exist in some of these HII galaxies (Pakull & Motch 1989). X-ray ionized nebulae albeit of smaller size have been discovered in the Large Magellanic Cloud (Pakull & Angebault 1986). Using the range of soft X-ray temperatures kTth and intrinsic column densities NH consistent with the ROSAT energy distribution of NGC 4861 (see below) we computed for a given emergent X-ray flux the number of He+ Lyman continuum photons present in a "typical" HII galaxy and the corresponding expected narrow Hell A4686 flux. We list in Table 1 the predicted Hell A4686 over ROSAT flux ratio versus kTth and the observed values. Low kTth allow a larger fraction of the X-ray radiation field to be trapped in the giant extragalactic HII region (GEHR) and hidden to detection in the ROSAT band. While high kTth cannot reproduce the observed ratios, radiation fields with slightly cooler temperatures may still be compatible with most observed ratios apart may be for I Zw 18. 208

Motch et al: Formation of Narrow Hell \4686 Emission in HH Galaxies

209

TABLE 1. X-ray ionized nebula model versus observations Model

Observed values

^narrow / ^emergent I 'Hell A4686/- P x(0.1-2.4 keV) |

IcT

2.0 1.0 0.5

0.3

0.0013 0.0023 0.0070 0.0300

rmarrow / ^emergent * H e I I A4686/- f x(0.1-2.4 keV)

I ZW 123 NGC 3125 NGC 5408

0.0044 0.006 - 0.027 0.0093 0.03 - 0.08 0.7- 1.0

NGC 4861 IZW18

NGC 4861

NGC 4861

Thermal bremsstrahlung

Power law

-2 •o

c c -4 o

o a. -6

-8

2 3 4 NH (102" cm" 2 )

2 3 4 NH (102* cm" 2 )

FIGURE 1. Allowed X-ray emission parameters derived from the ROSAT PSPC spectrum of NGC 4861 integrated over the entire galaxy. The three contours correspond to confidence levels of 68%, 95.5% and 99.7%.

Wolf Rayet stars: On the other hand, observations of extended Helll regions have indeed revealed population I objects with luminous He+ Lyman continua: these are certain Wolf-Rayet stages of massive stellar evolution, i.e. early WO stars (Davidson k Kinman 1982) and WN 1 stars (Pakull k Motch 1989; Pakull 1991).

3. ROSAT observations and X-ray spectrum of NGC 4861 NGC 4861 is a powerful soft X-ray source (Lx ~ 1O40 erg s"1; Fabbiano et al. 1982). NGC 5408 was found to be an exceptionally X-ray luminous GEHR for its optical luminosity (Lx ~ 1040 erg s"1; Stewart et al. 1982; Fabian k Ward 1993). We observed these two galaxies with ROSAT using the PSPC (NGC 4861) and the HRI (NGC 5408) detectors. Owing to the statistics, only simple models werefitto the count distribution of NGC 4861. Power law and thermal bremsstrahlung spectra give equally good fits whereas a Raymond Smith optically thin model failed to represent the data for metallicities larger than ss 0.05 solar.

210

Motch et ai: Formation of Narrow Hell \4686 Emission in HII Galaxies 90

NGC 4861 Nucleus 4200 FIGURE

4400

4600 4800 Wavelength (Angstroms)

5000

2. Optical spectrum of the nucleus of NGC 4861 showing the simultaneous presence of narrow and broad Hell A4686 line emission

We show in Figure 1 the allowed parameter space for the thermal bremsstrahlung and power law fits to NGC 4861. The X-ray spectrum is soft with YIth - 0.3 - 1.4 keV or a photon index smaller than -2. Galactic NH is RJ 1 1020 H atom cm" 2 , smaller than that allowed by the two fits suggesting that some absorption takes place in the H n galaxy.

4. Optical observations NGC 4861: A medium-resolution spectrum was obtained at Observatoire de Haute Provence (OHP) with the 1.9-m telescope and the CARELEC instrument. The slit covered NGC 4861 in its elongated direction and was aligned along the axis joining the brightest HII region at the SW edge of the galaxy and the galactic star located to the NE. NGC 4861 is one of the few cases, along with POX 4 (Vacca k Conti 1992), of GEHR exhibiting both narrow and broad Hell A4686 lines in emission as shown on Figure 2. The narrow Hell A4686 component has a FWHM of 3.8 A and is unresolved while the broad component has a FWHM of 35 A. The flux ratio of the narrow to broad component is as 1/3. We only detect narrow and broad Hell A4686 emission in the bright SW HII region. NGC 5408: NGC 5408 was observed at various positions using the red and blue medium resolution spectroscopic modes of EMMI on the ESO NTT. Narrow Hell A4686 line emission is detected from the two main HII regions 3 and 4 (see Figure 5). The bright notches 1 and 2 are stars in our Galaxy. The FWHM of the Hell A4686 line (2.4-4.6 A) is consistent with that of nebular lines and essentially unresolved with our instrumental setting. No broad Wolf-Rayet feature is conspicuous in our spectra. The narrow width of the Hell A4686 line strongly suggests nebular origin in both cases.

5. Spatial distribution of X-ray and optical emission NGC 4861: The ROSAT attitude error (« 8") was corrected for by identifying X-ray sources in the central PSPC field of view with objects on a digitized and astrometrically calibrated POSS plate using the MAMA at the Paris Observatory. Figure 3 shows PSPC

Motch et al: Formation of Narrow Hell \4686 Emission in HII Galaxies NGC 4861 BiJand»*J^PC 0.20-4.50 keV :::

•:::h'i:

211

KJQC 4861 B band/* PSPC 050-1.50 keV

S±&

1'

V

N

FIGURE 3. Spatial distribution of X-ray emission from NGC 4861 in two energy bands overlaid on a CCD B image. The lack of soft X-rays from the bright HII region may reflect intrinsic photoelectric absorption in the galaxy. At higher energies, the X-ray emission is elongated and its maximum is shifted with respect to the bright star forming region.

380

320

280 Spatial Position

240

200

FIGURE 4. Distribution of Hell A4686 narrow (thin line), H/? (dashed line) and X-ray (0.2-1.5 keV) (thick dotted line) emission along the main axis of NGC 4861 joining the nucleus and the galactic foreground star.

0.2-0.5 keV and 0.5-1.5 keV images overlaid on a B CCD image obtained at OHP with the 1.2-m telescope. The 0.5-1.5 keV X-ray emission appears elongated in the direction of the main axis of the galaxy with a maximum flux displaced by sa 17" from the main HII region. There is evidence for a lack of soft photons from the nucleus region (see Figure 3) which probably reflects enhanced X-ray absorption from the denser regions. Figure 4 shows that the maximum X-ray emission is shifted with respect to the main Hell A4686 emitting region which coincides with the bright HII region. Unless photoelectric absorption and/or spectral characteristics of the sources located in the bright HII region hide a luminous population to the view of ROSAT we may conclude that the lack of spatial coincidence between X-ray and Hell A4686 emission in NGC 4861 does not support the X-ray ionized nebula picture.

212

Motch et al.: Formation of Narrow Hell \4686 Emission in HII Galaxies

NGC 5408• V$and + ROSAT HRI'

10 ardsec FIGURE 5. NGC 5408: ROSAT HRI image overlaid on a V frame obtained with the ESO-MPI 2.2-m telescope and EFOSC. Profile analysis shows that the HRI source is likely to be unresolved. The main GEHRs (3,4) cannot be responsible for the X-ray emission.

NGC 5408: Several foreground X-ray active stars are detected in the HRI field of view and allow a direct check of the ROSAT attitude positioning. A bright, probably unresolved X-ray source located in the direction of NGC 5408 is detected with a count rate of 0.051 count s" 1 in the HRI. However, its position is displaced by « 10" from the main HII regions, significantly more than allowed from the combined attitude and internal position errors (see Figure 5). Therefore, we do not confirm the conclusion by Fabian & Ward (1993) that the luminous X-ray source originates from the GEHRs in NGC 5408. We may even consider the possibility that the X-ray source is unrelated to the HII galaxy, i.e. a foreground active star or a background AGN for instance. The only visible object on our V frames located within the HRI error circle i s a K a 16.5, F6-G4 late type-star (Object 1 in Figure 5) which does not exhibit any signature of X-ray activity (Ca H&K emission cores or A6708 Li absorption) and therefore is an unlikely counterpart to the X-ray source. We detect narrow Hell A4686 line emission from mainly HII regions at positions 3 and 4. Detailed inspection of our spectra at position 1 does not reveal any evidence for an enhancement of the A4686 line intensity toward the X-ray source. In fact we do not detect any Hell A4686 emission from the error circle of the HRI source. Our HRI map allows to put an upper limit of about 1/20 of the bright source to the total X-ray emission from the two main HII regions 3 and 4. This implies an F^rjri°A4686/^xTo.?-^ 4 keV) r a t i ° of « 0.2, much too high to be explained by X-ray ionization only. Therefore, whatever is the real nature of the bright X-ray source, we conclude again as for NGC 4861 that X-ray ionization is unlikely to account for the strength of the Hell A4686 line in NGC 5408.

6. D i s c u s s i o n Origin of the narrow Hell A4686 line emission: Our result does not imply that the X-ray ionization mechanism may not be at work in some cases yet to be discovered.

Motch et al.: Formation of Narrow Hell \4686 Emission in HII Galaxies

213

However, some other formation process has to be advocated for these two GEHRs. The presence of both, narrow and broad A4686 emission in several HII galaxies suggests that both phenomena may be intimately linked. We have in fact previously pointed out that a reasonable population of hot WN stars may be a sufficiently luminous source of E > 54 eV photons and have a broad stellar line which could remain undetectable with present instrumentation (Pakull & Motch 1989). Origin of the X-ray emission: Young accreting neutron stars exhibit power-law X-ray spectra with photon index in the range of 0 to —2 (White et al. 1983). This suggests that X-ray emission in NGC 4861 is not dominated by this particular kind of X-ray binary. Similar conclusions may apply to the source toward NGC 5408 which with kTth = 0.44 - 0.63 keV (Fabian & Ward 1993) seems incompatible with the expected radiation of a large population of young accreting neutron stars. Possible X-ray emitters are accreting black holes of stellar masses (kTth « 1 — 2keV), hot gas and SNRs. A collection of « 100 accreting black holes could account for the total soft X-ray output of NGC 4861. With a starburst age of « 5 106 yr, shorter than the time required for the formation of the first active massive accreting binaries (tevo\ ^ 9 106 yr; van den Heuvel, 1983), we may not yet be witnessing the X-ray binary active phase of these galaxies if no burst of star formation took place during the last few 107 yr. The ROSAT project is supported by the Bundesministerium fur Forschung und Technologie (BMFT). C.M. acknowledges support from a CNRS-MPG cooperation contract and thanks Prof. J. Triimper and the ROSAT group for their hospitality. The MAMA (Machine Automatique a Mesurer pour l'Astronomie), is developed and operated by CNRS/INSU. This work is partly based on observations collected at the European Southern Observatory, La Silla (Chile) with the 2.2.m telescope of the Max-Planck-Society, with the ESO-NTT telescope and on observations obtained with the 1.2-m and 1.9-m telescopes of the Observatoire de Haute-Provence, CNRS, France.

REFERENCES CAMPBELL, A. 1988 Ap. J. 335,

644.

Conti, P. S. 1991 Ap. J. 377, 115. DAVIDSON, K. & KINMAN, T. D. 1982 P. A. S. P. 94,

634.

FABBIANO, G., FEIGELSON, E. & ZAMORANI, G. 1982 Ap. J. 256, FABIAN, A. C. & WARD, M. J., 1993 M. N. R. A. S. 263,

397.

L51.

Gabler, R., Kudritzki, R. P. & Mendez, R. H., 1992 A. & A. 265, 656. KUDRITZKI, R. P. ET AL. 1991 In (eds.) Massive stars in Starbursts (ed. C. Leiterer et al.), p. 59. Cambridge Univ. Press. PAKULL, M. W. & ANGEBAULT, L. P. 1986 Nature 322, 511. PAKULL, M. W. & MOTCH, C. 1989 Extranuclear Activity in Galaxies (ed. E. Meurs &; B. Fosbury), p. 285. PAKULL, M. W. 1991 In Wolf-Rayet Stars and Interrelations with Other Massive Stars in Galaxies (ed. K. A. van der Hucht fc B. Hidayat), p. 391. Stewart, G. C , Fabian, A. C , Terlevich, R. J. k. Hazard, C.) 1982 M. N. R. A. S. 200, 61p. VACCA, W. D. & CONTI, P. S 1992 Ap. J. 401,

543.

VAN DEN HEUVEL, E. P. J. 1983 In Accretion driven stellar X-ray sources (ed. W. H. G. Lewin & E. P. J. van den Heuvel), p. 303. WHITE, N. E., SWANK, J. H. & HOLT, S. S. 1983 Ap. J. 270,

711.

Environmental Effects in Star-Forming Dwarf Galaxies By J. M. VILCHEZ Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain I present here some preliminary results of a study of the spectroscopic properties of a group of star-forming dwarf galaxies that have been selected in order to sample the range in physical conditions imposed by extreme density environments. This investigation is part of an ongoing project intended to evaluate the relative influences of the environment, and the initial conditions, on the evolution of galaxies with active star formation. It has been found that, on average, starforming dwarf galaxies located in nearby low density regions appear to present spectra with higher excitation, higher H/J equivalent widths and larger total H/? luminosities than similar objects located within high-density environments.

1. Introduction The influence of the environment on the mechanisms that control star formation is one of the most important subjects concerning the study of the origin and evolution of galaxies. Spiral galaxies in clusters, a high-density environment, have been used to trace the present-day star formation rate (SFR) in order to compare with field galaxies. Some studies of spirals in clusters suggest a reduced SFR with respect to field galaxies of the same morphological type, while others tend to favour a similar or higher SFR in cluster spirals; this question still remains open (Moss & Whittle, 1993). Current environmental effects which can be operating in galaxies as a consequence of the interaction with companions and with the intergalactic medium include, among others, tidal shaking, tidal stripping, ram pressure sweeping and evaporation. Interactions with big companions could be catastrophic for low-mass systems such as star-forming dwarf galaxies. These galaxies present the lowest mass surface density and rotation velocities measured for objects supporting active star formation (Gallagher & Hunter, 1989). They are therefore fragile systems easily affected by environmental factors. Thus, they can be ideal objects to test the influence of the environment on the star formation process. The variations observed in the number density of dwarf galaxies could be a direct consequence of their strong interaction with the environment. The observed large-scale distribution of galaxies shows that dwarf elliptical galaxies, dE, are associated with dense regions with a high pressure intergalactic medium (e.g. Binggeli et al. 1989; Thuan et al. 1990). On the other hand, dwarf galaxies located in low-density regions defined by Voids in the current distribution of galaxies, would remain very faint or undetected (cf Babul & Rees, 1992). It has been suggested that dwarf galaxies may be more numerous than it is currently assumed, and some studies have already concluded that dwarf HII galaxies are much less clustered than normal galaxies (Iovino et al. 1989). Dekel & Silk (1986) on the other hand, proposed that these galaxies, being absent within dense cluster cores, are in fact filling the Voids. On the other hand, Peimbert & Torres-Peimbert (1992) explored the possibility that emission line galaxies may correspond to a younger population with respect to lower SFR galaxies outside the Voids, and performed a spectroscopic study of a sample of 214

Vilchez: Environmental Effects in Star-Forming Dwarf Galaxies

215


1

3 2

. . .

1

. . .

1

1 . 1

Virgo Cor*

I . . .

I

Voldi

[OIII] \5007 A/H/3

FIGURE

1. Histograms of the number distribution of dwarf galaxies in different environments as a function of their observed excitation.

these galaxies in the Bootes Void. They conclude that some Bootes galaxies tend to present anomalous chemical abundances with respect to foreground nearby counterparts, and their excitation conditions appear also to be different. In the same direction, the spectrophotometric study of Gallagher & Hunter (1989) on Virgo cluster dwarf galaxies conclude that some real differences may be apparent between their derived SFR and the values observed in field galaxies. Within an ongoing project intended to evaluate the influence of the environment (and initial conditions) on the evolution of star-forming galaxies, we present here some preliminary results of a study of the spectroscopic properties of a sample of star-forming dwarfs that have been selected in order to map the conditions imposed by extreme density environments: LOW DENSITY REGIONS, including galaxies in three Voids and in the foreground of the Virgo Cluster periphery, and HIGHER DENSITY REGIONS, which include a subset of dwarf galaxies located close to the core of the Virgo cluster plus an isolated small condensation of galaxies refered to as the "Clump" in the following. 2. The sample Most galaxies have been selected from the sample of emission line galaxies of the University of Michigan (UM), survey lists IV and V (MacAlpine & Lewis, 1978; MacAlpine & Williams, 1981). In this sample, galaxies with absolute magnitudes of —16 and fainter are detected at redshifts in excess of 5000 km s~1, and contain a large fraction of low-

216

Vilchez: Environmental Effects in Star-Forming Dwarf Galaxies

, , | , , . . | . , 7

Clump

7

—

—

. . 1 . . , , 1 . ." u

a | z

•

• '

6

7

a 7

'

'

'

i

1

.

.

1

Virgo Core

Voids

- - J -

•

•

1 •

2

n

Log W(H0) (A)

FIGURE 2. Histograms of the number distribution of dwarf galaxies in different environments as a function of their observed H/? equivalent width.

luminosity galaxies compared with other surveys. From Salzer (1989) we have selected three major Voids apparent in the distribution of nearby star-forming galaxies. Within these regions, our search for catalogued emission line galaxies has produced a subset of 19 objects from the UM lists. Positional and velocity information was confirmed with the help of the Spectroscopic Catalogue of HII Galaxies (Terlevich et al. 1991). Within this subset, several galaxies clasified as Seyferts, plus a nuclear starburst, were identified from the catalogues, and subsequently eliminated of our study. The remaining objects, therefore, are expected to be true star-forming dwarf galaxies: 13 objects located deep inside the Voids plus other three which appear to be delineating the back wall of Void # 3 . In addition, a condensation of 11 emission line galaxies in a clump (Salzer 1989) have also been selected. From these, three of the objects that are catalogued as nuclear starbursts were subsequently eliminated from the final group. Candidates from regions of low to intermediate density were obtained selecting a group of foreground galaxies located in a zone of the Periphery of the Virgo cluster, the Local Supercluster region, some 13° south of the cluster core. From a total of 20 UM galaxies selected in that area, two of them classified as nuclear starbursts were discarded. On the other hand, the group of objects representative of high-density environments includes all the star-forming dwarf galaxies located close to the Virgo cluster core, having available spectroscopic data in the literature. Positions of the galaxies have been taken from the Virgo Cluster Catalogue VCC (Binggeli, Sandage & Tammann 1985).

Vilchez: Environmental Effects in Star-Forming Dwarf Galaxies

217

44

42 i

in

ao CD

~

.40

38

-

36 -10

,

-15

,

i

i

I

,

i

-20

FIGURE 3. The luminosity of H/3 as a function of the blue absolute magnitude Ms of the sample. Symbols: black circles, Virgo Core; black triangles, Clump; open squares, Virgo Periphery; open circles, Voids.

3. Preliminary Results Spectroscopic data from Gallagher k Hunter (1989) and Izotov &; Guseva (1989) have been used for the objects belonging to the VIRGO CORE subset. For the rest of galaxies belonging to the UM survey, blue magnitudes were taken from Salzer et al. (1989, pi); H/3 luminosities and spectroscopic ratios were derived from the Spectroscopic Catalogue of HII Galaxies of Terlevich et al. (1991). For a few objects with no data available in the catalogue, these ratios were computed from data in Salzer et al. (1989, pll). Figure 1 shows histograms of the number distribution of dwarf galaxies as a function of their observed excitation, measured by the ratio [0III] A5007 A/H/?, for each of the regions studied : VOIDS and VIRGO PERIPHERY (Low Density, right in the figure), CLUMP and VIRGO CORE (High density, left in the figure). The curve corresponding to the VIRGO PERIPHERY has been reproduced as a dashed line along with the VIRGO CORE one to illustrate the difference. Higher-excitation objects appear to be associated to lower density regions. The H/3 equivalent width is a measure of the strength and age of the starbursts. Larger values of W(H./3) are systematically found in recent and strong bursts, like those observed in HII galaxies. Figure 2 shows the histograms of the number distribution of star-forming dwarf galaxies as a function of their observed H/? equivalent width, W(K0) in A, for each of the regions studied : VOIDS and VIRGO PERIPHERY (Low Density, right in the figure), CLUMP and VIRGO CORE (high density, left in the figure). The curve corresponding to the VIRGO PERIPHERY has been highlighted as a dashed line, along with the VIRGO CORE one to illustrate the difference. It appears that larger values of W(Hf3) are associated to objects located in the lower-density regions. The luminosity of H/? is a measure of the global present-day SFR, and overall, it should

218

Vilchez: Environmental Effects in Star-Forming Dwarf Galaxies

.8

•

A

O O

o a

a

0

ro • .2

-.2

i

I

i

i

I

i

2.5

1.5 Log W(H/3)

FIGURE

4. [OIII]A5007 A/H/J as a function of the equivalent width of H/?, W(H/3). Symbols as in figure 3.

correlate with total luminosity. Figure 3 shows the luminosity of H/? as a function of the blue absolute magnitude MB for the programme galaxies. The straight line represents the fit to the whole UM survey. Black symbols, corresponding to objects located in highdensity environments, seem to be associated to lower values of the H/? luminosity, and SFR, for a given blue absolute magnitude. Therefore larger SFRs seem to be favoured in objects from low density regions. Finally, in Figure 4 we show the behaviour of the excitation, measured by [OIII] A5007 A/H/? as a function of the equivalent width of H/?, W(H0). Symbols as in Figure 3. It is clear from the figure that those galaxies showing high excitation appear to present also larger equivalent widths. This work has been partially financed by grants DGICYT Nos. PB91-0531 (GEFE), PB90-0570, and IAC P14/86.

REFERENCES BABUL, A. & REES, M. J. 1992 M. N. R. A. S. 255,

346.

B., SANDAGE, A. & TAMMANN, G. A. 1985 Astron. J. 90, 1681. BINGGELI, B., TARENGHI, M. & SANDAGE, A. 1989 Astron. k Astrophys. 228,42. DEKEL, A. & SILK, J. 1986 Astrophys. J. 303, 39.

BINGGELI,

GALLAGHER, J. S. & HUNTER, D. A. 1989 Astron. J. 98, 806. IOVINO, M. ET AL. 1989 In Large Scale Structure and Motions in the Universe, p. 371.

Yu. I. & GUSEVA, N. G. 1989, AstroMka. 30, 564. MACALPINE, G. M., LEWIS, D. W. & SMITH, S. B. 1977 Astrophys. J. Suppl. 35, 203.

IZOTOV,

Vilchez: Environmental Effects in Star-Forming Dwarf Galaxies

219

G. M. &; WILLIAMS, G. A. 1981 Astiophys. J. Suppl. 45, 113. Moss, C. & WHITTLE, M. 1993 Astiophys. J. Lett. 407, L17. PEIMBERT, M. & TORRES-PEIMBERT, S. 1992 Astron. & Astiophys. 253, 349. SALZER, J. J. 1989 Astiophys. J. 347, 152. SALZER, J., MACALPINE, G. M., BOROSON, T. A. 1989 Astiophys. J. Suppl. 70, 447 (pi). SALZER, J., MACALPINE, G. M. & BOROSON, T. A. 1989 Astiophys. J. Suppl. 70, 479 (pll). TERLEVICH, R., MELNICK, J., MASEGOSA, J., MOLES, M. & COPETTI, V. F. 1991 Astion. & Astiophys. Suppl. 91, 285. THUAN, T. X., ALIMI, J.-M., GOTT, J. R. & SCHNEIDER, S. E. 1990 Astiophys. J. 370, 25. MACALPINE,

Theory of Starburst and Ultraluminous Galaxies ByBRUCE G. ELMEGREEN IBM Research Division, T. J. Watson Research Center, P. O. Box 218, Yorktown Heights, NY 10598, USA Starbursts in four galaxy locations are discussed: on the periphery and in tidally ejected debris, in the main disk, in inner Lindblad resonance rings, and in the nucleus. Starbursts in dwarfs are also briefly mentioned. Possible reasons for the starbursts are summarized, mostly in the context of two theoretical models, one where star formation is initiated spontaneously by gravitational instabilities in disks, spiral arms or rings, and another where star formation is stimulated by high-pressure star clusters. The observed rates, efficiencies, and durations of star formation in all five regions follow from the models. We emphasize the importance of a critical density for star formation, which is approximately K2/G for epicyclic frequency K, and the importance of large-scale radial gas flows. Star formation tends to occur wherever the density exceeds the critical value. The rate of star formation is very large in inner rings and nuclear regions because the critical density is very high there. Normal galaxy disks have lower rates because of their lower K. This difference in rates implies that inner rings and nuclear regions of galaxies can maintain their star formation for much shorter times than the main disks following an episode of gas accretion that makes the density exceed the critical value. Thus only the inner regions will have major fluctuations in the star formation rate. Normal galaxy disks probably have fluctuations too, but with lower amplitudes and longer durations. Normal disks may maintain their critical gas densities, giving them semi-continuous mini-bursts, because of the continuous and slow accretion of disk gas driven by torques from spiral arms that are stimulated by remote or dwarf companions. Galaxies without companions have very low surface brightnesses because the gas is forced to remain subcritical. The duty cycle for a starburst is estimated from the theory. Starbursts in the ILR rings of non-interacting barred galaxies can repeat with a 10-50% duty cycle if gas from stellar evolution in the bar continuously drains to the center. The duty cycle for mild nuclear starbursts in non-barred non-interacting galaxies can be even higher. Variations in the efficiency of star formation in virialized clouds as a function of pressure and velocity dispersion are also estimated from the theory. High pressures and high velocity dispersions should both lead to greater efficiencies of star formation inside each virialized cloud.

1. Introduction Bursts of unusually intense star formation occur in various parts of galaxy disks following interactions with other galaxies. The first observations of such enhanced star formation were by Larson & Tinsley (1978), who based their conclusions on the integrated colors of galaxies with peculiar or tidal morphologies. Hummel (1980), Balzano & Weedman (1981), Weedman et al. (1981), Feldman et al. (1982), and Condon et al. (1982) soon realized that this enhancement usually occurs in the nuclei of galaxies, even though the interactions appear to take place along the periphery. The nearby interacting galaxy M82 offered a good example (Rieke et al. 1980), and many more were found in an extensive survey by Balzano (1983) of nuclear emission lines in galaxies. Soon other studies of individual galaxies, such as NGC 3690 (Arp 299) by Gehrz, Sramek & Weedman (1983), were able to fit star formation models to many aspects of the observed nuclear emission. The connection between starbursts and Seyfert activity or quasars was also made at this early stage (Weedman 1983). In the decade since these first studies of starburst galaxies, numerous observations of 220

Elmegreen: Starburst and Ultraluminous Galaxies

221

various types have been made to try to understand the conditions and processes that occur in the bursting regions. With these observations, we can begin to address some of the theoretical issues, such as how the burst begins, why it is located where it is, why the star formation rate is so large, why the initial mass function and star formation efficiency might differ from those in normal galaxies, and so on. This paper discusses these issues in the context of the gravitational instability model of star formation. In this model, star formation is proposed to begin spontaneously in a region of a galaxy after the gas collapses into virialized clouds because of energy dissipation and self-gravitational forces. The additional role of stimulated star formation, in which one generation of young stars compresses the surrounding interstellar gas and triggers another generation to form, is also discussed in the context of starburst galaxies. These two types of star formation are distinct and well-defined mechanisms, different in nature from other proposed mechanisms for star formation in starburst galaxies, such as cloud-collision-induced star formation (Vazquez k Scalo 1989; Olson k Kwan 1990; Habe k Ohta 1992; Jog k Solomon 1992) and the compression of pre-existing clouds in a high pressure environment (Jog k Das 1992). The primary motivation for applying the gravitational instability model of star formation to starburst galaxies is that it appears to be an important mechanism on a large scale in normal galaxies (see reviews in Elmegreen 1987, 1990a, 1993; Larson 1988, 1992; Balbus 1990; Franco 1991, 1993). This conclusion is based on the following observations: the largest cloud complexes in the Milky Way and nearby galaxies all have the theoretically predicted Jeans mass of several times 107 MQ (Viallefond, Goss k Allen 1982; Elmegreen k Elmegreen 1983), and this mass extrapolates in a self-consistent way to larger values (108 MQ) with increasing interstellar velocity dispersion in interacting galaxies (Elmegreen, Kaufman k Thomasson 1993; hereafter EKT93). The largest regions of star formation are separated along spiral arms by the Jeans length (Elmegreen k Elmegreen 1983). The sizes of the star complexes that form in these large clouds are approximately equal to one-half the Jeans length in the outer parts of the disks for all 221 galaxies in the Sandage k Bedke (1988) atlas of galaxies, when the disk instability parameter Q is set equal to 1 (Elmegreen et al. 1994). The star formation properties of galactic disks are sensitive to the threshold density that follows from the instability model (Quirk 1972; Guiderdoni 1987; Kennicutt 1989). Furthermore, gravitational instabilities form clouds faster than random cloud collisions (Elmegreen 1990b), and they also give the observed mass spectrum for the largest clouds (Elmegreen 1991) as well as the density dependence for the star formation rate (Elmegreen 1991). The gravitational instability is also triggered by spiral density wave compressions in the amount required to give the observed mild increase in the star formation rate per unit gas mass in spiral arms, and in a manner that gives the observed spacing between the star formation regions along the arms (Elmegreen 1979, 1994a; Balbus k Cowie 1985; Balbus 1988). Such instabilities are also likely to account for most of the star formation in galaxies without spiral waves (Elmegreen 1991; Elmegreen k Thomasson 1993). Detailed models for the galactic radial distributions of metallicity and star formation rates that are driven by gravitational instabilities are in Wang k Silk (1994). The current gravitational instability models are not without problems, however. The usual derivation of the instability, leading to the Q parameter or critical column density, assumes an initially uniform density and velocity dispersion over the length scales of interest. This assumption contradicts the common observation of a supersonically turbulent interstellar medium, in which the density and dispersion vary by large amounts over a wide range of scales. A more realistic theory including such variations has been addressed by Bonazzola et al. (1987) and Pudritz (1990).

222

Elmegreen: Starburst and UUraluminous Galaxies

The model also has an insufficient development of the equation of state, i.e. the relations between pressure, density, and their time derivatives. The equation of state is an essential component of all gas dynamics. Some reasonable equations of state include instabilities of their own, such as the macroscopic thermal instability, which help the general gravitational instability in this model. This problem was addressed by Cowie (1981), Struck-Marcell k Scalo (1984), Elmegreen (1989b), and others in the context of a cloudy fluid. A third problem with the instability model is the difficulty of observing it directly, as it happens. Other triggers of star formation leave more obvious clues, such as the large moving shells with peripheral star formation that appear in regions where high pressures from one generation of stars have triggered a second generation. There the whole history of the region can be mapped out in space and time. In the case of the instability, however, the only clues are the characteristic sizes and masses of the clouds that form, and perhaps their regular separation in ILR rings or along spiral arms. Other cloud formation processes could conceivably produce star formation on these scales too, so the available clues are not unambiguous. Moreover, the ISM usually displays a wide range of phenomena with no preferred scales - only power-law distributions. Because of this, the instability model can only address the largest scale, or the biggest clouds, which are somewhat better defined than just the scale of a "typical" star-forming cloud, because there is no typical star-forming cloud in a power-law distribution of sizes. The model then has to rely on a turbulent cascade or other process that is not directly related to the instability in order to give all of the smaller clouds in which stars form. Thus the instability model is vastly incomplete in terms of explaining all star formation. A fourth problem is the unknown importance of purely kinematic cloud formation mechanisms, such as those driven by compression and cooling in the presence of turbulent motions and stellar pressures. Purely stochastic models for such cloud formation were reviewed by Seiden k Schulman (1990). More deterministic models are in Chiang k Prendergast (1985), Chiang k Bregman (1988), and Passot et al. (1994). Usually these kinematic models have smaller characteristic scales (100 pc) than the gravitational instability models (1000 pc), in which case they could contribute to the internal structure of the large clouds that form first by the instability, or to the structure of the ISM between the large clouds. They could also dominate the instability, however, producing all of the observed structure even up to the largest scales. Kinematic models have not been developed enough to know their limitations. Regardless of these and other uncertainties about interstellar gas dynamics and the role of gravitational instabilities, there still appears to be ample evidence that such instabilities do have a role in the star formation process in normal galaxies, particularly on the largest scales. Theoreticians then face the challenge of defining and limiting the extent of that role. This paper suggests that, with a reasonable extrapolation of parameters, many aspects of the starburst phenomenon can also be explained by common gravitational instabilities. No doubt there are many other aspects too that cannot be so simply explained. One implication of the presumed extrapolation of normal galaxies to starburst galaxies is that the results support the prevalent idea that starbursts occur with a normal star formation process but with abnormally high gas densities, particularly in their central regions (e.g. Larson 1987; Jackson et al. 1991). Interacting galaxies such as starbursts might also have higher turbulent gas velocities. These two changes, high densities and velocity dispersions, appear to lead to essentially all of the peculiar properties of starbursts, including the large star formation rates and short gas consumption time scales, the small and dense nature of typical star-forming clouds, the prevalence of high-density

Elmegreen: Starburst and UHraluminous Galaxies

223

and high-temperature molecular transitions, and the proposed high-mass bias in the initial mass function. In this model, the key to the starburst phase is not any peculiarity in the star formation process itself, which may be normal, but the high gas density in either an inner Lindblad resonance ring for a barred galaxy or in the nuclear region for a non-barred galaxy. In what follows, Section 2 reviews the four main regions of a galaxy where tidal interactions have an effect on star formation, and it briefly discusses starbursts in dwarf galaxies, Section 3 gives the length and mass scales for star formation in each region, and Section 4 suggests why the starburst characteristics appear.

2. The location of the starburst Bursts of intense star formation occur in essentially four locations in an interacting galaxy: (1) on the periphery, in tidally ejected clouds and debris and in giant cloud condensations in tidal arms; (2) in the main disk in abnormally strong but galaxy-bound spiral arms; (3) in a ring at the inner Lindblad resonance (ILR) of a barred galaxy that contains such a resonance, and (4) in the nucleus when there is no ILR. The most intense star formation occurs in an ILR ring or nucleus because of the high gas densities there. 2.1. Peripheral star formation triggered by galaxy interactions Tidal arms are long-lived material features that readily break up by gravitational instabilities into Jeans-mass clouds containing 107-108 MQ of gas (Barnes & Hernquist 1992; EKT93). If the interacting galaxies have about equal masses, then many of these tidal arm clouds can be ejected into intergalactic space (EKT93). A tidal arm can also contain, in the form of a giant (108 — 109 MQ) pool or cloud of gas, most of the gas mass that was in the outer part of the original galaxy disk (EKT93), producing a giant burst of young stars at the tip of the tidal arm, as observed for the antennae and superantennae galaxies by van der Hulst (1979) and Mirabel et al. (1991; 1992), and possibly for NGC 520 by Standford (1990). Star formation begins in a tidal arm cloud after the internal energy dissipates, since these clouds are gravitationally self-bound from birth in the instability model. Their average densities are low (1-10 cm"3) in the outer galactic regions, so the delay before star formation begins, and the duration of the star formation phase, can be very long, perhaps lO(Gp)~1/2 ~ 108 yr or more (e.g. Sekiguchi & Wolstencroft 1993). This delay is comparable to the companion interaction time, so peripheral starbursts can begin after the time of strongest tidal forces, when the companion is already moving away, and they can last long enough to maintain star formation even after the clouds are ejected. Tidal arms stretch and wrap up because of their material (non-wave) nature. This stretching adds internal kinetic energy to the clouds that condense by instabilities. The gas in a tidal arm is also likely to be agitated directly by the interaction, so again the clouds that form should have higher than normal velocity dispersions. This higher dispersion for a given mass in a virialized cloud implies that the self-gravitational binding energy per gram is higher than normal, and that the cloud is harder to break apart by internal star formation at the normal efficiency (EKT93). Thus star formation in peripheral high-dispersion clouds should continue until a high fraction of the gas is converted into massive, destructive stars. If low-mass stars are also present and the star-to-gas ratio becomes high, then the resulting star cluster will be gravitationally bound, resembling a dwarf galaxy because of its high mass and peripheral location (Zwicky 1959; Zwicky & Humason 1960, 1961).

224

Elmegreen: Starbursi and UUraluminous Galaxies 2.2. Disk star formation triggered by an interaction

Intense star formation in the main disk of a galaxy can also be triggered by an interaction without necessarily leading to any tidal ejection of the gas (Larson & Tinsley 1978; Kennicutt et al. 1987). M51 is an example of a galaxy in which intense star formation (Vogel et al. 1993) takes place in unusually strong spiral arms (Elmegreen & Elmegreen 1984; Rix & Rieke 1993) that were probably triggered by the interaction (e.g. Howard & Byrd 1990; Hernquist 1990). M51 also has an unusually high gas surface density (Morris & Rickard 1982), and consequently a high pressure and molecular fraction in the disk (Elmegreen 1989a). The high surface density alone can account for the high star formation rate per unit area. The surface density is probably high because of mild gas accretion driven by strong spiral arms since the interaction. The gas need only accrete by 30% in radius to double the surface density. Such accretion corresponds to an inward drift speed of ~ 20 km s" 1 in the spiral arms. This inward drift speed, when divided by the sine of the spiral arm pitch angle, is comparable to the observed arm streaming motion of ~ 100 km s" 1 (Rand & Kulkarni 1990), so the inward drift could be present. All galaxies should respond to mild interactions by increasing their global star formation rates, i.e. in the main disks and not just in the nuclear regions. Nuclear starbursts may dominate disk starbursts only for strong interactions because the migration of gas to the center can be very fast for a strong interaction. But only a small fraction of galaxies are currently experiencing a strong interaction. Most have repetitive, weak interactions with numerous distant neighbors and dwarf companions. Each of these weak interactions should force a small amount of inward drift, and with the corresponding increase in the mass column density in the disk, trigger a mild disk starburst. Perhaps the most likely scenario is that all normal galaxies maintain a disk column density greater than the critical value, KC/(ITG) (cf. Section 4), because of continuous mild interactions and a continuous forced inward drift of gas that is driven by tidally activated spiral arm torques and turbulent viscosity. The star formation rate depends on how much the column density exceeds the threshold, so that star formation rates are highest during or after the interaction, and then they gradually decrease as the excess gas (i.e. above critical) is converted into stars. In this scenario, all galaxies should have variable star formation rates, with fluctuations of perhaps a factor of 10, corresponding to the presence or absence of a current or recent interaction. The Milky Way galaxy may be a good example of such variable star formation. Scalo (1987a,b) reviewed the evidence for this, noting in particular how the mini-starburst ages in the Milky Way that were derived by Barry (1988) using chromospheric ages of nearby stars are the same as the ages derived from features in the present-day mass function assuming a constant initial mass function. A more recent review is in Majewski (1993). Evidence for mini-starbursts also comes from stellar kinematics (Gomez et al. 1990), age and metallicity (Marsakov et al. 1990), the white dwarf luminosity function (Noh & Scalo 1990), the age distribution from photometry combined with theoretical isochrones (Meusinger 1990, 1991), and from stellar groupings based on line strength (Roman 1952; Yoss 1992) by Majewski (1993). Noh & Scalo (1990) also pointed out that the starburst history in the Large Magellanic Cloud is about the same as in our Galaxy, as if the mutual interaction between these two galaxies triggered both bursts. This continuous-burst model also explains why galaxies with a very low density of companions also have very low surface brightnesses and star formation rates. Presumably these galaxies never had an encounter that boosted the surface density above threshold (Zaritsky k Lorrimer 1993; Bothun et al. 1993; van der Hulst et al. 1993).

Elmegreen: Starburst and Ultraluminous Galaxies

225

2.3. Inner Lindblad resonance ring star formation The most intense star formation in starburst barred and interacting galaxies occurs in the inner regions (e.g. van den Broek 1993; Wynn Williams k Becklin 1993), either in an inner Lindblad resonance (ILR) ring when the galaxy has a bar, or in the nucleus when the galaxy does not have a bar (Telesco, Dressel, k Wolstencroft 1993). The star formation itself tends to occur in intense hotspots (Morgan 1958; Sersic k Pastoriza 1965, 1967; Pastoriza 1975; see recent observations of the hotspots in NGC 1068 by Bruhweiler et al. 1991 and in NGC 253 by Forbes et al. 1993). A large number of observers have studied such ILR ring starbursts (Telesco k Gatley 1981, 1984; Wolstencroft et al. 1984; Combes k Gerin 1985; Hawarden et al. 1986; Buta 1986a,b; Hummel, van der Hulst k Keel 1987; Arsenault, et al. 1988; Arsenault 1989; Feinstein et al. 1990; Tilanus et al. 1991; Wilson et al. 1991; Tacconi et al. 1991; GarciaBarreto et al. 1991a,b; Sofue 1991; Gerin, Casoli k Combes 1991; Wiklind k Henkel 1992; Combes et al. 1992; Hummel k Jorsater 1992; Sofue k Irwin 1992; Benedict et al. 1992, 1993; Telesco, Dressel k Wolstencroft 1993; Buta k Crocker 1993; Koribalski, Dickey k Mebold 1993; Boer k Schulz 1993). Reviews of nuclear rings and starburst galaxies are in Telesco (1988), Combes (1991), and Henkel, Baan, k Mauersberger (1991). A good example appears to be M82 (Telesco et al. 1991; Waller et al. 1992) The ILR ring of gas forms because gas outside the ILR experiences a decelerating torque from the slower rotating bar, and so falls inward along the bar shocks or dust lanes (Matsuda k Nelson 1977; Noguchi 1988; Barnes k Hernquist 1991). When the gas reaches the ILR, it no longer feels any net torque and tends to stay there in a ring. The gas might also gain energy from the resonant interaction with the bar and resist further accretion this way. Gas inside a second ILR, or inner inner Lindblad resonance, feels an accelerating torque from leading spiral arms that moves it outward to the ring (Combes 1988). Thus an ILR ring can accrete matter from either one side or both sides. For a review of this process, see Combes (1991). The same accretion from corotation to the ILR presumably works for non-barred galaxies too, but because the m=2 potential from the spiral usually becomes weak at or outside the ILR (e.g. there may be a wave-barrier outside the ILR or the wave may get absorbed at the ILR), the resonance interaction at the ILR may not be strong enough to maintain a dense ring there. Then the gas can accrete all of the way to the nucleus or to a wide (~ 500 pc) starburst region around the nucleus. An interaction can initiate the accumulation of gas into an ILR ring by first causing the gas to drain into the bar region through (non-corotating) tidal arm torques outside the bar. Recall that trailing spiral arm torques lead to accretion only inside corotation, so if the spiral region is to accrete to the bar region, and the bar ends near its own corotation, then the spiral has to have a slower pattern speed, i.e. a larger corotation radius, than the bar. Alternatively, the bar could be made by these torques (Noguchi 1987). If the outer disk gas can get into the bar region, or if the bar is newly formed, then the gas in the bar should quickly move inward and accumulate near the ILR. The time scale for this bar accretion is rapid if the bar potential is strong because the streaming motion will be fast compared to the orbit speed. Then each bar crossing produces a strong shock that dissipates a large fraction of the orbital energy. Radial streaming motions along the bar shocks can exceed 100 k m s " 1 (Matsuda k Nelson 1977; Athanassoula 1992). A large amount of gas can accumulate at the ILR, making an extremely large density and star formation rate there. For example, the gas in a bar region that extends for 3 kpc radius can accrete to a ring with 400 pc radius and 200 parsecs thickness (e.g. Hummel

226

Elmegreen: Starburst and Ultraluminous Galaxies

et al. 1987), increasing the gas density by the geometric factor of 75. Then the average gas density in the ring can be several hundred atom cm"3, and the ring mass 108 MQ, which is 10% of the gas mass in the galaxy. This ring density exceeds the critical value for self-gravity to exceed the shear and Coriolis forces, which is ~ K2/G SS 100 cm" 3 at this radius (for epicyclic frequency K), and then star formation can begin with a great burst (cf. Section 4). The mass in the ring turns into stars 10 times faster than the gas in the rest of the disk because of the factor of ~ 100 higher density, scaling approximately as p1!2 (Section 4). This makes the total star formation rate per unit volume 1000 times larger than in the disk (a p3^2) and places about half of the total galactic star formation in the ring. A more extensive discussion of starbursts that are initiated by the gravitational collapse of ILR rings is in Wada & Habe (1992) and Elmegreen (1994b). Relatively isolated barred galaxies (e.g. in loose groups) can have ILR hotspots and starbursts too (Sersic & Pastoriza 1965, 1967), with less massive and lower density rings corresponding to slower accretion rates. In some cases, the gas could have begun its migration into the bar region during an interaction some 2 x 109 yr ago, and then slowly accreted to the ILR ring. In other cases, the ILR could have obtained its gas from slow extragalactic infall to the bar region or from stellar mass loss in the bar region. In either case, the accretion increases the density in the ILR ring and the burst begins when the critical density is exceeded. ILR bursts could be repetitive in isolated galaxies, following long quiescent intervals when the ILR accretes residual gas from the bar region, and then bursting for a relatively short time when the critical density is reached. The duty cycle will be the ratio of the burst duration to the accretion time, and can be on the order of 10-50% (cf. Section 4). The short-term result of ILR ring starbursts is the formation of a blue ring of stars at the ILR. Such rings have been observed by Buta & Crocker (1993) and others. The long-term result could be ring dispersal by resonant scattering of stars. The scattered stars could fill a lens shape in the disk if the resonant forces are confined to the plane, or they could contribute to the bulge if the stars also acquire a perpendicular component of velocity (Sofue & Habe 1992). 2.4. Nuclear star formation If there is no ILR to accrete material from larger radii, then the gas will drain into the nucleus. The time scale and mass involved in this radial motion depend on the spiral arm strength. If the arms are strong and go all of the way from the periphery of the galaxy to the central regions, as might be the case for global tidal arms following a close, equalmass interaction or coalescence, then the nuclear accretion can be rapid and include a large fraction of the total disk gas mass (e.g. > 109 MQ). If the arms do not reach the central region but the waves reflect off a bulge or hot inner disk, then the spiral torques can only bring the gas into a large (several kpc radius) ring surrounding the bulge, where the rotation curve turns over (e.g. Dahlem, Bomans k. Will 1994). Further accretion requires a second step of angular momentum loss with perhaps separate spirals or viscous accretion in the nuclear disk. Dynamical models of interacting galaxies with intense star formation are in Mihos, Bothun & Richstone (1992, 1993). The strongest nuclear starbursts follow galaxy coalescence, when two galaxies merge and most of their gas falls to the common center (e.g. Noguchi 1988; Barnes & Hernquist 1991). These bursts are extremely luminous because there is a lot of gas involved, perhaps 1010 MQ. Evidence for large-scale shocks between the colliding interstellar media of two galaxies has been reported by Sargent & Scoville (1991) and van der Werf et al. (1993). In the first or these two studies, in Arp 299, three active star-forming knots are seen in the dense ridge of the shock. Sometimes the merger remnant can be an elliptical galaxy

Elmegreen: Starburst and Ultraluminous Galaxies

227

(e.g. Kormendy & Sanders 1992) or quasar (see references below) or both (Terlevich & Boyle 1993). The density in the nuclear disk can be very high because the geometric compression factor for the accretion can exceed 1000. The density will also be high because star formation has to wait until an extremely high critical density, K2/G, is exceeded. This critical density is high in the nuclear region because of the large tidal and Coriolis forces, making the epicyclic frequency 100 times larger than in the disk. When star formation finally does begin, it goes very rapidly, again because of the high density (see Section 4). Densities exceeding 104 cm" 3 are not unreasonable for nuclear star formation (e.g. Eckart et al. 1991; Mauersberger et al. 1991), and star formation rates per unit volume exceeding 106 times the average for a disk galaxy would follow from these high densities (scaling as p1 5, see Section 4). The burst would then last for a time that is 0.01 of the normal gas consumption time, or some 107 yr, as observed by Bernlohr (1993) and others.

2.5. Starbursts in dwarf irregular galaxies Dwarf irregular galaxies are small galaxies with relatively large blue patches of star formation. Their history of star formation is strongly dominated by bursts (Searle, Sargent & Bagnuolo 1973; Gerola, Seiden & Schulman 1980), Reviews of the observations are in Gallagher & Hunter (1984), Hunter & Gallagher (1986), and Hunter (1992). Perhaps the most successful theory of star formation in dwarf irregular galaxies utilizes stochastic self-propagating star formation (Gerola, Seiden & Schulman 1980). According to this theory, the burst-like nature of star formation in these galaxies is primarily the result of their smallness. The galaxy has to wait a long time for star formation to begin spontaneously, and when it does, the propagating mode quickly fills the disk with star formation. The associated heating then turns off further star formation until the gas cools and another spontaneous event starts the burst again. Smaller galaxies are more burst-like because a single event is more likely to propagation throughout the disk when the galaxy is small, and this allows the whole disk to get hot at the same time, squelching all further star formation until it cools. Small disks also have fewer possible sites for spontaneous events, and therefore longer wait times between bursts. A larger disk has only relatively small patches of propagating star formation, and it usually has reservoirs of cool gas somewhere to keep the star formation going in a globally steady fashion. The gravitational instability model can also explain the burst-like nature of star formation in small galaxies because the instability length scale is a large fraction of the galaxy size. When the gas is cool and has settled to a thin disk with a density exceeding the threshold of approximately K2/G for epicyclic frequency K, the ISM will collect itself spontaneously into one or two large clouds with masses comparable to the Jeans mass. These clouds form stars at a typical rate per unit mass (e.g. Wilson 1992), but because each is such a large fraction of the galaxy size, and because the escape speed from the galaxy is relatively small compared to the expansion speed of HII regions, supernovae, and wind-swept shells, the ISM heating and destruction that results from star formation in each cloud can be relatively severe. This heating puffs up the disk and lowers the density below critical, which turns off star formation globally. The process should resume when the disk cools again. This is the same control process that has been applied to larger galaxies (Goldreich & Lynden-Bell 1965), but it should be burst-like in small galaxies.

228

Elmegreen: Starburst and Ultraluminous Galaxies

3. Length and mass scales for star-forming clouds and young clusters The length and mass scales for gravitational instabilities depend primarily on the balance between the self-gravitational force density and the pressure gradient. Additional forces such as Coriolis and magnetic contribute in the linear regime only after the collapse motion starts. The Coriolis force, for example, depends directly on the perturbed velocity, and the perturbed magnetic force depends on magnetic field gradients which appear only after the gas moves and drags the field along with it. Thus the characteristic length and mass scales are essentially those that give neutral equilibrium before the motions begin. The rate of collapse and the critical density at which collapse occurs depend on rotation and other forces, but the mass does not. For a two-dimensional distribution of gas, the characteristic or fastest-growing wavelength comes from the wavenumber k = 2TT/A at which the gravitational and pressure terms in the dispersion relation have their peak difference; i.e. at the peak of k2c2 — 2irGark for effective velocity dispersion c and reduced-mass column density
(3.1) where the latter approximation is for typical ar =
M2D =

*£_ „ i£l2 „ 2.1 xiO*M0 f'Yf^—^) 1

G2a2

Ga

Vkms- /

~\

\ M Q pc~2/

(3.2)

For a long and thin distribution of gas with average density p , as in a spiral arm or piece of a ring, the self-gravitational acceleration per unit length is approximately AirGp(l — kArKi[kAr]) for Bessel function K\ and arm or ring half-thickness Ar. If the wavelength is larger than the arm or ring thickness, then this expression is approximately 2Gfik2\n(2/kAr) for mass per unit length \i = pirAr2. Thus the peak of 2 2 2 2G\ik ln(2//feAr) — k c determines the characteristic length and gives \lD = vAre0*1*'"'™ M1D = T2pAr3e°^1+c2^G^

~ 8.5Ar

(3.3)

~ 26.8?Ar 3

(3.4)

where the latter approximations are for C2/G/J, ~ 1. This dimensionless parameter c2/Gn is analogous to the ratio of the velocity dispersion to the virial velocity for a cloud. For a strongly self-gravitating spiral arm or ring, the parameter should be about 1 or larger, which corresponds to fi > 230(c/km s" 1 ) 2 M© p c " 1 . These results imply that peripheral star formation should make clouds in tidal spiral arms with masses of ~ 108 M Q (n/10 cm~ 3 )(Ar/225 pc) 3 and separations of ~ 8.5Ar ~ 2 kpc. The clouds in ILR rings can have about the same mass but a smaller dimension Ar (and separation) by a factor of 3 and a larger density by a factor of 30. The masses of clouds in nuclear disks can be much smaller, in inverse proportion to the surface density, unless the velocity dispersion is high. For a nuclear disk with 109 MQ inside a 100-pc radius, a ~ 3 x 10 4 M o pc~ 2 and M 2 D = 6 x 104(c/10 km s" 1 ) 4 MQ. The separation between these clouds would be only 3(c/10 k m s " 1 ) 2 pc.

Elmegreen: Starburst and Ultrahminous Galaxies

229

4. Time scales for star formation: why a burst results In the gravitational instability model of star formation, the star formation rate per unit volume on a large scale is given by the expression SFR sss euip

(4.5)

where e is the local efficiency of star formation, or fraction of the gas mass that turns into stars per cloud, w is the growth rate of the instability on the large scale, which is the cloud formation rate, and p is the average density of gas in all forms. The star formation rate is controlled on a small scale by the efficiency per cloud, e, which presumably results from cloud destruction by the stars it forms. Propagating star formation may be included in this e term by writing it as eo/(l — p) for an average number p of other star formation sites with a similar size that come from the first spontaneously generated site; eo is for a single site. If p fa 1, then this equation breaks down and other methods to get the SFR must be used, such as that developed by Seiden & Gerola (1982). The SFR is presumably controlled on a large scale by the instability growth rate, w, which depends primarily on the average density as (Gp)1'2. If cr is the mass per unit area in the disk, then the SFR per unit area is eua for the same e and w. We can write this u in terms of cr instead of p using the relation cr = ipH for scale height H, which now involves the velocity dispersion also. Then u ~ TTGCT/C. A different star formation law, suggested by Wyse (1986) and Wyse & Silk (1989), puts the star formation rate proportional to the molecular cloud column density, and the latter oc cr(HI)2(fi —fio).This law has been successfully applied by these authors to the radial distributions of molecular clouds and metallicities in galaxies. Wyse introduced this law because the star formation rate in galactic disks was observed to be roughly proportional to the molecular cloud density, and she conjectured that this molecular density depends on the square of the atomic hydrogen column density, cr(HI), because of molecular cloud formation by cloud collisions. The dependence on angular velocity, fi, comes from the rate of passage through spiral density waves, or any other radially-dependent kinematics; fio is a constant evaluated at the edge of the disk, or at corotation near the disk edge. More recent studies suggest that the star formation rate per unit area correlates better with the total H2+HI density than with either the H 2 density alone or the square of the HI density (Buat et al. 1989; Kennicutt 1989; Tenjes & Haud 1991). Galactic CO observations also suggest now that the CO abundance does not depend on the presence of a density wave (Stark et al. 1987). Nevertheless, density-squared empirical laws like this have successfully explained many of the bulk properties of galaxies. Equation 4.5 gives a similar result when w « (irGcr/c)(l — Q2)1?2 from the usual dispersion relation for radial instabilities in a non-magnetic isothermal disk, using Q = K.c/(irGa). Then the SFR per unit area is oc TtGcr2 /c times (1 — Q2)1^2, which has a a2 dependence multiplied by a function analogous to fi —fiothat varies slowly with radius and goes sharply to zero in the outer part of the disk. A recent study of star formation using a dynamically-based law like this is in Wang & Silk (1994). Another star formation law was used by Mihos et al. (1992, 1993) for N-body simulations of interacting galaxies with gas clouds and stars. They conjectured that the SFR in each cloud was oc Mp for cloud mass M and environmental density p. They find that this law has an effective column density scaling like a18. There is no column density threshold in this law, but they conjecture that at high velocity dispersions, when cr may be below threshold, the star formation process is dominated by cloud collisions. All of these laws are semi-empirical, following the original Schmidt (1959) law that has found many applications to normal galaxies. This is because there is no complete theory of star formation at the present time, only an idea of what processes are likely

230

Elmegreen: Starburst and Ultraluminous Galaxies

to occur. In this case, the empirical laws have a lot of value for reproducing common observations, but the problem with them is that they cannot be reliably extrapolated to regions where they have not yet been observed directly. Equation (4.5) is another such semi-empirical law, but it can be made arbitrarily detailed and realistic by including additional terms in the equations that determine ui, such as magnetic fields, turbulence, rotation, shear, viscosity, and so on, and by considering realistic heating and cooling laws for the equation of state. Perhaps the most elaborate system of equations that has been used to determine w is in Struck-Marcell & Scalo (1984) and Scalo & Struck-Marcell (1984), who consider cloud-collisional formation and destruction as well as the usual forces that drive cloud motions. A less detailed treatment in Elmegreen (1991) combines the gravitational, Parker, and thermal instabilities with rotation, shear and magnetic fields. The results of these theories are not simple expressions for u, which would be desirable, because the theories are too complicated. Instead the primary results so far are a better understanding of the relative importance of various physical processes. Yet the basic star formation law that comes from equation 4.5 using the simplest and most likely approximations has many of the features that are commonly observed in galaxies, such as the threshold density and the power-law scaling with density above the threshold. The threshold density comes from u, which gets very close to zero when the density is much less than K2/(2TTG). The scaling of the SFR with density comes from the scaling of w with p in the unstable regions, which is as p° for strong magnetic fields giving the Parker instability, or as p05 for weaker fields giving an essentially pure gravitational instability, or as p1 if the dominant cloud formation process is driven by cloud collisions or cloud-collisional cooling. These three mechanisms give a SFR oc pa for a = 1, 1.5 and 2, respectively, with intermediate a for mixtures of the instability. A typical mixture of the Parker and gravitational instabilities gives a SFR oc p13 (Elmegreen 1991), which is comparable to that observed by Kennicutt (1989). The scaling of the critical density with epicyclic frequency K, combined with the scaling of the SFR with density, help to explain many aspects of the starburst phenomenon in galactic rings and nuclei (see also Larson 1992). This leads one to believe that the starburst has essentially three phases: (1) the accretion phase, when gas from a bar or galactic disk accretes to the center, or possibly a ring around the center if there is an ILR; (2) the turn-on phase when the density reaches a high critical value proportional to K?/G for epicyclic frequency K, and (3) the burst phase, when star formation proceeds in a high-density environment at an extremely high rate. The stars form quickly once they begin forming - in a burst - because of the high density, and the threshold density is high because K is high at small radii in a galaxy. For a SFR oc p15 and p oc K2/G OC 1/r2, we get a SFR oc 1/r3, which can be 1000 times larger per unit volume at or inside 500 pc in the nuclear region than in the main disk. Most of this excess SFR occurs only because the density is higher (one power of p) but some of it is because the rate of conversion of gas into stars is higher also (w oc p1'2)The rate per unit area is ewer, which scales with a2 if the scale height is C/(2TTG/>) 1 ' 2 for constant velocity dispersion c. This scaling is because p =
Elmegreen: Starburst and Ultraluminous Galaxies

231

a consumption time of around 108 yr (as observed by van den Broek 1992) once the gas reaches the critical density. The star formation rate per unit mass is larger by this same factor of 10. Thus 50% of the galactic star formation can occur in 10% of the total gas mass if it is in an ILR ring. The SFR and consumption times are even more extreme at smaller radii (see review in Scalo 1987b). If there is no ILR to stall the accretion of gas to the center, and the gas reaches a very small disk around the nucleus, e.g. with a 50-pc radius, then the critical density for star formation is higher than in the main disk by a factor of nearly 104 (again scaling with K2), and the consumption time is smaller by a factor of 100, now approaching 107 yr. The SFR per unit volume will be higher by a factor of 10 6 , and per unit area by a factor of 104 if c is constant. The gas mass in the nuclear disk will be 100 times smaller than in the total galaxy if c is constant, because the critical column density scales with K OC 1/r and so the mass scales with r. If the velocity dispersion increases near the center so H is about the same as in the main disk, then the nuclear disk at critical density can hold as much mass as the entire disk, namely about 109 MQ or more. In that case an entire galaxy of gas mass can convert into stars in only 0.01 times the usual time, or 10 7 yr, making a luminosity 100 times larger in the nuclear region than in a whole galaxy. Such luminosities are around 10 12 LQ and the galaxy would be called ultraluminous. A good example of an ultraluminous galaxy with a powerful starburst and possible QSO core is IRAS F10214+4724 (Rowan-Robinson et al. 1991, 1993; Brown & Vanden Bout 1991; Soifer et al. 1991; Solomon et al. 1992; Lawrence et al. 1993; Clemens et al. 1993). Because of its large redshift, this object may be a galaxy in the process of formation. Other young starburst galaxies were discussed by Windhorst et al. (1991) and Eales et al. (1993). Ultraluminous infrared galaxies were reviewed by Mirabel (1991). Radio observations that support the starburst model for the ultraluminous energy source were reported by Condon et al. (1991) and Sopp & Alexander (1991). Evidence for AGNs in ultraluminous galaxies with single nuclei was discussed by Majewski et al. (1993). A recent molecular line study of ultraluminous galaxies was made by Downes et al. (1993). The final result of such intense nuclear star formation would be an extremely dense cluster of stars in the inner several tens of parsecs (Koornneef 1993). The nuclear cluster would presumably form or supplement a massive nuclear black hole because of star collisions and gas accretion (e.g. Weedman 1983; Barnes & Hernquist 1991). Because of further accretion on the black hole, the galaxy could then have an active nucleus, with rapid variability, a relativistic jet, etc. (see review in Osterbrock 1993). Quasars could have formed this way (Sanders et al. 1988; Stockton & Ridgway 1991; Mazzarella et al. 1991; Lipari & Macchetto 1992; Hamann & Ferland 1992). Many ultraluminous and starburst galaxies appear now to be having their second burst of nuclear star formation with a black hole already present, and they show characteristics of both the starburst and the active nucleus (Colina & Perez Olea 1992; Lipari et al. 1993). Usually the AGN part of the system shows a small nuclear radio source (Norris et al. 1990; Lonsdale, et al. 1992, 1993), although some of the standard AGN properties have been attributed entirely to the starburst phenomenon too (Jarvis & Melnick 1991; Terlevich et al. 1992; Cid Fernandes et al. 1992; Cram, North & Savage 1992). Some interacting galaxies show only starburst activity (Forbes et al. 1992; Smith & Kassim 1993). The short gas consumption times in the ILR ring and nuclear regions, scaling approximately as p~rit oc r, imply that the intense star formation rate is not likely to be continuous even if the accretion rate is continuous. Instead, the nuclear regions should be dormant for a while after a burst because the gas density is less than K2/G. Then

232

Elmegreen: Starburst and Ultraluminous Galaxies

accretion gradually increases the density above this critical value and star formation begins at a very high rate. When the gas is exhausted after a short time (oc r), star formation stops and the central regions are dormant again. Because of the requirement that a critical density be exceeded, and because of the rapid rate of star formation after it is, the star formation process should be burst-like at a small enough radius. We can estimate the duty cycle of the star formation, or the fraction of the time in the burst phase, as follows. Suppose the accretion rate of mass to the central region is a steady M. Then the mass in the central radius r, which is airr2, increases as Mt for time t since the last burst. When a equals the critical value for Q = 1, which is KC/(TTG), the burst begins. Thus the time between bursts is t — Kcr2/(MG). The duration of the burst is the gas consumption time given above, or l/(ew), where u> is the instability growth rate, which is u> = irGcr/c, but at the critical density, this is w = K. Thus the ratio of the duration of the burst to the time between bursts at the critical density is duration of burst MG MG . A „. 2 2 2 (4.6) time between bursts n cr e 4V ce where we have used K = 2V/r for solid-body rotation speed V in the central regions. This result equals , „ / M \ ( V \~ / c \~l / e \~1 duration of burst = 1.0 1 1 time between bursts ' I 10 M 0 yr~M \100 km s " / V l O k m s - / \0.1> (4.7) For this reasonable guess at parameters, the duty cycle is about 50%, i.e. the burst is on about half the time. If the velocity dispersion in the burst region is larger, and the efficiency is larger, as we have discussed elsewhere in this paper, then the duty cycle can be less, perhaps only a few per cent. This duty cycle should also be the fraction of barred galaxies with ring starbursts (because we have used V — 100 km s" 1 , which is typical for the ring region). Note that the duty cycle fraction depends on the radius of the gas concentration, r, through the circular velocity, V, which scales linearly with r in the inner regions of a galaxy. Because V is likely to be smaller for tiny nuclear bursts than for ring bursts, perhaps by a factor of ~ 10, the fraction of galaxies (without ILRs) that have nuclear bursts should be much larger than the fraction of galaxies (with ILRs) that have ring bursts. This prediction applies to non-interacting galaxies that have steady accretions of ambient gas to the nuclear regions. For interacting galaxies, the fractions with a burst depend more strongly on the fractions that are interacting, because a sudden increase in M from an interaction will lead to a burst more rapidly than the t we have just calculated. But note that binary galaxies tend to be early-type and barred (Elmegreen, Elmegreen k. Bellin 1990), which means that they are likely to have ILR rings, so the fraction of barred early-type galaxies with ILR starbursts can be high. Observations of a previous burst (10 9 yr ago) in a galaxy that also has a present-day burst were discussed by Davidge (1992).

5. Changes in the efficiency and IMF with increasing velocity dispersion and pressure The velocity dispersion in the interstellar gas presumably results from an equilibrium between turbulent or cloud-collisional cooling and heating. During a galaxy interaction, or during the accretion phase when the nuclear gas mass increases, the heating rate for the gas should increase enormously. The interaction leads to strong spiral arms and a

Elmegreen: Starburst and Ultraluminous Galaxies

233

general excess sloshing of gas in all three directions (see simulations in EKT93), and the ring accretion should contribute a large compressional energy PdV for specific volume V. Interacting ring galaxies should also have high gas turbulence because of the strong radial and perpendicular motions. The ILR region, where the ring forms in an early-type barred galaxy, should also be a region of continuous strong heating because of the resonant interaction between individual cloud epicycle motions and the bar motion. Such heating is well known for stars at a Lindblad resonance and may contribute to a high gas dispersion as well. Whatever the source of excess turbulent energy in the ring gas, the result of it is that the ring may be relatively thick, the unstable wavelength and mass large, and the velocity dispersion inside each cloud also large. This latter effect follows from the virial condition for a cloud that has just formed by the gravitational instability, which gives external ~ 0.4cexternai (EKT93). Clouds forming in this way are virialized from birth, and virialized clouds with large internal velocity dispersions have large gravitational binding energies per unit mass. This makes the clouds very difficult to destroy by internal star formation. For example, the ambient interstellar medium in our Galaxy has a velocity dispersion of 7 to 10 km s" 1 and the largest Galactic clouds have internal dispersions of 3 to 5 km s" 1 . These clouds are destroyed by star-induced motions of 10 km s" 1 (Leisawitz et al. 1989), which presumably limits the total efficiency of star formation in them; i.e. stars continue to form in a cloud until the combined pressures from all of the stars drive motions at greater than cloud-escape speeds. When the internal cloud dispersion is large, such as 20 km s" 1 or larger as observed by Casoli et al. (1991) and EKT93 in two interacting galaxy pairs, then normal galactic star formation efficiencies will not be able to generate enough pressures and motions to disrupt the cloud. Instead, star formation should continue to a higher total efficiency of massive star formation until the more tightly bound cloud is eventually destroyed. Thus there should be a link between the internal velocity dispersion in a cloud and the ratio of the total mass of massive stars that form in the cloud to the gas mass at the time of cloud destruction. An observational confirmation of this predicted link between efficiency and cloud velocity dispersion requires a comparison between final star formation efficiencies in clouds with very different velocity dispersions but the same mass. This equal-mass requirement is necessary to minimize the effects of sampling statistics on the number of massive stars present (small-mass clouds contain fewer stars and therefore fewer massive stars and this leads to less disruptive pressures per unit mass in some clouds). However, such a comparison is not possible for normal disk clouds because they all have about the same pressures and so their masses scale with their internal dispersions. The comparison may only be possible for clouds of the same mass that form in regions with very different pressures and velocity dispersions. A higher efficiency for massive star formation should also lead to a higher thermal temperature in the cloud cores because of the corresponding large total embedded stellar luminosity per unit cloud mass, which raises the dust temperature and the gas temperature through gas-grain collisions (high molecular temperatures in ILR rings were observed by Tilanus et al. 1991; Mauersberger k Henkel 1991; Bergman et al. 1992; Wild et al. 1992). Higher gas temperatures could also result from higher cosmic-ray fluxes per unit cloud mass (Suchov et al. 1993), which also follows from the higher efficiency. The thermal Jeans mass scales with T3/2/pJoud, so this characteristic mass could increase too if Pcioud is not too much larger. For example, an increase in gas temperature from 10 K in normal disk star-forming regions to 50 K in ILR rings would increase the thermal Jeans

234

Elmegreen: Starburst and Ultraluminous Galaxies

mass because the density in the ring is only larger by a factor of 100. An increase in T by a factor of 30, to several hundred degrees, would be enough to increase the thermal Jeans mass in starburst nuclei, where the density is larger than in the disk by a factor of 104, as discussed above. Presumably such an increase in the thermal Jeans mass would raise the lower mass limit in the initial stellar mass function. Various such shifts in the IMF have been reported by Rieke et al. (1980, 1993), Sage et al. (1991), Puxley (1991), Bernlohr (1992, 1993) and others. A review of IMF changes in starburst galaxies is in Scalo (1990). The changes in e and Jeans mass with velocity dispersion c and pressure P can be estimated from simple theory. The critical assumption is that star formation continues inside a cloud complex until either the embedded stellar luminosity exceeds the energy dissipation rate, or the gas is turned into stars. The energy dissipation rate in ergs s" 1 for the whole cloud is oc Mc3/R, which is the kinetic energy, oc Me2, divided by the crossing time, oc R/c. But for a virialized cloud, R oc c/(GpY^2, and if we take P = pc2, the energy dissipation rate oc Mc{GP)xl2. If we set this equal to the rate at which turbulent energy is pumped into the cloud, which is the product of the luminosity •Tatars from embedded stars and a luminous to kinetic energy efficiency factor £, then ^ s t a r S / M g a s » c{GPfl2 « C2{Gpf'2. This result gives the observed star formation efficiency in local clouds if £ is chosen properly. For local clouds, the velocity dispersion, c, and density, p, scale with cloud mass, M, and total external pressure, P, as (Elmegreen 1989)

(5.8) for P 4 = P/10 4 crrr 3 K. Thus

(^r>(c.ns-')°3--.

(5.9)

The luminosity-to-mass ratio for young stars ranges between 1 and 104 LQ/MQ depending on stellar mass. For a massive cloud where massive stars form, istars/Mstars ~ 103 LQ/MQ. This gives Mgas

08

/ Mgas \ 1

U

J

/ 4

45 / 8

'

( }

Then the observed star formation efficiency of several percent results if the efficiency of conversion from stellar photon luminosity to cloud turbulent energy, £, is around 5 x 10~ 7 . Some of this inefficiency comes from the conversion of stellar photons into heat and hot expanding gas in HII regions and winds, and some of it comes from the conversion of the high-temperature expansion into motions of the cold gas. Most of the stellar photon luminosity comes out of the cloud in the form of infrared photons. The important point here is the way the mass efficiency, M s t a rs/M gas , scales with cloud mass (= Mgas) a n d pressure. If we set the cloud mass equal to the characteristic value from the instability model, M g a s « c 3 / ( G 3 / y / 2 ) = cA/{G^2Pll2), then -l

Mgas This result suggests that gas disks with high velocity dispersions and/or high pressures should produce stars with a relatively high efficiency in the clouds that form by gravitational instabilities. (Higher efficiencies also follow if only low-mass stars form, giving

Elmegreen: Starburst and Ultraluminous Galaxies 1

235

3

10 LQ/MQ instead of 10 LQ/MQ.) This c and P dependence implies that e in some starburst galaxies can be high, increasing the SFR for a given density, and that the gas consumption time can be even lower than estimated above. It also implies that more star formation should lead to bound clusters, rather than unbound stellar groups, because high e puts a large fraction of the cloud's binding mass in the form of stars, which can therefore remain bound when the gas disperses. The result suggests furthermore that the efficiency of star formation in clouds of a given mass should decrease with galactocentric radius as the pressure decreases.

6. Changes in propagating star formation The previous sections discussed a mode of star formation in which the ambient interstellar medium collapses by itself into giant self-gravitating cloud complexes. Stars then form in the dense cores of these complexes after the internal turbulent energy dissipates. This is a spontaneous mode of star formation. Another mode of star formation occurs when pressures from an existing star formation region push on the surrounding gas and produce an expanding shell or ring. If this shell or ring can accumulate enough matter and last long enough, it will collapse along its periphery into new self-gravitating clouds which will eventually form more stars. This is a stimulated mode of star formation. Any pressure source that sufficiently disturbs the interstellar medium can stimulate star formation. A recent analysis of the stability of expanding shells and rings gives the time and size scales for when the collapse occurs (Elmegreen 1994c): l

1/2

M~1/2

i?shell = | v t = WM^'cno1'3

pc

1241/2

M~1

RrinS = ^ t = 2Ucno1/2

Myr

Myr pc

(6.12)

(6.13) (6.14) (6.15)

In the last parts of these expressions, c is the velocity dispersion inside the shell, in km s" 1 , M is the ratio of the expansion speed, V, to c, and «o is the ambient density in cm" 3 . In normal galaxies, where the ambient density is relatively low, perhaps 1 to 5 cm" 3 , the collapse occurs at a relatively late time and the effective Mach number M is low. Then the ring or shell size is large, perhaps 50 to 200 pc, and the generation-to-generation propagation time is long, perhaps several tens of millions of years. In the inner region of a starburst galaxy, where the density can be very high, this situation changes. The time and distance scales for propagation can be very small, such as 1 Myr and 1 pc, and M. can still be high at the time of collapse. This high M is primarily because the time scale for collapse of a pressure-driven shell can be less than the evolutionary time of a massive star when the ambient density exceeds 103 cm" 3 , so the shock-driving pressures can be high when the shell collapse occurs. The value of M at the time of collapse can affect the qualitative appearance of star formation in a nuclear disk compared to a normal galactic disk. Normally the propagation time is comparable to the time scale for spontaneous instabilities, both being approximately (Gp)" 1 / 2 for ambient density p and low M. in the stimulated case. Then

236

Elmegreen: Starburst and Ultraluminous Galaxies

these two modes are competitive and should occur with a comparable frequency in a normal disk, as is observed (i.e. big shells with peripheral star formation and big clouds with interior star formation both occur in normal galaxy disks). But when M is high at the time of ring or shell collapse, the propagation time can be much less than the spontaneous time and the morphology of star formation can shift to favor the propagating mode. This implies that the cloud structure could be more ragged, shell-like or filamentlike in star-forming regions with very high densities. This is unlike the situation at low densities where star formation is sometimes concentrated in large, often regularly-spaced clouds. Observations of small (1 pc) dense clouds in a starburst galaxy were made by Aalto et al. (1991). 7. Conclusions Galaxy interactions change a galaxy's structure, interstellar medium, and star formation properties. The star formation mechanisms could be the same as in normal galaxies, but at a much higher rate and overall efficiency because of the higher gas density in the starburst region. A higher velocity dispersion and pressure should lead to higher star formation efficiencies inside each cloud, higher gas temperatures in regions of star formation and to a higher lower-mass limit in the initial mass function. Propagating star formation could be relatively more important in starburst galaxies where the density is very high. These conclusions follow primarily from the theory of star formation, but many of them have yet not been demonstrated from observations. Some of the observational questions that should be answered are: (1) Are gas velocity dispersions high in the starburst regions of interacting galaxies? Is c2 ~ Gfi in ILR rings? (2) Are cr/(rcriticai and Q" 1 comparable to 1 or larger than in the starburst regions? (3) Are there large inflow velocities from rapid gas accretion, visible perhaps along the minor axes of interacting galaxies? (4) Are the star formation efficiencies and initial mass functions different in starburst regions and normal galactic disks? (5) What are the magnetic field strengths in ILR rings? (6) What fraction of the molecular mass in a starburst region is in the form of diffuse clouds, as opposed to self-gravitating clouds, simply as a result of the high pressure? (7) Does the morphology of star formation differ in starburst regions compared to normal galactic disks as a result of a shift in the proportion of stars that form by spontaneous and stimulated processes? These questions can probably be answered with instrumentation that is available today. Thanks to John Scalo for critically reading the manuscript and offering many valuable comments, and for help with references on the variability of star formation in the Milky Way.

REFERENCES AALTO, S., BLACK, J. H., BOOTH, R. S. & JOHANSSON, L. E. B. 1991 A. & A. 247, 291. ARSENAULT, R. 1989 A. & A. 217, 66. ARSENAULT, R., BOULESTEIX, J., GEORGELIN, Y. & ROY, J.-R. 1988 A. & A. 200, 29.

ATHANASSOULA, E. 1992 Mon. Not. Royal Asti. Soc. 259, 345. BALBUS, S. A. 1988 Ap. J. 324, 60.

Elmegreen: Starburst and Ultraluminous Galaxies

237

BALBUS, S. A. 1990 In The Interstellar Medium in Galaxies (ed. H. A. Thronson, Jr. &; J. M. Shull), p. 305. Kluwer. BALBUS, S. A. & COWIE, L. L. 1985 Ap. J. 297, 61. BALZANO, V. A. 1983 Ap. J. 268, 602. BALZANO, V. A. & WEEDMAN, D. W. 1981 Ap. J. 243, 756.

BARNES, J. E. & HERNQUIST, L. E. 1991 Ap. J. Lett. 370, L65. BARNES,

J. E. & HERNQUIST, L. 1992 Nature 360, 715.

BARRY, D. 1988 Ap. J. 334, 436. BENEDICT, G. F., HIGDON, J. L., TOLLESTRUP, E. V., HAHN, J. M. & HARVEY, P. M. 1992

A. J. 103, 757. BENEDICT, G. F. ET AL. 1993 A. J. 105, 1369. BERGMAN, P., AALTO, S., BLACK, J. H. & RYDBECK, G. 1992 A. & A. 265, 403. BERNLOHR, K. 1992 A. & A. 263, 54. BERNLOHR, K. 1993 A. & A. 268, 25. BOER, B. & SCHULZ, H. 1993 A. & A. 277, 397. BONAZZOLA, S., FALGARONE, E., HEYVAERTS, J., PERAULT, M. & PUGET, J. L. 1987 A. &

A. 172, 293. BOTHUN, G. D., SCHOMBERT, J. M., IMPEY, C. D., SPRAYBERRY, D. & MCGAUGH, S. S.

1993 A. J. 106, 530. BROWN, R. L. & VANDEN BOUT, P. A. 1991 A. J. 102, 1956. BRUHWEILER, F. C , TRUONG, K. Q. & ALTNER, B. 1991 Ap. J. 379, 596. BUAT, V., DEHARVENG, J. M. & DONAS, J. 1989 A. & A. 223, 42.

BUTA, R. 1986a Ap. J. Suppl. 61, 609. BUTA, R. 1986b Ap. J. Suppl. 61, 631. BUTA, R. & CROCKER, D. A. 1993 A. J. 105, 1344. CASOLI, F., DUPRAZ, C , COMBES, F., & KAZES, I. 1991 A. & A. 251, 1. CHIANG, W.-H. & PRENDERGAST, K. H. 1985 Ap. J. 297, 507. CHIANG, W.-H. & BREGMAN, J. N. 1988 Ap. J. 328, 427. Cm FERNANDES, R., JR., DOTTORI, H. A., GRUENWALD, R. B. & VIEGAS, S. M. 1992 Mon.

Not. Royal Astr. Soc. 255, 165. CLEMENS, D. L., VAN DER WERF, P. P., KRABBE, A., BLIETZ, M., GENZEL, R. & WARD,

M. J. 1993 Mon. Not. Royal Astr. Soc. 262, L23. COLINA, L. &c PEREZ OLEA, D. 1992 Mon. Not. Royal Astt. Soc. 259, 709.

COMBES, F. 1988 In Galactic and Extragalactic Star Formation (ed. R. E. Pudritz & M. Fich), p. 475. Kluwer. COMBES, F. 1991 In IAU Symp. 146, Dynamics of Galaxies and their Molecular Cloud Distributions (ed. F. Combes & F. Casoli), p. 225. Kluwer. COMBES, F. & GEREST, M. 1985 A. & A. 150, 327. COMBES, F. GERIN, M. NAKAI, N. KAWABE R. & SHAW, M. A. 1992 A. & A. 259, L27. CONDON, J. J., CONDON, M. A., GISLER, G. &; PUSCHELL, J. J. 1982 Ap. J. 252, 102. CONDON, J. J., HUANG, Z.-P. YIN, Q. F. & THUAN, T. X. 1991 Ap. J. 378, 65. COWIE, L. L. 1981 Ap. J. 245, 66.

CRAM, L. E., NORTH, A., SAVAGE, A. 1992 Mon. Not. Royal Asti. Soc. 257, 602. DAHLEM

M.,

BOMANS

D. J.,

WILL

J.-M. 1994 Ap.J. Submitted.

DAVIDGE, T. J. 1992 Ap. J. 397, 457.

D., SOLOMAN, P. M. & RADFORD, S. J. E. 1993 Ap. J.Lett. 414, L13. S. RAWLINGS, S. PUXLEY, P., ROCCA VOLMERANGE, B. &C KUNTZ, K. 1993 Nature 363, 140.

DOWNES, EALES,

238

Elmegreen: Starbursi and Ultraluminous Galaxies

ECKART, A., CAMERON, M., JACKSON, J. M., GENZEL, R., HARRIS, A. I., WILD, W. & ZINNECKER, H. 1991 Ap. J. 372, 67. ELMEGREEN, B. G. 1979 Ap. J. 231, 372.

B. G. 1987 In IAU Symposium No. 115, Star Forming Regions (ed. M. Peimbert & J. Jugaku), p. 457. Reidel. ELMEGREEN, B. G. 1989a Ap. J. 338, 178. ELMEGREEN, B. G. 1989b Ap. J. 344, 306. ELMEGREEN, B. G. 1990a In The Evolution of the Interstellar Medium, (ed. L. Blitz), p. 247. Astronomical Society of the Pacific Publishers. ELMEGREEN,

ELMEGREEN, B. G. 1990b, Ap. J. 357, 125. ELMEGREEN, B. G. 1991 Ap. J. 378, 139.

ELMEGREEN, B. G. 1993 In Protostars and Planets HI (ed. E. H. Levy & J. I. Lunine), p. 97. University of Arizona Press. ELMEGREEN, B. G. 1994a 'Supercloud Formation by Gravitational Collapse in the Crest of a Spiral Density Wave' Ap. J. Submitted. ELMEGREEN, B. G. 1994b 'Starbursts by Gravitational Collapse in the Inner Lindblad Resonance Rings of Galaxies' Ap. J. In press. ELMEGREEN, B. G. 1994C 'A Q Condition for Long Range Propagating Star Formation' Ap. J. Submitted. ELMEGREEN, B. G. & ELMEGREEN, D. M. 1983 Mon. Not. Royal Astr. Soc. 203, 31. ELMEGREEN, B. G., KAUFMAN, M. & THOMASSON, M. 1993 Ap. J. 412, 90. ELMEGREEN, B. G. & THOMASSON, M. 1993 A. & A. 272, 37.

ELMEGREEN, D. M. & ELMEGREEN, B. G. 1984 Ap. J. Suppl. 54, 127. ELMEGREEN, D. M., ELMEGREEN, B. G. &; BELLED, A. D. 1990 Ap. J. 364, 415. ELMEGREEN,

D. M., ELMEGREEN, B. G.,

LANG,

C. & STEPHENS, C. 1994 Ap. J. In press.

FEINSTEIN, C , VEGA, I., MENDEZ, M. & FORTE, J. C. 1990 A. & A. 239, 90. FELDMAN, F. R., WEEDMAN, D. W., BALZANO, V. A. & RAMSEY, L. W. 1982 Ap. J. 256,

427. FORBES, D. A., BOISSON, C. & WARD, M. J. 1992 Mon. Not. Royal Astr. Soc. 259, 293. FORBES, D. A., WARD, M. J., ROTACIUC, V., BLIETZ, M., GENZEL, R., DRAPATZ, S., VAN

DER WERF, P. P., KRABBE, A. 1993 Ap. J. Lett. 406, Lll.

J. 1993 Rev. Mex. Astron. Astiofis. 26, 13. J. 1991 In Chemical and Dynamical Evolution of Galaxies (ed. F. Ferrini, J. Franco &: F. Matteucci), p. 506. Pisa: ETS Editrice. GALLAGHER, J. S. & HUNTER, D. A. 1984 Ann. Rev. Astron. Astrophys. 22, 37. FRANCO,

FRANCO,

GARCIA-BARRETO, J. A., DOWNES, D., COMBES, F., GERIN, M., MAGRI, C , CARRASCO, L. & CRUZ GONZALEZ, I. 1991a A. & A. 244, 257. GARCIA-BARRETO, J. A., DETTMAR, R. J., COMBES, F., GERIN, M. & KORIBALSKI, B. 1991b

Rev. Mex. Astron. Astrofis. 22, 197. GEHRZ, R. D., SRAMEK, R. A. & WEEDMAN, D. W. 1983 Ap. J. 267, 551. GERIN, M., CASOLI, F. &; COMBES, F. 1991 A. & A. 251, 32. GEROLA, H., SEIDEN, P. E. & SCHULMAN, L. S. 1980 Ap. J., 242, 517. GOLDREICH, P. & LYNDEN-BELL, D. 1965 M. N. R. A. S. 130, 97. GOMEZ, A. E., DELHAYE, J., GRENIER, S., JASCHEK, C , ARENOU, F. &; JASCHEK, M. 1990

A. & A. 236, 95. GUIDERDONI, B. 1987 A. & A. 172, 27.

HABE, A. & OHTA, K. 1992 Pub. Astr. Soc. Japan 44, 203.

HAMANN, F. & FERLAND, G. 1992 Ap. J. Lett. 391, L53. HAWARDEN, T. G., MOUNTAIN, C. M., LEGGETT, S. K. & PUXLEY, P. J. 1986 Mon.

Not.

Elmegreen: Starburst and Ultraluminous Galaxies

239

Royal Astr. Soc. 221, 41p. HENKEL, C , BAAN, W. A. & MAUERSBERGER, R. 1991 A. k A. Rev. 3, 47.

HERNQUIST, L. 1990 In Dynamics and Interactions of Galaxies (ed. R. Wielen), p. 108. Springer. HOWARD, S. & BYRD, G. G. 1990 A. J. 99, 1798. HUMMEL, E. 1980 A. k A. 89, LI. HUMMEL, E., VAN DER HULST, J. M. & KEEL, W. C. 1987 A. & A. 172, 32. HUMMEL, E. & JORSATER, S. 1992 A. k A. 261, 85. HUNTER, D. A. & GALLAGHER, J. S. Ill 1986 P. A. S. P. 98, 5.

HUNTER, D. A. 1992 inStar Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto & F. Sanchez), p. 67. Cambridge University Press. JACKSON, J. M., ECKART, A., CAMERON, M., WILD, W., HO, P. T. P., POGGE, R. W. & HARRIS, A. I. 1991 Ap. J. 375, 105. JARVIS, B. J. & MELNICK, J. 1991 A. & A. 244, Ll. JOG, C. J. & SOLOMON, P. M. 1992 Ap. J. 387, 152. JOG,

C. J. &

DAS,

M. 1992 Ap. J. 400, 476.

KENNICUTT, R. C. 1989 Ap. J. 344, 685. KENNICUTT, R. C , KEEL, W. C , VAN DER HULST, J. M., HUMMEL, E. & ROETTIGER, K. A.

1987 A. J. 93, 1011. KOORNNEEF, J. 1993 Ap. J. 403, 581. KORIBALSKI, B., DICKEY, J. M. & MEBOLD, U. 1993 Ap. J. Lett. 402, L41.

KORMENDY, J. & SANDERS, D. B. 1992 Ap. J. Lett. 390, L53.

R. B. 1987 In Starburst and Galaxy Evolution (ed. T. X. Thuan, T. Montmerle & J. Tran Thanh Van), p. 467. Gif sur Yvette: Editions Frontieres. LARSON, R. B. 1988, In Galactic and Extragalactic Star Formation (ed. R. E. Pudritz & M. Fich), p. 459. Kluwer. LARSON, R. C. 1992, In Star Formation in Stellar Systems, (ed. G. Tenorio-Tagle, M. Prieto &; F. Sanchez), p. 381. Cambridge University Press. LARSON,

LARSON, R. B. & TINSLEY, B. M. 1978 Ap. J. 219, 46.

LAWRENCE, A. ET AL. 1993 Mon. Not. Royal Astr. Soc, 260, 28. LEISAWITZ, D., BASH, F. N. & THADDEUS, P. 1989 Ap. J. Suppl. 70, 731. LIPARI, S. &; MACCHETTO, F. 1992 Ap. J. 387, 522. LIPARI, S., TERLEVICH, R. & MACCHETTO, F. 1993 Ap. J. 406, 451. LONSDALE, C. J., LONSDALE, C. J. & SMITH, H. E. 1992 Ap. J. 391, 629. LONSDALE, C. J., SMITH, H. E. &C LONSDALE, C. J. 1993 Ap. J. Lett. 405, L9. MAJEWSKI,

S. R. 1993 Ann. .Rev. Astron. Astrophys. 31, 575.

MAJEWSKI, S. R., HERELD, M., KOO, D. C , ILLINGWORTH, G. D. & HECKMAN, T. M. 1993

Ap. J. 402, 125. MARSAKOV, V. A., SHEVELEV, YU. G. & SUCHKOV, A. A. 1990 Ap. Sp. Sci. 172, 51. MATSUDA,

T. &

NELSON,

A. H. 1977 Nature 266, 607.

MAUERSBERGER, R., HENKEL, C , WALMSLEY, C. M., SAGE, L. J. & WIKLIND, T. 1991 A. &

A. 247, 307. MAZZARELLA, J. M., GAUME, R. A., SOIFER, B. T., GRAHAM, J. R., NEUGEBAUER, G. & MATTHEWS, K. 1991 A. J. 102, 1241. MAUERSBERGER, R. & HENKEL, C. 1991 A. & A. 245, 457. MEUSINGER, H. 1990 Ap. Sp. Sci. 182, 19

MEUSINGER, H. 1991 Asti. Nachi. 312, 231. Mfflos, J. C , RlCHSTONE, D. O. &; BOTHUN, G. D. 1992 Ap. J. 400, 153. Mmos, J. C , BOTHUN, G. D. & RICHSTONE, D. O. 1993 Ap. J. 418, 82.

240

Elmegreen: Starburst and Ultraluminous Galaxies

MIRABEL,

I. F. 1991 ESO Messenger 63, 64.

MIRABEL, I. F., LUTZ, D. & MAZA, J. 1991 A. & A. 243, 367. MIRABEL, I. F., DOTTORI, H. & LUTZ, D. 1992 A. & A. 256, L19. MORGAN, W. W. 1958 P. A. S. P. 70, 364. MORRIS, M. & RICKARD, L. J. 1982 Ann. Rev. Astron. Astiophys. 20, 517. NOGUCHI, M. 1987 Mon. Not. Royal Astt. Soc. 228, 635. NOGUCHI, M. 1988 A. & A. 201, 37.

NoH, H.-R. & SCALO, J. 1990 Ap. J. 352, 605. NORRIS, R. P., ALLEN, D. A., SRAMEK, R. A., KESTEVEN, M. J. &. TROUP, E. R. 1990 Ap.

J. 359, 291. OLSON, K. M. &C KWAN, J. 1990 Ap. J. 361, 426. OSTERBROCK, D. E. 1993 Ap. J. 404, 551.

T., VAZQUEZ-SEMADENI, E. C. & Astrophysics (ed. J. Franco). In press.

PASSOT,

POUQUET,

A. 1994 In Numerical Simulations in

PASTORIZA, M. G. 1975 Ap. Sp. Sci. 33, 173. PUDRITZ, R. E. 1990 Ap. J. 350, 195.

PuXLEY, P. J. 1991 Mon. Not. Royal Asti. Soc. 249, l i p . QUIRK, W. J. 1972 Ap. J. Lett. 176, L9. RAND, R. J. & KULKARNI, S. R. 1990 Ap. J. Lett. 349, L43. RIEKE, G. H., LEBOFSKY, M. J., THOMPSON, R. L, LOW, F. J. & TOKUNAGA, A. T. 1980

Ap. J. 238, 24. RIEKE, G. H., LOKEN, K., RIEKE, M. J. & TAMBLYN, P. 1993 Ap. J. 412, 99. Rix, H.-W. & RIEKE, M. J. 1993 Ap. J. 418, 123. ROMAN, N. G. 1952 Ap. J. 116, 122. ROWAN ROBINSON,

M. ET AL. 1991 Nature 351, 719.

ROWAN ROBINSON, M. ET AL. 1993 Mon. Not. Royal Astr. Soc. 261, 513. SAGE, L. J., MAUERSBERGER, R. & HENKEL, C. 1991 A. & A. 249, 31.

A. & BEDKE, J. 1988 Atlas of Galaxies. U.S. Govt. Printing Office. Washington, D.C., NASA.

SANDAGE,

SANDERS, D. B., SOIFER, B. T., ELIAS, J. H., MADORE, B. F. MATTHEWS, K., NEUGEBAUER, G. & SCOVILLE, N. Z. 1988 Ap. J. 325, 74.

SARGENT, A. &; SCOVILLE, N. 1991 Ap. J. Lett. 366, LI.

J. M. 1987a In Starbursts and Galaxy Evolution (ed. T. X. Thuan, T. Montmerle & J. T. Thanh Van), p. 445. Editions Frontieres. ScALO, J. M. 1987b In Evolution of Galaxies, Proc.lOth European Regional Astronomy Meeting of the IAU, Vol.4 (ed. J. Palous), p. 101. Astronomical Institute of Prague. SCALO, J. M. 1990 In Windows on Galaxies (ed. G. Fabbiano, J. S. Gallagher &: A. Renzini), p. 125. Kluwer. SCALO,

SCALO, J. M. & STRUCK-MARCELL, C. 1984 Ap. J. 276, 60. SCHMIDT, M. 1959 Ap. J. 129, 243. SEARLE, L., SARGENT, W. L. W. & BAGNUOLO, W. G. 1973 Ap. J. 179, 427.

SEDDEN, P. E. & GEROLA, H. 1982 Fund.Cosmic Phys. 7, 241.

SEroEN, P. E. & SCHULMAN, L. S. 1990 Advances in Physics 39, 1. SEKIGUCHI, K. & WOLSTENCROFT, R. D. 1993 Mon. Wot. Royal Astr. Soc. 263, 349. SERSIC, J. L. & PASTORIZA, M. 1965 P. A. S. P. 77, 287. SERSIC, J. L. & PASTORIZA, M. 1967 P. A. S. P. 79, 152. SMITH, E. P. & KASSIM, N. E. 1993 A. J. 105, 46. SOFUE,

Y. 1991 Pub. Astr. Soc. Japan 43, 671.

Elmegreen: Starburst and Ultraluminous Galaxies

241

SOFUE, Y. & HABE, A. 1992 Pub. Asti. Soc. Japan 44, 325. SOFUE, Y. & IRWIN, J. A. 1992 Pub. Astr. Soc. Japan 44, 353.

SOIFER, B. T. ET AL. 1991 Ap. J. Lett. 381, L55. SOLOMON, P. M., DOWNES, D. & RADFORD, S. J. E. 1992 Ap. J. Lett. 398, L29.

SOPP, H. M. ALEXANDER, P. 1991 Mon. Not. Royal Astr. Soc. 251, 112. STANFORD, S. A. 1990 Ap. J. 358, 153. STARK, A. A., ELMEGREEN, B. G. & CHANCE, D. 1987 Ap. J. 322, 64. STOCKTON, A. & RIDGWAY, S. E. 1991 A. J. 102, 488. STRUCK-MARCELL, C. & SCALO, J. M. 1984 Ap. J. 277, 132. SUCHOV, A., ALLEN, R. J. & HECKMAN, T. M. 1993 Ap. J. 413, 542. TACCONI, L. J. TACCONI-GARMAN, L. E. THORNLEY, M. & VAN WOERDEN, H. 1991 A. & A.

252, 541. TELESCO, C. M. 1988 Ann. Rev. Astron. Astrophys. 26, 343. TELESCO, C. M. & GATLEY, I. 1981 Ap. J. Lett. 247, Lll. TELESCO, C. M. & GATLEY, I. 1984 Ap. J. 284, 557. TELESCO, C. M., CAMPINS, H., JOY, M., DIETZ, K. & DECHER, R. 1991 Ap. J. 369, 135. TELESCO, C. M., DRESSEL, L. L. & WOLSTENCROFT, R. D. 1993 Ap. J. 414, 120. TENJES, P. & HAUD, U. 1991 A. k A. 251, 11.

TERLEVICH, R. J. & BOYLE, B. J. 1993 Mon. Not. Royal Astr. Soc. 262, 491.

TERLEVICH, R., TENORIO-TAGLE, G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. Royal Astr.

Soc. 255, 713. TILANUS, R. P. J. ET AL. 1991 Ap. J. 376, 500. VAN DEN BROEK, A. C. 1992 A. & A. 261, Ll. VAN DEN BROEK, A. C. 1993 A. k A. 269, 96. VAN DER HULST, J. M. 1979 A. k A. 71, 131. VAN DER HULST, J. M., SKILLMAN, E. D., SMITH, T. R., BOTHUN, G. D., MCGAUGH, S. S. & DE BLOK, W. J. G. 1993 A. J. 106, 548.

VAN DER WERF, P. P., GENZEL, R., KRABBE, A., BLIETZ, M., LUTZ, D., DRAPATZ, S., WARD, M. J. & FORBES, D. A. 1993 Ap. J. 405, 522. VAZQUEZ, E. C. & SCALO, J. M. 1989 Ap. J. 343, 644. VIALLEFOND,

F., Goss, W. M. & ALLEN, R. J. 1982 A.
VOGEL, S. N., RAND, R. J., GRUENDL, R. A. & TEUBEN, P. J. 1993 P. A. S. P. 105, 666.

WADA, K. & HABE, A. 1992 Mon. Not. Royal Astr. Soc. 258, 82. WALLER, W. H., GURWELL, M. & TAMURA, M. 1992 A. J. 104, 63.

WANG, B. & SILK, J. 1994 Ap. J. In press. WEEDMAN, D. W. 1983 Ap. J. 266, 479. WEEDMAN, D. W., FELDMAN, F. R., BALZANO, V. A. RAMSEY, L. W., SRAMEK, R. A. &

Wu, C.-C. 1981 Ap. J. 248, 105. WIKLIND, T. &; HENKEL, C. 1992 A. k A. 257, 437. WILD, W., HARRIS, A. I., ECKART, A., GENZEL, R., GRAF, U. U. JACKSON, J. M., RUSSELL, A. P. G. & STUTZKI, J. 1992 A. k A. 265, 447. WILSON, C. D. 1992 Ap. J. 391, 144. WILSON, A. S., HELFER, T. T., HANIFF, C. A. & WARD, M. J. 1991 Ap. J. 381, 79. WINDHORST, R. A. ET AL. 1991 Ap. J. 380, 362.

WOLSTENCROFT, R. D., TULLY, R. B. & PERLEY, R. A. 1984 Mon. Not. Royal Astr. Soc.

207, 889. WYNN-WILLIAMS, C. G. & BECKLIN, E. E. 1993 Ap. J. 412, 535.

242

Elmegreen: Starburst and Ultraluminous Galaxies

WYSE, R. F. G. 1986 Ap. J. Lett. 311, L41. WYSE, R. F. G. & SILK, J. 1989 Ap. J. 339,

700.

Yoss, K. M. 1992 A. J. 104, 327. ZARITSKY, D. & LORRIMER, S. J. 1993 In Evolution of Galaxies and their Environment (ed. D. J. Hollenbach, H. A. Thronson &; J. M. Shull), p. 82. NASA CP-3190. ZwiCKY, F. 1959 Handbuck dei Physik 53, 373. ZWICKY, F. & HUMASON, M. L. 1960 Ap. J. 132,

627.

ZWICKY, F. & HUMASON, M. L. 1961 Ap. J. 133,

794.

Colliding Galaxies, Shocked Gas, and Violent Star Formation 1 By SUSAN A. LAMB1, RICHARD A. GERBER ! 2 AND DINSHAW S. BALSARA {

departments of Physics and of Astronomy, University of Illinois, Urbana, IL 61801, USA 2

Department of Physics and Astronomy, Johns Hopkins University, Baltimore, MD 21218, USA

Collisions of galaxies are often observed to produce increases in the far-IR flux and star formation rates, as compared to those seen in isolated galaxies. It is expected that the star formation taking place in these systems occurs under more violent circumstances than those found in the spiral arms of disk galaxies, for example. We will present some results from a study in which we have produced combined N-body/Smooth Particle Hydrodynamics simulations of collisions of galaxies, looking for regions in which shocks develop and where the gas density gets high enough that new star formation might be expected to take place. These results give us a preliminary idea of what the properties of the shocked, high density regions might be and provide a basis for the future inclusion of violent star formation into such calculations with a view to eventually explaining the observed enhancements.

1. Introduction In seeking out places in the present Universe where star formation might be expected to occur under more violent circumstances than those that occur in the arms of spiral galaxies one is drawn to consider the interiors of gas-rich galaxies that are undergoing a collision with another galaxy. Although not extremely common at this epoch, such systems of galaxies have been detected and extensively studied during recent years. There are now considerable data available to suggest that these systems commonly experience increased star formation, this taking place most often in the nuclei of the disk galaxies, rather than in their spiral arms. Surveys of optically selected interacting galaxy systems in which at least one of the original galaxies had a disk show strong evidence for increased star formation in many cases (see Keel et a/.1985; Bushouse 1986, 1987), although not in all. Approximately 30% show no significant star formation in the last 109 yr. On average the IRAS 25—//m and 60—/im fluxes from galaxies in these same optically selected samples is double that of isolated spirals (Kennicutt et al. 1987; Bushouse, Lamb &, Werner 1988), although, as with the optical indicators of star formation, there is sufficient spread within the samples to be consistent with very little recent star formation in some galaxies. (Also see Lonsdale, Persson, &; Matthews 1984 with IR results from the IRAS minisurvey). Thus increased star formation is not an automatic consequence of a strong interaction. However, in some interacting or merging systems the increase in the rate of star formation appears to be greatly increased, as in the ultraluminous IRAS galaxies (see Sanders et al. 1988). A question of interest and great relevance to the subject of current enquiry, namely violent star formation, is whether the star formation taking place in the strongly interacting galaxies is taking place under circumstances sufficiently different from those observed at other locations to produce a detectable difference in such parameters as the upper and f Present address: NASA Ames Research Center, Theoretical Physics Branch, MS 245-3, Moffett Field, CA 94035, USA. t Present address: National Center for Supercomputer Applications, Beckman Institute, University of Illinois, Urbana, IL 61801, USA. 243

244

Lamb et al.: Violent Star Formation

lower mass cut-offs, the slope of the initial mass function (IMF), and the morphological distribution of stars on the scale of tens of parsecs. This subject awaits further exploration. It may be, for example, that the 'violence' of the environment (where I here take 'violence' to be a measure of the energy content in mass motions of the ambient material) has a crucial effect on the distribution of new stars in space but none on the distribution in mass; that is, if stars can form at all they do so in such a way that the IMF is universal. There are many interesting questions here concerning star formation, and the study of the conditions of the gas and the nature of the stars produced in interacting and merging galaxies will help to illuminate this subject in general by providing a somewhat different environment within which to study the phenomenon. The nuclei of galaxies in which most of the star formation in colliding galaxies and mergers is taking place are crowded and sometimes very confused morphologically. This crowding makes it difficult to get information on some aspects of the star formation history in such regions, especially concerning possible sequential triggering of star formation from one star forming region to another, or effects on the star formation timescale due to changes in the large-scale environment. However, there is one type of collision product, the ring galaxies, that has a very simple morphology and extensive star formation away from the nucleus, in a comparatively luminous material ring. These rare galaxies are ideal candidates for the study of 'violent' star formation because modelling shows that the gas streaming motions in these rings can become very high and can produce very strong shocks (see Gerber 1993; Gerber, Lamb & Balsara 1992, 1994a; Lamb, Gerber k Balsara 1993, 1994). 2. Collision-induced ring galaxies As a class, the collision-induced ring galaxies also includes galaxies which exhibit an arc of high density material and enhanced star formation (produced by an off-center collision). Sometimes a nucleus is present within the ring or arc, often displaced from the center or focus, respectively. In many cases the ring appears empty of a nucleus, although evidence suggests that a nucleus may still exist, buried in the arc or ring, as viewed by us. This general class of ring galaxies has been designated P by Few & Madore (1986) and as RING by de Vaucouleurs et al. (1976). The class has been further subdivided by Few k Madore according to morphological characteristics, such as dumpiness in the ring or arc. These subclassifications may be very useful when detailed comparisons are made between observations and models in future work, but we will not discuss them further here. Lynds k Toomre (1976) were the first to use numerical simulations to demonstrate that ring galaxies are a natural outcome of some galaxy collisions. Their early work has subsequently been amplified and built upon by Appleton k Struck-Marcell (1987a), Huang k Stewart (1988), Struck-Marcell k Lotan (1990), Struck-Marcell k Higdon (1993), Hernquist k Wiel (1993), Gerber, Lamb k Balsara (1992, 1994a), and Gerber k Lamb (1994). Observational evidence for the hypothesis that some ring galaxies are collisionally induced was obtained by Theys k Spiegel (1976) when they found that bright ring galaxies usually have a nearby second galaxy lying along a line aligned with the apparent minor axis of the ring. Further study has found that vigorous star formation is taking place in these rings (Theys k Spiegel 1976, 1977; Fosbury k Hawarden 1977; Thompson k Theys 1978; Jeske 1986; Appleton k Struck-Marcell 1987b; Schultz et al. 1991; Higdon 1992; Marcum, Appleton k Higdon 1992). Elevated far-IR colors, bright HII emission, and blue spectral colors are jointly taken as evidence for recent star formation. A color gradient between the ring and center (blue to red) has been found in the Cartwheel

Lamb et al.: Violent Star Formation

245

galaxy (see Higdon 1993) which has been attributed to excess star formation occuring in the ring as it propagated outwards through the disk. As mentioned earlier, these bright ring galaxies are rare. An interpenetrating collision of two galaxies, of which at least one must be a gas-rich disk galaxy, with a mass ratio of very approximately 10:1 or less, is required to produce the ring/arc geometry. Depending on the central mass concentration of the two galaxies involved, an impact parameter of no more than approximately half of the disk radius is needed. Further, the trajectory of the galaxies before collision cannot be inclined at a very large angle to the normal to the disk plane. Add to this that the intrinsic timescale for the propagation of the ring disturbance out through the disk is a few times 108 years (this number depends upon the mass and size of the galaxy involved, see Gerber, Lamb & Balsara 1994a) and we can understand the paucity of such systems at this epoch. (The impact parameter range for the formation of ring and arc structures is investigated using a kinematic model by Gerber & Lamb (1994), and the connection between these morphologies and the more usual two-armed spiral structure and the rare one-armed spiral structure is explored.) In summary, both the observations and the numerical simulations of the bright ring galaxies strongly suggest that they provide a unique laboratory for star formation studies. For example, from numerical simulations we can determine such things as the time scales for ring formation and propagation, the connections between the different observed morphologies of rings and arcs and the collision parameters (impact parameter, angle of attack, and galactic mass ratio), the possible locations of 'missing nuclei', and the disposition of the gas (density and velocity as a function of position and time). Dynamical and kinematic model simulations, the former including a representation of the gas, can then be compared to observations of a variety of these systems. In particular, one can look for regions of activity in real galaxies, that is, regions with excess star formation or emission from hot gas, and compare their locations and properties with those regions in the models that one might expect to provide a good environment for star formation because of high density or shocks. It is anticipated that such comparisons will lead to a clearer understanding of the conditions needed to produce the excess star formation observed in these galaxies and allow a comparison between these conditions and those observed in our nearby environment, i.e. in star forming regions in our own galaxy, in the spiral arms of neighbours, and in such giant star forming regions as 30 Dor in the LMC.

3. Model galaxy collisions Both collisions between an elliptical and a disk galaxy and between two disk galaxies can lead to the formation of rings or arcs in the disk(s). In order to explore the dynamics of the formation of a ring it is necessary to have only one disk in the system, so we here limit our discussion to collisions of the former type. The strongest effects are obtained when the two galaxies are of equal mass, so here we focus on such models. We simulated the collisions using a combined N-body/Smooth Particle Hydrodynamics (SPH) code (see Balsara 1990) to follow stars, "dark matter", and gas in the encounter. The calculations were performed on a Cray 2 computer and were fully three dimensional, with all components gravitationally active. We used a King (1966) model to represent the stars and dark matter of the elliptical. It is spherical and gasless and can be thought of as an E0. The model for the halo of stars and dark matter in the disk galaxy were also produced from a King model. An exponential disk of stars was introduced into the King potential slowly, so that the overall potential could gradually adjust. The disk galaxy can be thought of as an Sc. The disk has a mass two fifths that of the halo and comprises

246

Lamb et al: Violent Star Formation

both stars and gas, the latter having one tenth of the disk mass. The disk galaxy is rotating and the elliptical impacts the disk at varying impact parameters (between 0 and 0.9 of the disk radius) along a trajectory parallel to the rotation axis of the disk. We incorporate a very simplified representation of the gas in that we include only one phase in which the density is given by 0.025 amu cm" 3 < r g a s (in p i ane ) < 2 amu cm" 3 , and for which the temperature is given by 8000 K < T/fj, < 6 x 105 K. It is assumed that the gas is isothermal, that is, that the time scale for radiation processes that cool the gas behind shock fronts are shorter than the calculational time step. Full details of the computational method, the calculations, and the results can be found in Gerber (1993). In order to discuss our results in a manner that is helpful when comparing them to observations we introduce a scaling. In a comparison of the models with observations the scale is chosen, in general, to fit some known parameter of a particular system of galaxies, such as a mass or a timescale. Here we will scale our galaxies' masses to approximately that of the Milky Way. We take M ga i = 1.75 x 1011 M©, then M ha i o = 1.25 x 10 11 M 0 , Mdisk = 5 x 10 10 MQ , and M g a s = 5 x 109 M©. With this mass scaling, the computational time unit = 2 x 106 yr, the initial separation of the galaxies is approximately 30 kpc, and the relative velocity is approximately 450 km s" 1 .

4. Results Our simulations confirm that impacts perpendicular to the plane of the disk galaxy or at some modest angle to that direction lead to the formation of a full ring (head-on), or partial ring (off-center), in both the stars and the gas of the disk. Shortly after closest approach, a ring or arc of excess density forms and propagates outwards. This is not just a density wave but is a combination of a wave motion with a material outflow in the disk. Both high densities and high relative velocities are achieved in the ring or arc because it is in this region that outflowing material from the inner disk impacts still infalling material from the outer disk. Thus the variation in the dynamical response time across the disk from center to edge is crucial in producing the ring formation. A way of discussing this phenomenon is in terms of the epicyclic frequency, /c, of orbiting particles which decreases with r, the original radius of an orbit. The passage of the intruder galaxy through the disk galaxy causes an impulse toward the impact point (the easiest to visualise is for a head-on collision, in which case the impact is at the center of the disk). Because of the temporary increase in the gravitational pull toward that point, the whole galaxy contracts towards it. Once the intruder has passed through, the gravitational pull is reduced to its original value and the orbits respond by expanding outwards, past their equilibrium positions, so that the galaxy eventually becomes distended as a result of the interaction. The timescale for the contraction and subsequent expansion at any orbit radius is given by K. The dependance on r gives rise to orbit crowding and this leads to a ring of excess density. In other words the inner orbits will reach their outer extent while orbits further out are still contracting. Such phenomena have been described by Struck-Marcell (1990) and Struck-Marcell & Lotan (1990) in terms of 'cusp' formation. The disturbance propagates outward through the galaxy disk on a similar timescale to the average rotation period, in our chosen scaling this is a few times 108 yr. This leads to a coupling between the two effects which can be observed in the simulations. That is, in a time sequence of the models that comprise a simulation one sees a spiralling of stars inwards and then outwards. In somewhat asymmetric situations, the spiralling can be easily detected in gross features of the morphology. An important result of our calculations is the determination of the three-dimensional

Lamb et al: Violent Star Formation

247

effects in the flows. We detect structure in the ^-direction which was not studied previously. If we picture the intruder impacting the disk from the top, we find that the outflowing, inner disk material flows over the still inwards falling, outer material during the propagation out through the disk of the ring or arc structure. This leads to a somewhat complicated thdimensional structure in which there are at least two shock fronts, one on the underside of the ring and the other towards the top and the leading edge of the ring. As a consequence, the exact strengths of the shocks in these regions are difficult to determine but at least Mach 10 is achieved. If star formation is closely connected in space in these galaxies with the locations of shock fronts, having in mind the circumstances explored in Mouschovias, Shu & Woodward (1974), then the star formation may be triggered primarily near the site of the stronger, lower shock front which is near to the inner edge of the ring. This is highly speculative, but it does point out that the three-dimensional structure in the ring is complicated and that simple assumptions about the potential location of star formation triggers might be inappropriate. The three dimensional structure also leads to a lack of correlation between the locations of high surface density and those of high volume density in the gas. This is of particular importance when considering the locations of possible star formation. The highest surface density remains at the nucleus throughout the collision. However, the highest volume densities are achieved in the ring or arc, once these structures are formed. There is a relative compression in the ring but the nucleus expands very shortly after closest approach. Thus on this basis alone, one would predict from the models that increased star formation resulting from the collision would occur in the ring rather than the nucleus as in many other impacted galaxies. The volume density in the gas can increase by a factor of about 100 in the ring while the surface density in this region increases by a factor 4-8, over the original density in the disk (the maximum is about a factor of 20 increase). For the stars in the ring the increase in surface density is only a factor 3-4. As one might expect, the gas ring is more peaked and somewhat narrower than the stellar ring and the two rings move outwards together at a velocity of about 100 km s" 1 (using the scaling given), the gas only lagging the stars towards the end of the simulation, when both the gas and stellar rings have broadened and are close to merging back into the disk. A peak inflow velocity of approximately 160 km s" 1 is achieved which gives rise to a relative velocity between the inflowing and outflowing streams of approximately 320 km s" 1 . This is equivalent to a Mach 10 shock under the relevant physical conditions. The maximum azimuthal velocity of 450 km s - 1 (to be compared with the average initial rotation rate of approximately 250 km s" 1 ), occurs at a time 2.2 x 107 years after closest approach; close to the time when maximum infall velocity is achieved. The maximum azimuthal velocity at any time is always in the ring. At the end of the simulations there is a flat rotation curve over much of the disk interior to the ring, typically with a velocity of about 130 km s~1. Our simulations also illustrate that off-center collisions lead to off-center nuclei, which can appear to be buried in the ring, for a reasonably large range of viewing angles (see Gerber, Lamb & Balsara 1992). In general, we show that with a full range of impact parameters in the collisions and a wide range of viewing angles the gross morphology of P/RING galaxies can be obtained (see Gerber, Lamb & Balsara 1994b). The full results of this investigation are given in Gerber (1993), and in Gerber, Lamb k Balsara (1992, 1994a, 1994b). SAL gratefully acknowledges the support of NASA grant NAG5-1241 and University of Illinois Research Board grants. RAG thanks NASA for a graduate student training

248

Lamb et al: Violent Star Formation

grant, NGT 70041. The numerical calculations were performed at the National Center for Supercomputing Applications at the University of Illinois at Urbana-Champaign, which is funded by the National Science Foundation.

REFERENCES P. N. & STRUCK-MARCELL, C. 1987a Ap. J. 312, 566. APPLETON, P. N. & STRUCK-MARCELL, C. 1987b Ap. J. 318, 103. BALSARA, D. S. 1990 PhD thesis, University of Illinois at Urbana Champaign. APPLETON,

BUSHOUSE, H. A. 1986

A. J. 91, 255.

BUSHOUSE, H. A. 1987

Ap. J. 320, 49.

BUSHOUSE, H. A., LAMB, S. A. & WERNER, M. W. 1988

Ap. J. 335, 74.

DE VAUCOULEURS, G., DE VAUCOULEURS, A. &; CORWIN, H. 1976 In Second Reference Catalogue of Bright Galaxies University of Texas Press. FEW, J. M. A. & MADORE, B. F. 1986

M. N. R. A. S. 222, 673.

FOSBURY, R. A. E. & HAWARDEN, T. G. 1977

M. N. R. A. S. 178, 473.

GERBER, R. A. 1993 PhD thesis, University of Illinois at Urbana Champaign. GERBER, R. A., LAMB, S. A. & BALSARA, D. S. 1992 GERBER, GERBER, GERBER,

Ap. J. Lett. 399, L51.

R. A. & LAMB, S. A. 1994 Ap. J. Submitted. R. A., LAMB, S. A. & BALSARA, D. S. 1994a M. N. R. A. S. Submitted. R. A., LAMB, S. A. &; BALSARA, D. S. 1994b. In preparation.

HERNQUIST, L. & WIEL, M. L. 1993

M. N. R. A. S. 261, 804.

HlGDON, J. L. 1993 PhD thesis, University of Texas at Austin. HUANG, S. & STEWART, P. 1988

A. & A. 197, 14.

JESKE, N. A. 1986 PhD thesis, University of California, Berkeley. KEEL, W. C , KENNICUTT, R. C , HUMMEL, E. & VAN DER HULST, J. M. 1985 A. J. 90, 708. KENNICUTT, R. C , KEEL, W. C , VAN DER HULST, J. M., HUMMEL, E. & ROETTIGER, K. A.

1987 A. J. 93, 1011. KING, I. R. 1966

A. J. 71, 64.

S. A., GERBER R. A. & BALSARA D. S. 1993 In The Evolution of Galaxies and Their Environment (Proc. of the Third Tetons Conference on Astrophysics). NASA Conf. Pub. 3190, p. 225. LAMB, S. A., GERBER R. A. & BALSARA D. S. 1994 Astr. & Space Sci. In press.

LAMB,

LONSDALE, C. J., PERSSON, S. E. & MATTHEWS, K. 1984 LYNDS, R. k TOOMRE, A. 1976

Ap. J. 287, 95.

Ap. J. 209, 382.

MARCUM, P. M., APPLETON, P. N. & HIGDON, J. L. 1992

Ap. J. 399, 57.

MOUSCHOVIAS, T. CH., SHU, F. H. & WOODWARD, P. R. 1974

A.& A. 33, 73.

SANDERS, D. B., SOIFER, B. T., ELIAS, J. H., MADORE, B. F., MATTHEWS, K., NEUGEBAUER,

G. & ScovnxE, N. Z. 1988 Ap. J. 325, 74. SCHULTZ, A.B., SPIGHT, L.D., RODRIGUE, M., COLEGROVE, P.T. & DISANTI, M.A. 1992

A. A. S. 23, 953. STRUCK-MARCELL, C. 1990

A. J. 99, 71.

STRUCK-MARCELL, C. & APPLETON, P. N. 1987 STRUCK-MARCELL, C. & HIGDON, J. L. 1993 STRUCK-MARCELL, C. & LOTAN, P. 1990

Ap. J. 323, 480.

Ap. J. 411, 108.

Ap. J. 358, 99.

THEYS, J. C. & SPIEGEL, E. A. 1976

Ap. J. 208, 650.

THEYS, J. C. & SPIEGEL, E. A. 1977

Ap. J. 212, 616.

THOMPSON, L. A. & THEYS, J. C. 1978

Ap. J. 224, 796.

B.

Violent Star Formation in Merger Remnants ByUTA FRITZE - VON ALVENSLEBEN Universitatssternwarte, Geismarlandstr. 11, D-37083 Gottingen, Germany In mergers of gas-rich spirals powerful starbursts are triggered, in the course of which a secondary population of globular clusters (GCs) may be formed. We present results from our chemical and spectrophotometric evolutionary models and show that even in the case of an old merger remnant like NGC 7252 the star formation (SF) history can be determined quite exactly, if only enough observational data are available. About a Gyr ago, NGC 7252 went through a starburst that, over (l — 5)108 yr, increased its stellar mass by 20 — 50% and created a number of new GCs detected with HST. Young GCs may serve as a tracer for star formation efficiency (SFE). Our models predict metallicities for a secondary population of GCs which should allow to identify Sp-Sp merger remnants among ellipticals. In the case of NGC 7252, follow-up spectroscopy of the two brightest young GCs confirmed our metallicity prediction. We show that once the metallicity is known, very exact age dating of these GCs becomes possible.

1. Introduction A most violent mode of SF is observed in mergers of massive gas-rich spirals. These starbursts have been considered as the near-by analogues of violent SF during the initial collapse at galaxy formation. Luminosities of 1012—1014 LQ are observed, predominantly emitted in the IR. All the IR-UL galaxies and many of the luminous IRAS galaxies have by now been shown to be in an advanced stage of merging. The physical processes governing this violent mode of SF seem to be fundamentally different from those of the normal SF mode observed in spiral, irregular, and non-interacting starburst galaxies, e.g. BCDGs. In the violent mode, SF efficiencies seem to be 1-2 orders of magnitude higher, and the same seems true for the fraction of very high density molecular gas (seen in HCN and CS lines) relative to the total interstellar gas content. All the range from galaxy-wide (as e.g. in NGC 3310) to purely nuclear starburst activity (e.g. in NGC 7714) is observed. The origin of these differences in starburst extension must, of course, be determined by the interplay of gas transport to the centre and gas consumption by SF, but it is by no means understood at present. A possible connection with non-thermal nuclear activity - feeding of a black hole, either preexisting or formed during the merger - might be suspected. By means of a chemical and spectrophotometric evolutionary synthesis model we describe the evolution of undisturbed galaxies of various spiral types, the starburst induced by a merger between two of them, and the future evolution of the merged post-burst system (Fritze - v. Alvensleben & Gerhard 1994a). Comparison with the wealth of observational data for the famous old merger remnant NGC 7252 allows us to seriously constrain the parameters of our model (Fritze - v. Alvensleben & Gerhard 1994b). Here, we focus on the SFE and on the metallicity of stars and GCs formed during the burst. In a second step, we calculate the photometric evolution of GCs using stellar evolutionary tracks and colour calibrations for various metallicities. Comparison with V - I colours for ~ 40 young GCs around NGC 7252 identified with HST shows that - by resolving the age-metallicity-degeneracy - quite accurate age-dating of these GCs becomes possible. 249

250

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants

2. Model and parameters Our 1-zone model can most briefly be characterised as a kind of a synthesis of a photometric and spectral evolutionary synthesis code in the spirit of Tinsley and Bruzual with a detailed chemical evolution model for a series of individual elements a la Matteucci. We start from a gas cloud of given mass and initial chemical composition. Our input data base comprises complete sets of stellar evolutionary tracks for 6 metallicities in the range 10~4 < Z < 2-ZQ with consistent colour calibrations in UBVRIK (Z) (see Einsel et al. 1994), high resolution stellar spectra observed from the UV to the NIR for solar metallicity, gaseous emission line fluxes for ~ 30 lines as a function of the number of Lyman continuum photons, lifetimes, yields and remnant masses for ~ 40 stellar masses (0.01M© < m+ < 85M 0 ). Basic parameters of our chemical and spectrophotometric model are TYlu

• the IMF, which we take from Scalo (1986) with a normalisation / <&(m)-m-dm = 0.5 mi

(Bahcall et al. 1992) that yields the correct M/L values, and • the SFRs as a function of time for which we use a linear Schmidt law with different characteristic timescales for SF, tt, for the different spiral types: tt = 2, 3, 10, 16 Gyr for 5a, Sb, Sc, Sd galaxies (cf. Sandage's 1986 empirical determinations). Ellipticals are described by an SFR decreasing exponentially with time and t* = 1 Gyr. The results of our models for the evolution of undisturbed galaxies have been extensively compared to spectrophotometric samples of nearby and high-redshift galaxies (Fritze - v. Alvensleben 1989; Kriiger et al. 1991, 1992, 1994), to chemical observations of the Milky Way and to QSO absorption line observations (Fritze - v. Alvensleben et al. 1989, 1991). A Merger in terms of our very simple model is described by • the undisturbed evolution of two spirals until some time ij, • a starburst of duration n - assumed to be triggered by the interaction at t\, - and • the subsequent fading of both the aging spiral and starburst populations. The IMF is assumed to be the same in the burst as for undisturbed evolution. This very simplified approach is subject to severe restrictions: we only consider mergers of two spirals identical in mass, Hubble type, and age, limit our description to global properties, and neglect any dynamical effects. The 4 parameters of our merger models, then, are the Hubble type and age of the progenitor spirals and the strength and duration of the starburst. We have presented an extensive study of our model parameters in terms of • 2-colour diagrams and colour tables from U .. .L, • luminosity evolution, and • chemical abundances and element ratios Though somehow more detailed and extended, this study is similar in spirit to the work of Schweizer & Seitzer (1992). It provided us with large grids for the classification and age dating of mergers and merger remnants.

3. NGC 7252 (=Arp 226) The "Atoms for Peace" galaxy NGC 7252 is the most extensively studied among old merger remnants. Despite its long tidal tails it has advanced far off towards becoming a "normal" elliptical galaxy. It shows a de Vaucouleurs profile over > 15 kpc, its fundamental plain parameters perfectly fit into the range of ellipticals. Among the available observational data are UBVRI luminosities and colours for various apertures (Schweizer

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants

251

1982ff; Hibbard et al. 1994), an optical spectrum (3600 - 5600 A, kindly provided to us by Schweizer) showing a nuclear Hp emission feature, information about the content of cold, warm and hot gas (e.g. Dupraz et al. 1990; Hibbard et al. 1994), and total mass estimates. Dynamical modelling has been done by Borne & Richstone (1991), and the first photometric modelling is presented in Schweizer & Seitzer (1992). The appearance of the tidal tails, their length, symmetry and gas richness already point to a pair of late spiral-type progenitor galaxies of similar mass. 3.1. Some results from photometric and spectral modelling Comparing the observed UBVR colours with our grid of models allows the determination of a cell in the 4-dimensional parameter space which is, however, not very small yet. Taking further into account the detailed information contained in the observed spectrum and comparing with our synthetic model spectra considerably narrows down the permitted parameter space. It becomes possible to trace back in detail the SF history of this merger remnant - despite its advanced age - and to predict a lot of observationally testable quantities. Here, we only briefly pick out some of the results presented in Fritze - v. Alvensleben & Gerhard (1994b): • Two fairly massive ~ 12 Gyr old spirals of type Sc with MB ~ —21.3, Mtot ~ 1 • lO n M0, each, and a gas content of ~ 40 % have been the progenitors of NGC 7252. • The burst age is 1.3+^ Gyr. • The burst duration (the least restrictable parameter) was (1 — 5) • 108 yr, similar to the lifetime of the IR-UL phase (Carico et al. 1990). • The burst strength 6, defined as the increase of the stellar mass during the burst relative to the mass of stars present before, must have been > 0.2, and, more probably, as high as 0.5. This corresponds to an SFE rj (mass of luminous stars formed in the burst devided by available gas mass) in the range 0.2 < r) < 0.45, again in agreement with SFEs estimated from typical parameters of IR-UL galaxies. This burst strength implies a peak SFR during the burst of 300 - 540 MQ yr ~* and a peak bolometric luminosity of the order of 1 • 1O12£0. About a Gyr ago and with a sufficient amount of dust at that time, NGC 7252 could well have looked like an IR-UL galaxy. 3.2. Evidence for violent SF in NGC 7252 The evidence for violent SF and the high SFE comes from two independent lines of reasoning. First, the strength of the Balmer lines and the broad-band colours tell us unambiguously that the burst must have ended > 1 Gyr ago. Still, however, the luminosity of the system is high, especially when compared to the luminosity range of late-type spirals in Virgo (Sandage et al. 1985). Even if two spirals from the bright end of the luminosity distribution have merged, a very strong burst (6 = 0.5) is needed to explain the high present-day luminosity of the system. If the progenitor spirals were not both from the bright end, an even higher SFE would have to be postulated. The high total mass of the system also supports the progenitors' luminosities. The empirical M/L agrees very well with that resulting from our models, and so does the presently observed gas content. No evidence for an important mass loss has been found with ROSAT (Hibbard et al. 1994), and no wind is expected to have occurred on the basis of a primitive static energy balance estimate (cf. Matteucci & Tornambe 1987) starting from our results. Second, the spectral fit we obtain for b = 0.5 is definitely better than that for 6 = 0.2, especially with respect to the Balmer lines. We conclude that an SF episode comparable in violence and duration to the one observed in situ in luminous IRAS galaxies should have occurred in the past of NGC 7252. Still now, 1.3 Gyr after the onset of the burst, several pieces of evidence (HST colour

252

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants

profiles, nuclear Balmer emission features) possibly point to ongoing SF at a rate of > 1.5 MQ yr" 1 within the central 1-2 kpc region, this SFR is included in our post-burst modelling. The late high luminosity of NGC 7252 also argues against a strong truncation of the IMF at low masses. Only burst stars with m* < 2 MQ are still alive and can account for the observed luminosity. 3.3. Chemical evolution of NGC 7252 Our study of the chemical evolution of NGC 7252 is restricted in the sense that we only consider a 1-phase ISM which we assume to be perfectly mixed at any time. The time delay due to the finite lifetimes of the stars restoring enriched material through winds, PN ejections, and SNe is, however, taken into account, i.e. we do not use the Instantaneous Recycling Approximation. While the global ISM metallicity, Z, as well as element abundances, [O/H] etc., increase strongly during the strong burst, stars formed in the burst should essentially reflect the metallicity of the spiral progenitor ISM. Abundance ratios [O/Fe], [C/O], . . . should be typical of spiral galaxies though slight enhancements of SN II products through selfenrichment during the burst cannot be excluded. Our models give burst star metallicities of > ifo ZQ for mergers of Sa through Sd galaxies happening not more than 5 Gyr ago. For NGC 7252, our best-fit model predicts Z > 0.008, [O/H]> -0.45, [Fe/H]> -0.73.

4. GC formation and evolution in mergers 4.1. GC formation in mergers The work of Burkert et al. (1993) has shown that for the formation of GCs as bound systems a high SFE rj > 0.2 is required. The results from our models that 0.2 < 77 < 0.5 led us to suspect that in the case of NGC 7252 a secondary population of GCs might have formed. Normalised to the stellar mass of the galaxies, the mean specific GC frequencies around elliptical and spiral galaxies are (TGC)sp ~ 2.2 and (TGC)E ^ 5, suggesting that, if elliptical galaxies are single-event mergers of two spirals, the number of GCs should increase by a factor ~ 2 during the merger-induced starburst (Zepf & Ashman 1993). If GCs are formed in such a major merger, their metallicities should be as high as that of the burst stars, i.e. > !/3 ZQ. The resulting bimodal metallicity distribution of the final GC system will still testify to a merger origin of the parent galaxy when other suspicious features, like blue colours, tails, etc., have long disappeared. Bimodal metallicity distributions of GCs have indeed been observed around several elliptical galaxies by Ostrov et al. (1993) and Zepf k Ashman (1993), and the mean ([Fe/H]) ~ -0.4 of the secondary peak agrees nicely with our predictions (cf. Sect. 4.3). Young populations of GCs around suspected merger remnants have been detected by direct high-resolution imaging around NGC 3597 by Lutz (1991) and around NGC 1275 by Holtzman et al. (1992). Whitmore et al. (1993) have recently reported the HST discovery of some 40 young GC candidates in NGC 7252; Schweizer & Seitzer (1993) present spectra of the two brightest of these. From these spectra, they derive Mgb values of 3.5 and 3.7, which correspond to Mg2 ~ 0.22 (Burstein et al. 1984) and to [Fe/H] 0.5 (Einsel 1992; see also Buzzoni et al. 1992), to be compared to our model prediction of [Fe/H] > —0.7 for NGC 7252.

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants

253

I.O-

Z = 1*10"(-4)

1.6-

z = 1»10"(-3) z = 1*10"(-2) z = 4»10"(-2)

1.41.21.01

0.80.60.40.20.0*•"

-0.27.0

8.0

9.0

10.0

11.0

Log Time [yr] FIGURE 1. Colour evolution in V — I of GCs of various metallicities. 4.2. Colour evolution for GCs of various metallicities

From the observed colours and line strengths, Schweizer & Seitzer (1993) determine ages for these GCs using Bruzual k, Chariot's (1993) single-burst models for solar metallicity. The resulting ages (10-500 Myr) are much lower than the burst age of 0.8-1.3 Gyr we derive from global properties or dynamical ages of the merger. That is why we calculate the colour evolution of GCs (t» = 106 yr, Scalo-IMF = 0.1.. .60 MQ) with stellar evolutionary tracks and colour calibrations for different metallicities. Here, we only briefly report some preliminary results, a more careful analysis and thorough discussion will be presented in Fritze - v. Alvensleben & Burkert (1994). As can be seen from Figurel, we find enormous age differences atfixedcolour between GCs of different metallicities. A stellar population showing e.g. V-I = 0.8 can have ages of3.410 8 , 1.1109, 5.0-109, and 8.9-109 yr for metallicities Z = A-10~2, 110" 2 , M 0 " 3 , and 1 • 10~4, respectively. For this V — I and a Scalo-IMF, Bruzual & Chariot's models yield an age of 5 • 108 yr for ZQ = 2 • 10~2, which correctly falls between our ages for 2ZQ and i/a-^oTaking from Whitmore et al. (1993) mean V-I colours for the outer (R > 2 kpc, 36 GCs) and inner (R <2 kpc, 9 objects) sample, and from Schweizer & Seitzer (1993) the V-I values for the two brightest GCs, W3 and W30, and comparing to our GC model for Z = 10~2, the metallicity closest to our model prediction of Z > 0.008 for the burst population of NGC 7252, we derive the GC ages presented in Table 1. The high ages of W3 and the 36 GCs in the outer sample are in good agreement with our global burst age of 1.3 ^

The mean age of the inner sample of 9 objects,

254

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants Object outer sample inner sample

V—I 0.83 ±0.25 0.71 ±0.35 0.84 ±0.05 0.53 ±0.05

Ages [Gyr]

1-371J1

0.5+1.37 ±0.3 0.3 ±0.1 1. Ages of young GCs in NGC 7252 as derived from our Z - 0.01 model for the W3 W30

TABLE

Distance from centre > 2 kpc < 2 kpc 7 kpc 5 kpc

observed (V — /) colours.

which - contrary to the outer sample clusters - have effective radii about twice as large as those of typical Galactic GCs, and show greatly varying V — I colours, is not well denned. Whitmore et al. already mention the possibility that the large Rejf objects from the inner sample may rather be young OB associations or 30 Dor-like regions rather than GCs. The well-determined low age of 0.3 ± 0.1 Gyr for W30, which has Rejj = 7.6 pc, agrees with the upper limit of Schweizer & Seitzer's estimate, and it is not understood at present if this is due to a lower than average metallicity or to a locally prolonged violent SF episode. 5. Conclusions We have shown that detailed spectrophotometric modelling together with a wealth of observational data allows the rather precise selection of one out of a large number of evolutionary paths that led to the present appearance of NGC 7252 and the statement that NGC 7252 must have had a very violent burst comparable to the ones observed today in luminous IRAS galaxies with a high star formation efficiency 0.2 < 77 < 0.5. A secondary population of GCs formed during this burst, removing the last objection against a spiral-spiral merger origin of - some - elliptical galaxies. The metallicity of these GCs is Z > XJZZQ. Any secondary GC population formed from Sp + Sp —> E mergers is expected to "eternally" show up as a second peak in the GC metallicity distribution and represents an indication of the height of the star formation efficiency. Calculation of the photometric evolution of GCs for their appropriate metallicity removes the age-metallicity degeneracy and allows very accurate age-dating. I thank A. Burkert, K. J. Fricke, and F. Schweizer for many helpful discussions, the organisers for this highly inspiring meeting, and many of the participants for stimulating discussions. Financial support from the Verbundforschung Astronomie through BMFT grant WE-010 R 900-40 is gratefully acknowledged.

REFERENCES BAHCALL, J. N., FLYNN, C. & GOULD, A., 1992 Ap. J. 389, 234. BORNE, K. D. & RICHSTONE, D. O. 1991 Ap.J. 369, 111. BRUZUAL A., G. & CHARLOT, S. 1993 Ap.J. 405, 538. BURKERT, A., BROWN, J. & TRURAN, J. W. 1993 In The

Globular Cluster Galaxy Connection, 11 th Santa Cruz Summer Workshop in Astronomy and Astrophysics (ed. G. H. Smith &; J. P. Brodie). In press.

BURSTEIN, D., FABER, S. M., GASKELL, C. M. & KRUMM, N. 1984 Ap. J. 287, 586. BUZZONI, A., GARIBOLDI, G. & MANTEGAZZA, L. 1992 A. J. 103, 1814.

Fritze - von Alvensleben: Violent Star Formation in Merger Remnants

255

CARICO, D. P., GRAHAM, J. R., MATTHEWS, K., WILSON, T. D. & SOIFER, B. T., NEUGEBAUER, G. & SANDERS, D. B. 1990 Ap. J. 349, L39. DUPRAZ, C , CASOLI, F., COMBES, F. & KAZES, I. 1990 A & A 228, L5.

ElNSEL, C. 1992 Diploma Thesis, University of Gottingen. EINSEL, C ,

FRITZE - v. ALVENSLEBEN, U., KRUGER, H. & FRICKE, K. J. 1994, A & A.

Submitted. FRITZE - v. ALVENSLEBEN, FRITZE - v. ALVENSLEBEN, FRITZE - v. ALVENSLEBEN, FRITZE - v. ALVENSLEBEN,

U. 1989 PhD Thesis, Gottingen. U. & BURKERT, A. 1994 In preparation. U. & GERHARD, O. E. 1994a A & A In press. U. & GERHARD, O. E. 1994b A& A In press.

FRITZE - v. ALVENSLEBEN, U., KRUGER, H., FRICKE, K. J., LOOSE, H.-H. 1989 A & A 224,

Ll. FRITZE - v. ALVENSLEBEN, U., KRUGER, H. & FRICKE, K. J. 1991 A & A 246, L59.

HlBBARD, J. E., GUHATHAKURTA, P., VAN GoRKOM, J. H. & SCHWEIZER, F. 1994 A.J. 107, 67. HOLTZMAN, J. A. et al. 1992 A. J. 103, 691. KRUGER, H., FRITZE - v. ALVENSLEBEN, U., LOOSE, H.-H. &; FRICKE, K. J. 1991 A & A 242,

343. H., FRITZE - v. ALVENSLEBEN, U., FRICKE, K. J. & LOOSE, H.-H., 1992 A & A 259, L73. KRUGER, H. & FRITZE - v. ALVENSLEBEN, U. 1994 A & A Submitted. KRUGER,

LUTZ, D. 1991 A & A 245, 31. MATTEUCCI, F.'& TORNAMBE, A. 1987 A & A 185, 51. OSTROV, P., GEISLER, D. & FORTE J. C. 1993 A. J. 105, 1762. SANDAGE, A. 1986 A & A 161, 89.

A., BINGGELI, B. & TAMMANN, G. A. 1985a, A. J. 90, 395. SCALO, J. M. 1986 Fund. Cosm. Phys. 11, 1.

SANDAGE,

SCHWEIZER, F. 1982 Ap. J. 252, 455. SCHWEIZER, F. k, SEITZER, P. 1992 A. J. 104, 1039. SCHWEIZER, F. & SEITZER, P. 1993 Ap.J. 417, L29. WHITMORE, B.C., SCHWEIZER, F., LEITHERER, C , BORNE, K. &; ROBERT, C. 1993 A. J. 106,

1354. ZEPF, S. E. & ASHMAN, K. M. 1993 M. N. R. A. S. 264, 611.

UV Variability of IRAS 13224-3809 ByJ. M. MAS-HESSE 1 , P. M. RODRIGUEZPASCUAL2,! L. SANZ FERNANDEZ DE CORDOBA1 AND TH. BOLLER3 1

Laboratorio de Astrofisica Espacial y Fisica Fundamental, POB 50727, E-28080 Madrid, Spain 2 3

ESA IUE Observatory, POB 50727, E-28080 Madrid, Spain

Max-Planck-Institut fur extraterrestrische Physik, D-85740 Garching, Germany

Cross-correlation of the ROSAT All Sky Survey and the IRAS Point Source Catalog has provided a sample of 244 galaxies with strong emission at both far-infrared and soft X-ray ranges. IRAS 13224-3809 appeared as an outstanding object within this sample due to its extreme X-ray luminosity (Lx = 3 • 1044 erg s" 1 ), steep X-ray spectrum and rapid X-ray variability, with a doubling timescale of only 800 s (Boiler et al. 1993). We have performed repeated IUE observations of this object in January, February and May 1993, looking for variable features in its spectrum, having detected a strong variability in the Lya line. While a relatively broad Lya component (FWHM ~ 5000 km s"1) remains essentially constant over the three IUE observations, the initially strong and narrow core emission component vanishes completely becoming a strong absorption. A maximum variation of 50% has also been detected in the UV continuum level. IRAS 13224-3809 has a deficit of UV emission when compared to Seyfert 1 galaxies. The UV-X-ray energy distribution suggests that the UV bump frequently found in these galaxies might be present at higher energies, well within the ROSAT band (0.1-2.4 keV). If this bump is due to thermal emission of a heated accretion disk, as proposed by several authors, its temperature should be significantly higher than in other similar objects (blackbody temperature kT ~ 100 eV). On the other hand, the strong far-infrared emission, the X-ray and Lya profile variability, the absence of broad Balmer-line components and the high Ha/H/? ratios could be explained by assuming the presence of a nucleus that has become recently "active", being still surrounded by large amounts of gas and dust which would obscure selectively different emitting regions. A recent merging episode could possibly be at the origin of this scenario, since it would provide large amounts of gas and dust and would trigger energetic episodes in the nuclei involved. IRAS 13224-3809 could therefore provide some insights into possible evolutionary links between Infrared Ultraluminous Galaxies and AGNs. A more detailed discussion of the observational properties of IRAS 13224-3809 will be presented elsewhere (Mas-Hesse et al. 1994).

REFERENCES BOLLER, T H . , TRUMPER, J., MOLENDI, S. FINK, H., SCHAAEIDT, S., CAULET, A. & DEN-

NEFELD, M. 1993 Astion. & Astrophys. 279, 53. MAS-HESSE,

J.M., RODRIGUEZ-PASCUAL,

P.M., SANZ FERNANDEZ DE CORDOBA,

BOLLER, TH. 1994 Astron. & Astrophys. In press. t Affiliated to the Astrophysics Division, Space Sciences Department 256

L. &

257

Mas-Hesse et al: UV Variability of IRAS 13224-3809

E on
X 3

1300

1400

1500

1600

1700

1800

1900

wavelength (A) FIGURE 1. WE spectra of IRAS 13224-3809 at observed wavelength frame and with no correction for galactic reddening. The vertical dotted lines mark the redshifted Lya and N V A1243 lines wavelength. 0.8

0.6

-15000

-10000

-5000 0 5000 Velocity (km/s)

10000

15000

FIGURE 2. (a) Lya line profiles with the velocity scale centered at (1 + z)1216 A; (b) January and February Lya spectra after subtracting the May one. Note the variations in the narrow-line core. The continuum emission as measured over 1400-1800 A was first removed from each spectrum in (a) and (b).

Star Formation in Polar-Ring Galaxies By V. RESHETNIKOV 1 2 AND F. COMBES 1 'DEMIRM, Observatoire de Paris-Meudon, F-92195 Meudon, France 2

Astronomical Institute of St.Petersburg University, 198904 St. Petersburg, Russia

We report the results of new spectral and photometric observations of a sample of northern polar-ring galaxies (PRGs). We found that the nuclei of galaxies with a forming polar ring have both higher Ha equivalent width and luminosity and show enhanced star formation. We also find an excess of active nuclei (Sy or LINER) among forming PRGs. From the optical colours and Ha luminosities of the rings, we infer that the rate and IMF of star formation in the polar rings are similar to those in normal late-type spiral galaxies.

1. Introduction Polar-ring galaxies (PRGs) are among the most interesting and spectacular relics of gravitational interactions between galaxies. The existence of prominent rings of gas, dust and stars aligned roughly in a perpendicular orientation with respect to the major axis of the main galaxy has provided a new opportunity to study the dynamical structure and evolution of galaxies. According to the most popular point of view, the observed structure of PRGs is the result of galaxy interactions ranging from a gas accretion to a complete merger. With the goal of a detailed study of PRGs we undertook observations of a restricted sample of northern PRGs from the catalogue by Whitmore et al. (1990) using the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences for the CCD photometry in the B, V, and Rc bands (this work is now in progress) and the 1.93-m telescope of the Observatoire de Haute-Provence (spectral observations of 16 PRGs in the range 6200-7170 A). In accordance with the optical appearance and internal kinematics we roughly classified the objects into two groups: "true" PRGs and currently forming PRGs due to accretion of matter or merging with a companion (see Figure 1). Here we will briefly discuss (in combination with other work in the literature) observational evidence of ongoing star formation in the central galaxies and rings of PRGs. (Throughout this work, all distance-dependent quantities are calculated using Ho = 75 km a"1 Mpc"1.) 2. Central galaxies 2.1. Nuclear emission-line spectra The nuclear spectra were classified into HII regions and low- (LINER) or high-ionization (Seyfert) spectra. We found that about half of the galaxies with forming rings due to current accretion or merging demonstrate evidence of active (Sy or LINER) nuclei. In spite of the small volume of our sample (16 PRGs), this conclusion is quite reliable since we have investigated about half of all northern PRGs and good or possible candidates in PRGs from Whitmore et al. (1990). An excess of active nuclei among PRGs become more evident with the inclusion of two well-investigated northern PRGs - NGC 660 and NGC 2685 - which also demonstrate LINER-type spectra (van Driel et al. 1993; Willner et al. 1985). Comparing the median properties of our sample [WOk(Ha) = 20 A, logi(Ha) = 39.0 258

Reshetnikov & Combes: Polar-Ring Galaxies

FIGURE

259

1. 5-band images of UGC 7576 ("true" PRG, left) and UGC 4261 (forming PRG, right) with arbitrary orientation and scale.

(in erg s 1), log SFR(Ra) = —7.2 (in M© yr x pc 2)] with the characteristics of various galaxy samples according to Keel (1983), Keel et al. (1985), and Bushouse (1986) we found that in general the nuclear properties of PRGs are close to those for strongly interacting galaxies with comparable absolute magnitudes: PRG nuclei have both higher Ha equivalent width and luminosity and also show more extensive star formation than field spirals (this conclusion becomes more significant if we consider the nuclei of spiral PRGs only). 2.2. Far-infrared (FIR) luminosities Excluding the ultraluminous infrared galaxy UGC 5101, the mean value of L(FIR) for our sample PRGs is 2 1010 L 0 . For the subsample of galaxies with forming rings by accretion (UGC 5600, NGC 6285, NGC 3808B), the corresponding value is 510 10 L Q . These values exceed significantly the mean FIR luminosities of isolated spirals and are close to those for strongly interacting galaxies (Bushouse 1987; Bushouse, Lamb & Werner 1988). Converting FIR luminosities into SFRs according to Hunter et al. (1986) we find that, excluding UGC 5101, the mean value of global SFR for our sample PRGs is 10.5 ± 5.5 MQ yr - 1 (for the subsample of forming rings - UGC 5600, NGC 6285, NGC 3808B - the corresponding value is 26 ± 13 MQ yr" 1 ). Excluding UGC 5101, the mean ratio of L(FIR) to L(B) is 12 ± 5 (the median ratio is 6.4), which is also close to that for interacting galaxies (Bushouse 1987, Bushouse et al. 1988). Therefore, since the L(FIR)/L(B) ratio measures the ratio of current to more long-term (several Gyr) SFRs, we find that star formation in PRGs (as in interacting galaxies) is characterized by a strong increase in recent time.

2.3. CO emission CO(1->0) emission has been detected in three PRGs: NGC 2685 (Taniguchi et al. 1990), NGC 660 (Combes et al. 1992), and ESO 603-G21 (Arnaboldi et al. 1993). In NGC 660 and ESO 603-G21 the presence of molecular gas in the central parts of the galaxies probably reflects ongoing star formation. In NGC 2685 the circumnuclear molecular gas may have only little relation to the star-forming activity (Taniguchi et al. 1990). In the rings of NGC 660 and ESO 603-G21 the total H2 mass is about 109 MQ. For the ring of NGC 660 the HI/H2 mass ratio is 3.6 (Combes et al. 1992), which is typical for late-type disk galaxies (Thronson et al. 1989).

260

Reshetnikov & Combes: Polar-Ring Galaxies

FIGURE 2. Spatial structure of Ha and [Nil] emission lines with the slit along the rings of UGC 4261, UGC 7576, and UGC 9796 (the brightest knots are marked by numbers). The distances between the two large dashes are 23" (vertical axis) and 20 A (10 A for UGC 7576 - horizontal axis).

3. Polar rings 3.1. Optical colours

The polar rings are bluer than the central galaxies in all well-studied cases. The mean B- V colour of the rings of 6 kinematically confirmed "classic" PRGs (A 0136-0801, ESO 415-G26, NGC 2685, UGC 7576, NGC 4650A, and UGC 9796) is +0.49 ± 0.12 (s.e.m.) The colours were corrected for the Galactic absorption only. Inclusion in the sample of the data for the rings of other PRGs or forming PRGs does not change the mean value = +0.49 ± 0.06 for 12 rings. This value, typical for the late-type spirals and also for the tidal tails of galaxies, indicates ongoing star formation in the rings. 3.2. Ha emission The spatial structure of the emission lines with the slit orientated along the polar rings of UGC 7576, UGC 9796 and along the forming ring of UGC 4261 is shown in the isophotal maps shown in Figure 2. As one can see, the distribution of Ha emission along the rings of UGC 7576 and UGC 9796 shows knotty structure. The Ha and [Nil] emissions in UGC 4261 show only one prominent condensation associated with the NE side of the forming ring. Taking the Ha luminosities of the condensations into account (0.3 — 3 1039 erg/s), their linear sizes ( < 2 kpc) and absolute luminosities (Mjj = —13™3 ± 0™8), we conclude that they are giant HII regions. Figure 3 plots the location of the mean characteristics of the detected HII regions (crosses) in the log W-^Ha + [Nil]) - B - V plane, where W\(Ra + [Nil]) is the equivalent widths of the joint Ha + [Nil] emission and B — V is the corrected for the Galactic absorption colour. The solid lines in Figure 3 restrict the location of the characteristics of a model disk galaxy for three different initial mass functions (IMFs) according to Kennicutt (1983). (The upper bounds represent the absorption-free models, and the lower bounds correspond to a depression of a factor 2 in the equivalent width due to dust absorption.) As one can see in this figure, the standard (Salpeter) IMF seems to reproduce the observed star-forming region properties satisfactorily, within the uncertainties introduced by observational errors and internal absorption. The total observable Ha luminosities of the rings of UGC 7576 and UGC 9796 obtained by integrating emission along the slit are 1040 erg s" 1 and 5 1039 erg s" 1 , respectively. Therefore, global SFRs in these rings (1.8 10~9 and 1.2 10~9 M© yr" 1 pc~2) correspond to the normal values for isolated spirals (Bushouse 1987). The global SFR in the ring of UGC 4385 (4.4 10~9 M 0 yr" 1 pc~2) is close to the mean value for the sample of interacting

Reshetnikov & Combes: Polar-Ring Galaxies

261

FIGURE 3. Comparison of the model Ha+[NII] equivalent widths and B — V colours with the mean characteristics of HII regions in the polar rings (crosses). The three regions correspond to the different IMFs: the Miller-Scalo function (1), Salpeter function (2), and the shallow m~2 IMF (3)

41 40

36

14 15 16 17 18 19 20 21 22 MB FIGURE

4. (1) - UGC 4261, (2) - UGC 7576, (3) - UGC 9796.

spirals (Bushouse 1987). Taking into account that the polar rings under consideration are nearly edge-on structures, the obtained estimates of SFRs/area may represent only lower limits due to internal extinction. Figure 4 shows the dependence of the mean luminosity of the three brightest HII regions on the absolute magnitudes and Hubble types of their parent galaxies (Kennicutt 1988; the solid line is a linear fit for the field Sc galaxies). The circles represent the data for the PRG HII regions (open circles show the total absolute magnitudes of the galaxies, and solid circles the absolute magnitudes for polar ring components only). We find that the observed HII regions of the PRGs are located in the same area as HII regions for a normal late-type spiral galaxy. It should be noted also that the general characteristics of the rings of UGC 7576 and UGC 9796 (total Ha+[NII] luminosity and absolute blue luminosity) are located in the corresponding plane exactly along the mean dependence for Sc-Irr galaxies (Kennicutt et al. 1987).

262

Reshetnikov & Combes: Polar-Ring Galaxies

4. Discussion We have found that the forming PRGs demonstrate enhanced nuclear star formation and an excess of active nuclei with respect to normal galaxies. The majority of our sample objects are the result of strong interaction with a companion galaxy. The presence of two kinematically distinct subsystems of gas gives evidence for gas transfer or complete merging between galaxies. The mass transfer from one galaxy to another is probably favourable for nuclear starbursts and active nucleus formation. The totality of observational data suggests active star formation in polar rings. We have found that in the first-order approximation the IMF and the global SFR in the rings are similar to normal late-type spiral galaxies. Different mechanisms can probably contribute to the observational properties of star formation in the polar rings. For example, if ring material originates in a gas-rich star-forming galaxy, then the star formation in some forming rings can be decaying with respect to the initial rate of star formation. Besides, in the course of simulations of settling gaseous disks Christodoulou & Tohline (1993) have noticed that, as a settling disk approaches its preferred planar orientation, the gas in the disk gets compressed against the preferred plane. Star formation is favoured in the disk or ring at this stage of evolution. Moreover, in a radially extended ring spiral instabilities may exist. These instabilities can create transient spiral patterns with enhanced gas density and, probably, with enhanced star formation. Our photometric observations of the ring of UGC 7576 give interesting evidence in support of the latter mechanism. We found local reddening maxima associated with the inner part of the ring. If these maxima are due to the presence of dust, it means that dust in the ring of UGC 7576 concentrates preferably towards the inner part of the ring. Such a dust distribution can be explained by the spiral structure of the ring.

REFERENCES ARNABOLDI, M., CAPACCIOLI, M. &; COMBES, F. 1993 In AGN conference. Canberra. BUSHOUSE, H. A. 1986 A. J. 91, 255.

BUSHOUSE, H. A. 1987 Ap. J. 320, 49. BUSHOUSE, H. A., LAMB, S. A. & WERNER M. W. 1988 Ap. J. 335, 74. CHRISTODOULOU, D. M. &; TOHLINE, J. E. 1993 Ap. J. 403, 110. COMBES F., BRAINE J., CASOLI F., GERIN, M. & VAN DRIEL, W. 1992 A. & A. 259, L65. HUNTER, D. A., GILLETT, F. C , GALLAGHER, J. S., RICE, W. L. & Low, F. J. 1986 Ap. J.

303, 1986. KEEL, W. C. 1983 Ap. J. Suppl 52, 229. KEEL, W. C , KENNICUTT, R. C , HUMMEL, E. & VAN DER HULST, J. M. 1985 A. J. 90, 708. KENNICUTT, R. C. 1983 Ap. J. 272, 54. KENNICUTT, R. C. 1988 Ap. J. 334, 144. KENNICUTT, R. C , KEEL, W. C , VAN DER HULST, J. M., HUMMEL, E. &; ROETTIGER, K.

A. 1987 A. J. 93, 1011. TANIGUCHI, Y., SOFUE Y., WAKAMATSU K.-I. & NAKAI N. 1990 A. J. 100, 1086. THRONSON, H.A. ET AL. 1989 Ap. J. 344, 747. VAN DRIEL,

W.

ET AL.

1993 A. & A. Submitted.

WHITMORE, B. C , LUCAS, R. A., MCELROY, D. B., STEIMAN-CAMERON, T. Y., SACKETT, P. D. & OLLING, R. P. 1990 A. J. 100, 1489. WILLNER, S. P., ELVIS, M., FABBIANO, G., LAWRENCE, A. & WARD, M. J. 1985 Ap. J. 299,

443.

Infrared Spectroscopy of IR-Luminous Galaxies By A. STERNBERG 1 , M. BLIETZ 2 , M. CAMERON 2 , R. GENZEL 2 , A. KRABBE 2 AND L. J. TACCONI 2 School of Physics and Astronomy, Tel Aviv University, Ramat Aviv, 69978, Israel 2

Max-Planck-Institut fur Extraterrestrische Physik, D-85740 Garching bei Munchen, Federal Republic of Germany

Near-infrared spectral imaging observations of the starburst galaxy NGC 1808 and of the Seyfert galaxy NGC 1068 are briefly discussed.

1. Introduction Most of the presentations at this meeting have focussed on optical, ultraviolet, and X-ray observations of starburst galaxies (SBGs) and active galactic nuclei (AGN) and their interpretation. In this contribution I draw attention to the utility of infrared array spectroscopy and millimeter-wave interferometry to the study of energetic galaxy nuclei. Infrared spectral observations are useful because they probe objects with large internal or foreground extinctions. Many interstellar sources such as photon-dominated regions in molecular clouds or non-dissociative shocks release energy at primarily infrared wavelengths. Millimeter spectroscopy provides information about the molecular medium which is not readily observable at optical, UV or X-ray wavelengths. In this article I discuss infrared observations of the starburst galaxy NGC 1808 and of the Seyfert galaxy NGC 1068 carried out by members of the MPE group (Blietz et al. 1994; Krabbe, Sternberg & Genzel 1994; Taconni et al. 1994). Most of this work was carried out using the MPE infrared array spectrometer FAST (Krabbe et al. 1993). 2. N G C 1808 NGC 1808 is a nearby (10.9 Mpc, for Ho = 75 km s" 1 Mpc"1) morphologically peculiar spiral galaxy (Sersic and Pastoriza 1965). Optical images show that several dust filaments protrude from the nucleus out into the galactic halo. Optical spectroscopic measurements have shown that intense emission lines lines (Ha, H/?, [SII] AA6717, 6731 A, [Nil] AA6548, 6583 A and [Oil] A3702 A) are produced within the central few kpc of the galaxy (Phillips 1993). NGC 1808 is a moderate IRAS source with a total far-infrared luminosity of LIRAS = 1.8 X 1010 LQ. NGC 1808 is a prominent radio source, and 6-cm radio VLA observations reveal a family of compact radio sources within the central ~ 750 pc "circumnuclear region". The radio knots are not well correlated spatially with the optical emission line (e.g. Ha) hot spots. We have produced l"8-resolution images of 2.17-//m HI Br7 and 2.12-/im, and 2.15-/im continuum emission, of the "circumnuclear region" of NGC 1808 (see Figures 1 and 2). Our images show that the Br? emission is extended and is produced in several distinct emission complexes. The 2.15-/im continuum emission is distributed smoothly through the Br7 emission region and peaks sharply at the nucleus. The Br7 emission knots are well aligned and almost coincident with several compact 5 GHz radio sources in the circumnuclear region. From a comparison of our Br7 and 263

Sternberg et al:

264

15

1

' ' i ' ' ' ' rr

i

i i i i i

i i i i

i I i

NGC 1808 Bry 2.17 urn 10

o V

- 1 0

- 1 5 - 1 5

I M M I

-1O

-5

I

I I I I I I I I I I I I I I I I !

0

5

1O

IS

R.A. Offset [arosec]

FIGURE 1.

a Ha image (see Krabbe et al. 1994). we conclude that the Br7 and radio emission sources trace the actual locations of HII regions and supernova remnants formed in distinct star-forming clusters, while the optical spots primarily trace directions of low foreground extinction. Our Br7 line data show that between 10% and 50% of the 5-GHz radio emission is thermal free-free emission produced in dense HII regions. We have analyzed the individual star forming-clusters within the circumnuclear region using models of evolving star-clusters (see Krabbe et al. 1994). Our analysis shows that the Br7, 2.15-//m continuum, radio and far-infrared IRAS data are consistent with starformation activity which has been proceeding steadily for between 5xlO 7 and 108 yr in the nucleus, and at most ~ 107 yr in the luminous circumnuclear star-forming clusters. The data are also consistent with "normal" stellar initial mass functions (IMFs) which vary as M ~ 2 5 where M is the stellar mass. The observations imply that the IMFs are likely truncated at upper mass limits of ~ 30 MQ. The IMFs may extend to small (~ 0.1 MQ) lower-mass limits, and there is no evidence for "biased" massive star formation in NGC 1808. For an assumed lower-mass IMF limit equal to 1 M© the inferred starformation rates range from ~ 0.5 to ~ 1 M© yr" 1 in the star-forming clusters. Much of the extended 2.15-^m continuum emission is probably produced by old late-type stars in the preexisting galactic bulge. However, it appears that the circumnuclear star-forming clusters are just becoming visible in the 2.15-^m continuum. Most of the intense nuclear

Sternberg et al.:

15

i

• \

i

i

i

i

r

265

•-r-r-|-

NGC 1808 2.15 urn

10

o •ss

o

o d

a - 5

c -

I

I

I

I

-10

1 I

I

I

I

I

1 __}

I

I

L I 1

-5 0 5 R.A. Offset [arcsec]

10

15

FIGURE 2.

2.15 /xm-continuum emission is probably produced by the large numbers of red giants and supergiants present in the evolved star-forming cluster in the nucleus. Our analysis suggests that in NGC 1808 the ratio between the supernova rate and synchrotron radio luminosities of the supernova remnants is equal to ~ 2.5 x 104 yr" 1 LQ1 HZ. Our analysis also shows that about half of the IRAS luminosity of NGC 1808 is probably produced in the nuclear star-forming cluster, and that the remainder is produced in the circumnuclear clusters. Our near-infrared observations do not indicate the presence of an active galactic nucleus (AGN) in the center of NGC 1808.

3. NGC 1068 NGC 1068 is a very prominent and nearby (14 Mpc) Seyfert 2 galaxy, and is one of the best studied systems containing an active galactic nucleus. The central luminosity of NGC 1068 is dominated by a compact 10 ^m source (Cameron et al. 1993) which emits 1.5 x 1011 LQ. The galaxy nucleus contains a "narrow-line region" (NLR) about 4" in diameter, and a star-formation ring with a radius of about 15". A bipolar 13" synchrotron radio jet is centered on the nucleus (Wilson and Ulvestad 1987). We used FAST at the William Herschel Telescope to produce high spatial resolution

Sternberg et al.:

266

NGC 1068

1"

0

-1" 2"

1"

- 0

T -6

100 pc I

S G H z

RADIO CONTINUUM

0

-2"

-i,"

-6"

R.A. - OFFSET FIGURE 3.

(0"8) infrared images of 2.12-^m H2 1-0 S(l), 1.64-/un [Fell] 4Z>7/2 - 4 ^9/2, and 2.17-/zm Br7 line emission from the NLR of NGC 1068 (Blietz et al. 1994). Our observations show that intense infrared iron line emission is produced in an unresolved source at the nucleus and in a "flare-region" associated with one of the lobes of the radio jet. The [FeII]/Br7 intensity ratio is very large in both the nucleus and flare and suggests that the gas phase abundance of iron is enhanced in the NLR of NGC 1068 compared with Milky Way values. The nuclear iron emission may be produced in dense clouds photoionized by X-rays from the central source, and the iron emission in the flare region may be produced in fast J-shocks in the jet outflow (Blietz et al. 1994). Our H2 image superimposed on an image of the radio jet is shown in Figure 3. Most of the vibrational H2 emission is produced within or in the immediate vicinity of the optical NLR. This is consistent with millimeter interferometric observations of CO and HCN J = 1 —• 0 rotational line emission which show a clear concenration of molecular gas within the central 3" (Planesas, Scoville and Myers 1991; Tacconi et al. 1994). The CO observations imply a total molecular mass of about 3 x 107 M 0 within the central 200 pc of NGC 1068, and suggest that most of the gas in the NLR is actually molecular. Our H2 image consists of several emission knots, including one located close to the near-infrared

Sternberg ti al:

267

continuum peak at the galactic nucleus. The molecular cloud traced by the central H2 emission knot likely contributes significantly to the obscuration of the embedded broadline region revealed by the optical scattering-polarization observations of Antonucci and Miller (1985). The central H2 emission knot may be produced in dense gas heated by UV and X-ray radiation from the embedded AGN (Rotaciuc ti al. (1991). The more extended H2 emission may be produced in shocked gas at the leading edge of a stellar bar (Tacconi ti al. 1994). The molecular clouds in the NLR of NGC 1068 appear to be very unusual. The HCN/CO line brightness-temperature ratio is very large and exceeds unity at the nucleus. The nuclear HCN/CO intensity ratio is larger than seen anywhere in our Galaxy (except in much smaller dense cores of active star-forming clouds), and is much larger than the ratio observed in other galaxies including the most powerful IRAS galaxies. The large HCN/CO intensity ratio suggests that most of the molecular gas is contained in clouds with densities exceeding 105 cm" 3 , but also that the HCN/CO abundance ratio is unusually large (10~3 - 10"1) in these clouds (Tacconi ti al. 1994). The large HCN/CO abundance ratio may be indicative of a severe depletion of gas phase oxygen in the NLR. An oxygen depletion would result in an increased HCN abundance and reduced CO abundance in the molecular clouds compared with typical Milky Way values (Sternberg ti al. 1994). Evidence for an oxygen depletion in the NLR of NGC 1068 has recently been presented by Marshall ti al. (1993) in an X-ray study of the ionized component.

REFERENCES ANTONUCCI, R. R. J. & MILLER, J. S. 1985 Ap. J. 297, 621. BLIETZ, M., CAMERON, M., DRAPATZ, S., GENZEL, R., KRABBE, A., VAN DER WERF, P. V.

D., STERNBERG, A. & WARD, M. 1994 Ap.J. In press. CAMERON, M., STOREY, J. W. V., ROTACIUC, V., GENZEL, R., VERSTRAETE, L., DRAPATZ, S., SIEBENMORGEN, R. & LEE, T. 1993 Ap. J. 419, 136. KRABBE, KRABBE,

A. et al. 1993 Publ. Ast. Soc. Pac. In press. A., STERNBERG, A. &; GENZEL, R. 1994 Ap. J. In press.

MARSHALL et al. 1993 Ap. J. 405, PHILLIPS, A. C. 1993 A. J. 105,

168. 486.

PLANESAS, P., SCOVILLE, N. & MYERS, S. T. 1991 Ap. J. 369,

364.

ROTACRIC, V., KRABBE, A., CAMERON, M., DRAPATZ, S., GENZEL, R., STERNBERG, A. &;

STOREY, J. W. V. 1991 Ap. J. Lett. 370, L23.

SERSIC, J. L., & PASTORIZA, M. 1965 Publ. Ast. Soc. Pac. 77, 287. STERNBERG, A. et al. 1994 In preparation. TACCONI, L. J., GENZEL, R., BLIETZ, M., CAMERON, M., HARRIS, A. I. &; MADDEN, S. 1994

Ap. J. Lett. In press.

Application of the Multiphase Model to the Galactic Bulge By M. MOLLA 1 :

A. I. DIAZ 1 AND F. F E R R I N I 2

Dep. de Fisica Teorica, Universidad Autonoma de Madrid, 28049-Cantoblanco, Spain

2

Dipartimento di Astronomia, Universita di Pisa, Piazza Torricelli 2, 56100-Pisa, Italy

We present the application of the multiphase model to the Galactic bulge to assure that this model may be used in all regions. Results show a star formation rate that is higher in the bulge. The logical consequence is a higher metallicity due to the burst of star formation at early phases of the evolution. But when a comparison is made between models applied to an elliptical galaxy and a bulge, a different chemical evolution results: the relation [O/Fe] versus [Fe/H] is not the same for elliptical galaxies and bulges.

1. The multiphase model The bulge is the central region of the Galaxy (Frogel, 1988) with a radius of 1-2 kpc and a total mass of ~10 10 M©. In the outside region, the bulge connects with the spheroidal halo and with the adjacent disc. Characteristics for the stellar population of this bulge are obtained from observations: stars are old (Terndrup 1988; Lee 1992), metal rich (Gratton & Ortolani 1986; Rich 1988) and spatially distributed in two components (Rich, 1990; IRAS results).There is also a spheroidal metal-poor component without rotation and a more centrally concentrated and more metal-rich component with rotation. The observed metallicity distribution is well reproduced by a simple "closed box" model. It implies no slow "infall" of gas in this region. Other models for the bulge are those of Arimoto & Yoshii (1987), based on their elliptical galaxy models, and Matteucci & Brocato (1990). We use the multiphase model, which has been used for the Galactic disc (Ferrini et at. 1994; hereafter FMPD) with good results. This same model has been used for the bulge (Molla k Ferrini, 1994; hereafter MF), with the possibility of Supernova (SN) winds; that is, there is a mass loss from the bulge, a fraction of which flows to the adjacent disc. A SN wind appears when the thermal energy produced by SN explosions is larger than the binding energy. It is assumed that only a fraction, ICB, of the affected mass is lost. The mass of an external shell of the bulge may be lost, but the internal mass is in a medium of larger density which is only heated by SN explosions. There are four zones in the model: bulge, adjacent disc, halo on the bulge and halo on the disc. There is a mass transport from each region to the others: the halo (on the disc) gas collapses and forms out the disc (infall, /#); the halo (on the bulge) gas collapses and forms out the bulge (infall, /#); and the bulge is connected with the adjacent disc if there is a SN wind which transports the mass lost from the bulge. In this case a fraction of it goes through the halo to the intergalactic medium and another fraction goes to the disc region. The other characteristics of the model are described in FMPD and MF. Studied regions are modelled as two cylindrical regions surrounding a spherical bulge. Parameter values (fi and the other parameters related on the star formation processes) on every region i are calculated by the model taking into account the radial dependence obtained from the disc study through the volume of each region. 268

Molla et al.: Application of the Multiphase Model to the Galactic Bulge

269

0.5

-1

+model 1 • model 2 x model 3 omodel 4 Amodel 5 • model 6 -2

-1

[Fe/H] FIGURE 1. The relation [O/Fe] vs. [Fe/H] for 6 different models.

2. Results A set of 6 models has been run with different bulge masses, MB , and different massloss fractions, kB, with Mo = 13(109 MQ) between R = 2 kpc and R = 4 kpc, ko = 0.0, and ks = 0.50 when kB ^ 0.0. Models are: 1) MB = 10, 2) 20, and 3) 40 10 9 M o with kB = 0.0, and 4) /fcB=0.25, 5) 0.50 with MB = 20 10 9 M o . Model (6) has MB = 200 109M© andfcB = 5.0. Comparing a "bulge" model (number 2) with a simulating elliptical galaxy model (number 6), we see an important difference in the relation [O/Fe] vs. [Fe/H]: the "plateau" is larger for the model of the elliptical, and the overabundance [O/Fe] decreases when the metallicity reaches a larger value, [Fe/H] ~ —0.8 vs. [Fe/H] ~ —1.2, the value for the bulge. We computed other models without mass loss and different MB (1, 2 and 3) and also models with the same mass and different mass loss fraction values (2, 4 and 5) to compare. Results (Figure 1) show: a) different kB value models show the same [O/Fe] vs. [Fe/H] behaviour, the effect of mass loss on the relation [O/Fe] vs.[Fe/H] being very small; and b) the value [Fe/H] at which [O/Fe] decreases, increases with increasing mass. Therefore, the relation [O/Fe] vs. [Fe/H] is not the same for the bulge and for elliptical galaxies: the "plateau" of the [O/Fe] overabundance is longer for elliptical. The [O/Fe] decreases at a lower metallicity [Fe/H] for the bulge. This result is different from Matteucci & Brocato's (1990) prediction, who suggest the same relation for both cases. This different behaviour seems to be due to larger mass and density in elliptical galaxies.

REFERENCES ARIMOTO, N. &C YOSHII, Y. 1987 Astr. Astrophys. 173, 23. FERRINI, F., MOLLA, M., PARDI, M. C. & DIAZ, A. I. 1994 Astrophys. J. In press. (FMPD) FROGEL, J. A. 1988 Ann. Rev. Astr. Astrophys. 26, 51. GRATTON, R. G. & ORTOLANI, S. 1986 Astr. Astrophys. 169, 201. LEE, Y.-W. 1992 Astron. J. 104, 1780. MATTEUCCI, F. & BROCATO, E. 1990 Astrophys. J. 365, 539. MOLLA, M. & FERRINI, F. 1994 In preparation. (MP) RICH, R. M. 1988 Astron. J. 95, 828. RICH, R. M. 1990 Astropiiys. J. 362, 604. TERNDRUP,

A.M. 1988 Astron. J. 96, 884.

Stellar Populations and Population Gradients in Spiral Bulges By MARC BALCELLS AND REYNIER F. PELETIER Kapteyn Astronomical Institute, Postbus 800, 9700 AV Groningen, Netherlands Galactic bulges probably formed in starbursts such as those studied during this conference. We study the population contents and spatial structure of bulges to learn on the star formation history of these systems. From broadband optical CCD images we derive mean colours and colour gradients for a complete sample of ~ 40 bulges of edge-on galaxies. After excluding objects with dust, the colours trace age and metallicity of the stellar populations. Metallicities inferred from the colours seldom reach solar values. This result is consistent with recent metallicity estimates for the Milky Way (MW) Bulge derived from spectroscopy of K giants. Colour profiles have negative slopes, i.e. colours become bluer outward. The derived gradients are similar to those observed in elliptical galaxies. Gradients of more luminous bulges are steeper than those of smaller bulges.

1. Introduction Our notions on the stellar content of spiral bulges are largely derived from studies of a few nearby cases, notably our own MW Bulge. We know that the spectra of bulges often resemble more the strong-lined spectra of elliptical nuclei than those of metal-poor halo stars (Whitford 1978). It is also known that the MW Bulge, at Baade's window (BW) at least, contains stars of very high metallicity (Rich 1988). Often disregarded is the large width of the distribution of BW K giant metallicities, and the fact that the mean metallicity is just about solar. Systematic studies of external bulges have focussed on early types; Visvanathan k Sand age (1977) find that the integrated colours of SOs are the same as those of ellipticals of the same luminosity. For types later than SO, population diagnostics based on colours are difficult to obtain. Reddening due to dust in the disk seen in projection makes colours (and even absorption line strengths) difficult to interpret. Dust makes the study of population gradients in bulges particularly hard. The early result of Wirth k. Shaw (1983), which showed that bulges of late-type spirals have stronger colour gradients than ellipticals and SOs, is likely to be strongly affected by dust, as the galaxies they observed are all of high inclination. Thus, little or no data exist to answer even simple questions such as: what are the colours of the stellar populations of bulges? do bulges show colour gradients? what are the bulges' ages and metallicities? do these vary with bulge luminosity? with galaxy type? are the bulges of late-type galaxies bluer than those of early-types? do bulge colours correlate with disk colours? how do bulge colours compare with disk colours? are colours and gradients similar to those of ellipticals? is there evidence from the colours that bulges formed before or after the disk? The general availability of CCDs with reliable stability and flatness makes it possible partly to circumvent the problem of dust, and quantify the reddening, provided one works on a large, well-defined sample. We have undertaken a study of colours and colour gradients in bulges of early- to middle-type spirals, aiming at uncovering colours and colour gradients that give meaningful information about stellar populations. A comprehensive report of this work is given in Balcells k Peletier (1994); a related study on the colours of the disk components is given in Peletier k Balcells (1994). 270

Balcells &; Peletier: Colours of Bulges

271

2. Sample and observations While a face-on orientation is convenient for studies of the nuclear regions of bulges, the measurement of extended surface-brightness profiles of bulges requires high-inclination orientations to ensure that the faint outer parts of the bulges project away from the disk. We therefore selected all high-inclination (i > 50°) SO-Sbc galaxies from the UGC with diameter larger than 2 arcmin and blue magnitude brighter than 14.0, within a given RA-dec range. We excluded barred galaxies as their analysis is more complicated. Observations were carried out with the INT PF camera equipped with a coated GEC CCD (scale 0.549 arcsec pixel"1). We took UBRI images of about 40 objects. We had photometric weather throughout the run, which allowed us to secure a very homogeneous data set. Uniformly-illuminated fields were flattened to better than 0.2%. The effective seeing was in the range 1.0-1.5 arcsec FWHM. Surface-brightness profiles were derived over wedge-shaped apertures opening toward the galaxies' two semi-minor axes. Care was taken to register the UBRI images of each object using stars in the frames; wedges for the images were then centred on a common origin, which we took to be the luminosity peak of the /2-band frame. This was important as the dust in the highly-inclined disks causes the luminosity peak in each band to shift slightly in relation to the peak in the other bands. We ran a bulge-disk decomposition program to determine the relative contribution of the disk to the minor-axis light. We did not subtract the disk light, but measured colour profiles only out to a point where 50% of the light comes from the disk. For more details on the derivation of the profiles, refer to Balcells & Peletier (1994).

3. Quantifying dust reddening Central to our strategy is to determine when a measured colour is a reliable measure of the colour of the stellar population, and when the colours are instead largely affected by dust reddening. We did not attempt to deredden the colours, but rather to identify a subsample of bulges with colour profiles not affected by dust. Visual inspection allows us to identify patchy dust, but not distributed dust. But this is highly subjective. To quantify the presence of dust in an objective manner we used the ratio of disk scaleheights in the (/ and R bands, since we can expect dust to increase the disk scale-length in the bluer bands relative to the redder bands. Also, bulges considerably redder than the reddest elliptical galaxies were deemed dusty. Out of the 40 objects observed, 12 were too irregular for a disk-bulge decomposition, and another 12 did not fulfil the dust acceptance criterion, leaving 18 objects with colour profiles suitable for population studies. Of these latter we used the sides that had smooth radial colour profiles, always the side away from the dust lane, and sometimes the average of both sides.

4. Bulge colours Colours of bulges are somewhat bluer than those of ellipticals. This is true when the comparison is carried out at a given local surface brightness, as well as when mean colours are compared for a given integrated absolute magnitude. In Figure 1 we plot the colourmagnitude relation for the bulges in our study. For comparison, we plot the relation for an elliptical sample from Peletier et al. (1990, hereafter PDIDC), and the mean colourmagnitude relation for "old" ellipticals from Schweizer k Seitzer (1992). Bulges fall on or below the colour-magnitude relation for ellipticals. The distribution of deviations from the mean elliptical colours (for a given absolute magnitude) is shown in Figure 2.

Balcells & Peletier: Colours of Bulges

272 2.6

2

-

1.5

-

OS I

-16

-16

-2O

-22

-24

-18

-18

M R Bulge

-2O

-22

-84

M R Bulge

FIGURE 1. B — R colours for the 18 bulges without dust (open symbols). For each galaxy we plot the colour at 0.5 re, or at 5 arcsec if 0.5 r e is smaller than this. The number of sides in the symbols code the galaxy type, early- to late-type, from SO (circles) to Sbc (triangles). Filled symbols correspond to the ellipticals in PD1DC. The solid line is the colour-magnitude relation for old ellipticals.

a

r

-' —U-R . — B-R

i

:

1

1

1

4

-

1 1

•

2 •

r

1 —

1 1

1 '

m

O -O.8

-O.4

-0.2

O.2

FIGURE 2. Histogram of displacement of bulge colours with respect to the mean colourmagnitude relation for ellipticals. Shown are U — R (solid line) and B — R (dashed line).

From this figure we measure the colour differential of bulges with respect to ellipticals: A(B - R) - -0.08 ± 0.08 and A((/ - R) = -0.15 ± 0.19. Although it is tempting to ascribe the bluer colours to the contribution of bluer disk light, this is seen not to be the case, as bluer colours are found both at high and low bulge surface brightness. Metallicities derived from the colours, using single-burst models with ages ranging from 12 to 20 Gyr, are clearly subsolar. This is contrary to the belief that bulges have high metallicities. Careful examination shows that there is no contradiction; the Baade's Window region, whose super-metal rich K giants did much to establish the notion of the high metal content of bulges, has a mean rnetallicity of almost exactly solar, as derived

Balcells k. Peletier: Colours of Bulges

273

I

-16

-18

-20

-22

MRBulge FIGURE

-24-16

-18

-20

-22 -24

MRBulge

3. Colour gradients as a function of absolute magnitude for the bulges without dust (symbols as in previous figures) and for ellipticals from PDIDC.

from the same K giants (McWilliam k. Rich 1994). It is inaccurate to believe that the MW bulge as a whole is metal rich.

5. Colour gradients For the bulges without dust, the colour profiles are usually linear when plotted against logr. We measured gradients by fitting straight lines to the colour profiles out to the point where the disk contributes 50% of the light. We show A(fl — R)/A log r against the bulge absolute magnitude in Figure 3. Gradients for a sample of ellipticals are plotted for comparison. We find that the colour gradients in bulges show two distinct behaviours as a function of luminosity. For M^u'se < —20.0, gradients become increasingly negative with luminosity. These gradients are comparable to those of ellipticals, and so is the trend with luminosity (see PDIDC). The steepening of the colour profiles with luminosity is a specific prediction of dissipational models of galaxy formation, see e.g. Carlberg (1984). Dissipation may be the origin of the gradient-luminosity relation we have found. Fainter bulges {MRU 9" > —20.0) deviate from this relation. Gradients are generally much more pronounced and the scatter is large. While these measurements are difficult due to the small sizes of these bulges, errors in either the gradients or the absolute magnitudes cannot explain the large deviation from the relation observed for brighter bulges. It appears that these bulges are so small that they are easily affected by external influences, like e.g. the disk. It should be noted that a similar behaviour is seen for faint ellipticals (Vader et al.1988). The trends of colour gradient with luminosity contrast with the lack of trends between colour gradients and galaxy type (Balcells & Peletier 1994). The luminosity of the bulge, and not the morphology of the entire galaxy, seems to control the population gradients in spiral bulges. In any case, colour gradients in bulges are comparable to those of elliptical galaxies. Our result is in agreement with some individual measurements of large bulges like those of M31 and NGC 4594. It appears as if the disk forms later than the bulge, and that its formation only very slightly affects the stellar populations of the bulge. Whether this is also true for the small bulges of later type spirals remains to be seen.

274

Balcells k Peletier: Colours of Bulges

REFERENCES BALCELLS, CARLBERG,

M. & PELETIER, R. F. 1994 Astron. J. 107, 135. R. G. 1984 Astrophys. J. 286, 403.

MCWILLIAM, A. & RICH, M. R. 1994 In press.

R. F., BALCELLS, M. 1994 In preparation. R. F., DAVIES, R. L., ILLINGWORTH, G., DAVIES, L. & J. 100, 1091.

PELETIER, PELETIER,

CAWSON,

M. 1990 Astron.

RICH, S. M. 1988 Astron. J. 95, 828.

F. &; SEITZER, P. 1992 Astron. J. 104, 1039. VADER, J. P., VIGROUX, L., LACHIEZE-REY, M. &I SoUViRON, J. 1988 Astron. Astrophys. 203, 217. VlSVANATHAN, N. & SANDAGE, A. 1977 Astrophys. J. 216, 214. WHITFORD, A. E. 1978 Astrophys. J. 226, 777. SCHWEIZER,

WIRTH, A. & SHAW, R. 1983 Astron. J. 88, 171.

Implications of Galaxy Alignment for the Galaxy Formation Problem By WLODZIMIERZ

GODLOWSKI

Astronomical Observatory, Jagiellonian University, Orla 171, 30-244 Krakow, Poland The process of galaxy formation is one of the crucial problems of modern astronomy. Galactic alignments are important as a test of the various available scenarios for galaxy origin which predict different types of alignments. A method for investigating the galactic rotational axes is applied to two samples of galaxies chosen from the UGC, ESO and NGC catalogs for testing different models of galaxy formation. In the whole Supercluster the planes tend to be oriented perpendicularly to the Local Supercluster (LSC) plane. The effects strongly depend on the supergalactic coordinates. We compare the observed distribution of galactic rotation axes with theoretical models. Our results support the so-called "pancake" or "hedgehog" galaxy formation scenario and exclude the "turbulence" models. Moreover, we have some evidence on the importance of membership of clusters belonging to the LSC.

1. Introduction Galactic alignments are a crucial problem for understanding the process of galaxy formation. Various scenarios of galaxy origin predict different types of galaxy alignments within superclusters. Analysis of LSC galaxies (Flin & Godlowski 1986; Godlowski 1991, 1992, 1993) has shown that the preferred orientation of the galactic plane is perpendicular to the LSC plane, and that the projection of the rotational axis on the LSC plane tends to be directed towards the Virgo Cluster center. The distributions of face-on and edge-on galaxies are different. The orientation depends on supergalactic latitude B. In the clusters of the LSC the alignment is normally observed; however the direction of departures from isotropy is different for various clusters. Here, we compare the position angles of the rich clusters in the LSC, to detect any non-random trends. Afterwards, we compare the observational distributions of galaxy planes with theoretical models.

2. The method and observational data Historically, two main methods for studying galaxy orientation were proposed. The first one, by Hawley & Peebles (1975, hereafter HP), consists of an analysis of the observed position angles of the galactic image major axes. These angles yield reliable information about orientations of the galactic planes only for edge-on galaxies with small absolute values of B. The other approach, proposed by Jaaniste and Saar (1977), uses face-on galaxies. They also considered the galaxies' inclination, i, with respect to the observer's line of sight. In fact, we analyse the distribution of two angles, given the orientation of the vector normal to the galactic plane: 6, the angle between normal to the galaxy and LSC plane and rj, the angle between the projection of this normal on the LSC plane and the direction towards the LSC center. We use two independent samples of galaxies. The first sample id based on UGC and ESO catalogs of galaxies and second is taken from Tully's Nearby Galaxy Catalogue (1988). Data from this catalogue contain the most probably distance to the galaxy. It also allows us to test the "hedgehog" model, given various orientations of angular momentum depending on the position of galaxies with respect to the Virgo Cluster. Do distributions of the angles 6 and r\ agree with 275

276

Godlowski: Implications of Galaxy Alignments for Galaxy Formation

theoretical predictions? We applied different statistical tests (originally proposed by HP) to the detailed analysis in Godlowski (1993). 3. The results We concluded that the preferred orientation of galactic plane distribution is perpendicular to the LSC plane. The projection of rotation axis on the LSC plane tends to be directed toward the Virgo Cluster center. We now tested whether the distribution of angles agreed with theoretical models. We obtained the theoretical distribution by Monte Carlo methods (including random variation of rotation axis from position given by the "clear" model). We tested a large number of possible models. We successfully fitted three: model 2, with galaxy planes perpendicular to the LSC plane and projection of rotational axes parallel to the direction toward our galaxy with the LSC center; model 3, as for model 2, but with rotational axes projected parallel to the direction towards a particular galaxy with the LSC center; and model 5, with the rotational axis for each galaxy pointing towards the LSC center. These models can explain the observed dependence. In models 3 and 5 we can explain the differences between face-on and edge-on galaxies, so these models are better than model 2. All three models can explain the dependence of the galactic orientation on the supergalactic longitude while only model 5 can also explain dependence on the supergalactic latitude B. We observed alignment in the orientation of galaxies of the rich cluster of the LSC in most cases, but direction departures from isotropy were different in the various clusters. Moreover, when we analysed the position angles of that group we can concluded most of them were about 90 degrees. This suggests the important role of clusters in the galaxy formation process. 4. Conclusions Evidence shows that galaxy planes tend to be perpendicular to the LSC plane. Rotational axis projections on the supergalactic plane tend to point toward the Virgo Cluster center. Because of the shape of the LSC a small number of galaxies have high B. This may result from the perpendicularity of galaxy planes to the vector towards the Virgo cluster center according to the "hedgehog" model. This model explains the dependence of the orientation of galaxies on the supergalactic coordinates and accounts for the observed differences between face-on and edge-on galaxies. Because of the shape of the LSC, it has not been possible till now to finally decide between the two main probable models: "pancake" scenarios and the "hedgehog" model. Analysis of the orientation of galaxies within the clusters, together with alignments of the cluster position angles, suggests the major role of clusters of galaxies in the galactic formation process. REFERENCES FLIN, P. & GODLOWSKI, W. 1986 Mon. Not. Roy. Astr. Soc. 222, 525.

GODLOWSKI, W. 1991 In Galaxy Environments and the Large Scale Structure of the Universe (ed. G. Giuricin), p.146. SISSA-ISAS. GODLOWSKI, W. 1992 In . The feedback of chemical evolution on the stellar content of galaxies.

Ill DAEC Meeting (ed. D. Alloin & G. Stasinska), p.350. Observatoire de Paris. GODLOWSKI, W. 1993 Mon. Not. Roy. Astr. Soc. 265, 874. HAWLEY, D. L. & PEEBLES, P. J. E. 1975 A. J. 80, 477. JAANISTE, J. & SAAR, E. 1977 Tartu Obs. Preprint A-2. TULLY, R. B. 1988 Nearby Galaxy Catalog. Cambridge University Press.

Annular Structure Analysis of the Starburst Spiral Galaxy NGC 7217 By A. M. VARELA 1 , M. PRIETO 1 , A. K. VIVAS 2 AND C. MUNOZ-TUNON 1 1

Instituto de Astrofisica de Canarias, Via Lactea, E-38200 La Laguna, Tenerife, Spain

2

Centro de Investigaciones de Astronomia, Apartado 264, Merida, Estado de Merida, Venezuela

Circumnuclear starbursts can occur in galaxies with no notable companions. One of the consequences of this mechanism is to sweep out the star-forming gas from the few inner kiloparsec around the nucleus causing annular structures. Repeated starbursts (and perhaps the initial starburst itself) require a mechanism to feed new gas towards the axis of rotation. A dynamical process for this refuelling is based on the effects of departures from axisymmetry in the gravitational potential of the bulge, which can cause a net gaseous inflow, seen at its densest along a bar or bar-like structure, but also present for more generalized oval distortion. In an initial phase of the present work we have analysed the bulge component of a sample of spiral galaxies obtained using the 4.2-m WHT and the 2.5-m INT at the Observatorio del Roque de los Muchachos (La Palma). These galaxies present a starburst signature or exhibit evidence of residual phenomena from a previous starburst (see Beckman et al. 1991; Varela 1992). Seeing values oscillated between 0-5 to 1-4. Data reduction was performed using FIGARO and IRAF standard programs. We carried out a photometric analysis which enable us to perform a bulge/disc emission decomposition, in order to reconstruct that associated only with the bulge (Prieto et al. 1992a,b; Varela 1992). For each individual object under study, we intend to model a three-dimensional structure capable of reproducing the observations (Simonneau et al. 1993; Varela et al. 1994) and assuming a M/L ratio, we obtain the multipolar moments of the gravitational potential associated with the bulge component. Here we present initial photometric results for the isolated Liner galaxy NGC 7217.

1. Analysis and preliminary results Figures l a and b show the isophotal images of NGC 7217 in the B and / bands. In the B filter we clearly see the inner and bluer ring at 600 pc from the nucleus and the external and redder one at about 6 kpc. However, this structure is not present in / . Figures 2a and b gives the B — I colour mean profile and the bulge and disc brightness profile decomposition in the / band respectively. For the bulge re = 3.1 kpc and the scale length of the disc is 2.06 kpc. The B — I colour image of NGC 7217 (not shown) displays a series of rings of different angular sizes. The size increases with radius, suggesting the presence of a density wave, which becomes relaxed when far from the nucleus. In Figures 3a and b are plotted the ellipticity and position angle profile in both filters against radius. We see a sharp rise in ellipticity at 5" from the nucleus in both filters, and a sharp fall to low values between 0.03 and 0.06 in the first 5" from the nucleus. It rises toward 0.9 at around 8" and finally reaches its disc value of 0.12. The position angle behaviour shows a clear difference between the bulge and disc regions. These characteristics indicate a non-axisymmetric bulge. Otherwise it would show a steady rise from its near-nuclear value to the disc value, suggesting that we are seeing essentially the effects of inclination of the galaxy as we pass from the spheroidal bulge to the planar disc. 277

278 Varela et al.: Annular Structure Analysis of the Starburst Spiral Galaxy NGC 7217

N

J

FIGURE 1. Isophotes of our CCD frames of NGC 7217: a) In the B band; b) In the / band. The outer contours correspond to 22.68 and 21.20 mag arcsec2 respectively and the spacing used is 0.5 mag. The scale is 38.3 pc arcsec"1.

0

20

40

60

80 100

0

20

40

orcseconds

60

80 100

orcseconds

FIGURE 2. a) B — I mean colour profile and 6) mean radial brightness profile in the / filter of NGC 7217. .filter B

b)

)

50

100

f] 0

orcsecond* FIGURE

.

****a$$t«tia«i

u _

d

.filter I

50

100

arcttcond*

3. a) Position angle and 6) ellipticity vs. radius for NGC 7217 in B and / bands.

2. Conclusions The presence of non-axisymetric structures in the bulge of NGC 7217 has been demonstrated by using precision photometry with a CCD camara under outstanding optical conditions. This galaxy shows strong non-circular gas motions in its nucleus. The orientation of the nuclear region relative to the galactic disc indicates that it has a bulge without axial symmetry. NGC 7217 presents a multiple annular structure like an expanding wave from the nucleus, with a filamentary appearence. REFERENCES BECKMAN, J. E., VARELA, A. M., MUNOZ-TUNON, C ,

VILCHEZ, J. M. &; CEPA, J. 1991

Astron. Astrophys. 245, 436. PRIETO, M., LONGLEY, D. P. T., PEREZ, E., BECKMAN, J. E., VARELA, A. M. & CEPA, J.

1992a Astron. Astrophys. Suppl. Ser. 93, 557. M., BECKMAN, J. E., CEPA, J. & VARELA, A. M. 1992b Astron. Astrophys. 257, 85.

PRIETO,

SIMONNEAU, E., VARELA, A. M. & MUNOZ-TUNON, C. 1993 J. Q. S. R. T. 49, No.2, 149.

VARELA, A. M. 1992 PhD thesis, University of La Laguna, Tenerife, Spain. VARELA, A. M., SIMONNEAU, E. & MUNOZ-TUNON, C. 1994 Astron. Astrophys. Submitted.

How a dust concentration mimics dynamical signatures around the nucleus of NGC 7331 By F. PRADA 1 - 2 , J. E. BECKMAN 1 AND C. D. McKEITH 2 1

Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain 2

Queen's University of Belfast, Physics Dept., Belfast BT7 INN, U. K.

Where the phenomenon of wavelength-dependent kinematics has been observed in galaxies, it has been well modelled as a dust extinction effect. Therefore "compact mass features" in rotation curves close to the nuclei of dusty inclined galaxies may not be caused by dynamical effects, but by differential extinction by circumnuclear dust. NGC 7331 exhibits this phenomenology.

1. Introduction The kinematics of the zones around the nuclei of galaxies are of exceptional interest, especially in the context of tests for very compact central objects (black holes). A number of nearby galaxies show striking kinematic features around their nuclei and have been well modelled by adding a central point mass to an otherwise smoothly varying bulge distribution, yielding the steep velocity gradients, dispersions and "shoulders" in their rotation curves (Bower et al. 1993). Given the exceptional interest in the presence of supermassive compact objects, however, it may not be surprising that, in at least some cases, remarkable circumnuclear kinematical effects may have been overinterpreted. In previous spectroscopy of dusty, highly inclined galaxies we have found systematic steepening of velocity curves from the near-UV to the near-IR (McKeith et al. 1993), which is convincingly modelled via dust extinction (Prada et al. 1994). 2. Data and results To see how the observable rotation curve of the dusty spiral inclined galaxy NGC 7331 varies with wavelength we took long-slit spectra with the ISIS spectrograph on the 4.2-m WHT (La Palma), with the slit along the major axis centred on the optical nucleus. The observed rotation curves are shown in Figure 1. The curves in Ha and [Nil] A6584 A in emission from interestellar gas (HII regions) and Ball A6496 A in absorption from stars show a very steep central gradient, with sudden turnover and shoulders at either side of the nucleus between 350 pc and 700 pc from the centre, followed by a slower, bumpy rise to a steady "plateau" velocity 250 km s" 1 some 2.8 kpc from the centre. There are some differences between the curves, but there is not a dichotomy between gas and stars. In Figure 1 we compare Call A8542 A with [Nil] A6584 A and we can see that the curve in Call does not show the shoulders at 350 pc seen in [Nil], Ha and Ball.

3. Discussion The most coherent explanation of the velocity curves in Figure 1 is that the central 350 pc of the galaxy is essentially unobscured, allowing us to see the kinematics which truly reflect the dynamics; we compute a mass of 2 109 M© within this radius. The kinematic structure seen in the "red" curves beyond this radius is, however, an artefact due to 279

280

Prada et ai: How a dust concentration mimics dynamical signatures 1 ' '

-

' 1 ' '

1

1

,

,

,

i

,

,

,

,

,

,

,

,

.

•

•

.

1100 A

Ha 6562 A.

D

D 1000 --

A

Nil 6584 A,

gas

D

-

900

; o

Ball 6496 A stars

D

Call 8542 A stars

_--

rj^A ^

gas

A&

A*

SE

-

J-JDOD AA^ Q Q A«*
A^

A

4^e M

-

AT D

- 800 -

-

-

A A) 700 - -

A/^^AA

NW

A

-

-

g&

a a

_ -

p.a,167°

600 : -

500 -80

1

1

-60

-40

,

,

,

1

-20

,

i

,

, , 1 , , , 1 , , , 1 , , ,"

0 Position (arcsec)

20

40

60

80

FIGURE 1. Observed rotation curves, along the major axis of NGC 7331, in Ha A6562 A , [Nil] A6584 A interestellar emission and in Ball A6496 A and Call A8542 A stellar absorption. The "shoulders" disappear in the Call curve due to the greater light penetration through the dust at longer wavelengths.

extinction by a toroidal, though somewhat patchy, dust distribution. The agreement between the red curves (Ha, [Nil], Ball) and the difference for the near-IR curve (Call) shows that the effect is wavelength dependent rather than component dependent. There is not one kinematic regime for the gas and another for the stars. The greater dust transparency at longer wavelength allows the Call IR triplet curve to show a closer approximation to the dynamical rotation curve than do the curves for the other species. Qualitatively similar effects of dust blocking, stronger in the red than in the infrared, cause the shoulders in Fig. 1, which are not, on this interpretation, due to strong concentration of mass near the centre, as claimed by Afanasiev et al. (1989). We thank Dr. Robin Clegg for the observations taken on a service night (August 1992) at the Spanish Observatorio del Roque de los Muchachos of the IAC. We thank Alicia Gordillo for valuable comments on the manuscript.

REFERENCES AFANASIEV, V. L., SIL'CHENKO, O. K. & ZASOV, A. V. 1989 A. & A. 213, L9. BOWER, G. A., RICHSTONE, D. O., BOTHUN, G. D. & HECKMAN, T. M. 1993 Ap. J. 402, 76. MCKEITH, C. D, CASTLES, J., GREVE, A., & DOWNES, D. A. & A. 272, 98. PRADA. F., BECKMAN, J. E., & MCKEITH, C. D., CASTLES, J., & GREVE, A. 1994 Ap. J.

Lett. In press.

UGC 5101: An Ultraluminous IRAS Galaxy with Circumnuclear Star Formation By V. RESHETNIKOV 1 - 2 AND F. COMBES 1 ^EMIRM, Observatoire de Paris-Meudon, F-92195 Meudon, France 2

Astronomical Institute of St. Petersburg University, 198904 St. Petersburg, Russia

New spectral and photometric observations of the ultraluminous infrared galaxy UGC 5101 (IRAS 09320+6134) suggest that a ring of giant HII regions surrounds the Seyfert nucleus.

1. Observations and results Prime-focus photometry of UGC 5101 in B, V, and Rc was performed with a CCD at the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences. Figure 1 shows the isophotal map in B (the faintest contour corresponds to a surface brightness of 26 magarcsec"2, the step between isophotes is 0.75 mag). The general photometric characteristics of the galaxy are summarized in the Table (the blue luminosity and colours are corrected for Galactic absorption according to RC3). £ ( t f o = 75 kmsr 1 Mpc- 1 )

160 Mpc

(B - Vft

15.5 ±0.10 +0.72 ±0.03 +0.55 ±0.02

M/L B (< 14")

-20.5 [2.5 1010L©(5)] 200 km s" 1 4M0/L0(5)

a (HB = 26) b/a (HB = 26)

75" (58 kpc) 0.61

MB

vmax

The spectral observations (6200-7170 A) of the galaxy were carried out at the Cassegrain focus of the 1.93-m telescope of the Observatoire de Haute-Provence, using the CARELEC spectrograph and a TK512 CCD. The dispersion was 1.78 A pixel"1, the slit width 2-5. We obtained two spectra - with the slit along the major axis of the galaxy (P.A.=85°) and at P.A.=0°. In Ha and [Nil] emission the galaxy demonstrates fast (Vmax « 200 km s"1) rotation around the minor axis. The spectrum with the slit along the minor axis shows that the gas in the central part (r < 10") of the galaxy rotates around the major axis at Vmax w 50 km s"1 and probably belongs to the luminous arc or ring (see Figure 1). The spatial structure of the Ha and [Nil] AA6548,6583 emission lines at P.A.= 85° and 0° is shown in the isophotal maps reported in Figure 2. 2. Discussion The complicated morphology and kinematics observed suggest that UGC 5101 is a merger remnant. The results of our surface photometry also support this conclusion. 281

282

Reshetnikov & F.Combes: UGC 5101: An Ultraluminous IRAS Galaxy

N

FIGURE 1. Isophotal contour map of UGC 5101 in the B filter.

P.A. - 85

FIGURE

o

P.A. - 0

2. Spatial structure of Ha and [Nil] emission lines at two position angles of the slit (bars indicate a length of 10 arcsec).

Using the equivalent luminosity profiles for the averaging of the general photometric structure of the galaxy, we found that at req > 4" the equivalent profiles fit the de Vaucouleurs law well with Re = 5.25 ± 0.4 kpc (mean value of three passbands). Such an elliptical structure is typical for merger remnants. It is interesting also that the derived equivalent characteristics of the galaxy are located in the plane of effective parameters (/ie - Re) almost exactly on the standard relation for bright normal galaxies. There are indications that the objects belonging to this "bright family" have probably undergone merging phenomena (Capaccioli et al. 1992). As one can see in Figure 2a, the Ha emission demonstrates clear double-peaked structure in the central (r < 3") region of the galaxy. With care, three separate emitting regions can be distinguished: a central Seyfert nucleus with [Nil] A6583/Ha = 1 . 1 and FWHM(Ha) = 660 km s" 1 , and two regions at approximately 2" (1.5 kpc) on either side of the center, in which the line emission is dominated by the Ha line. Note that in the regions dominated by the Ha emission the lines are relatively narrow (FWHM < 300 km s"1) in comparison with the central active nucleus. Given the size of these Ha dominated regions (« 1.5 kpc) and their total luminosities (L(Ha) « 2 — 31040 erg s" 1 ), we suppose that they are giant HII complexes. Figure 2b shows that at P.A.= 0° the Ha line does not display such a multi-component structure. This probably implies that we observe a nearly edge-on ring of strongly enhanced star formation surrounding the Seyfert nucleus of the galaxy. A similar circumnuclear star formation has been recognized in several other Seyfert galaxies (for instance, Boer & Schulz 1993).

REFERENCES BOER, B. & SCHULTZ, H. 1993 A. & A. 277,

397.

CAPACCIOLI, M., CAON, N. & D'ONOFRIO, M. 1992 M. N. R. A. S. 259,

323.

The Stellar Content of Nearby and Distant Starburstsf ByCLAUS LEITHERER Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA We studied the content of massive stars in the nearby HII region 30 Doradus and in the distant nuclear starburst NGC 7552. Ultraviolet imaging and ultraviolet spectroscopy with the Hubble Space Telescope and the International Ultraviolet Explorer have been obtained. The observational data have been compared with evolutionary population synthesis models in an attempt to constrain the star-formation history and mass spectrum of the two starburst regions. Despite the different physical and chemical conditions, the high-mass end of the mass functions in 30 Dor and NGC 7552 are remarkably similar.

1. 30 Doradus and NGC 7552 — basic properties 30 Doradus in the Large Magellanic Cloud is the closest giant extragalactic HII region (GEHR). The combination of its relative proximity and its location outside the plane of our Milky Way makes 30 Dor a prime laboratory for the study of the formation and evolution of the most massive stars. Walborn (1991) reviewed the basic properties of 30 Dor, designating it the "Starburst Rosett". 30 Dor is metal-deficient with respect to the Sun (Z = 0.3 ZQ) and has only a moderate interstellar reddening of E(B — V) = 0.4. At a distance of 0.05 Mpc, 1" corresponds to a linear distance of 0.25 pc. The most massive stars are concentrated within a few arcsec of the center of 30 Dor. Only the Hubble Space Telescope (HST) provides sufficient spatial resolution to do the crowdedfield photometry necessary to investigate the stellar content. HST data suggest the presence of an extremely young (T < 4 Myr) stellar cluster rich in stars with zero-agemain-sequence (ZAMS) masses above 50 MQ (Heap et al. 1992; De Marchi et al. 1993; Malumuth, this meeting). NGC 7552 is an SBbc galaxy with a nuclear starburst, whose properties have been described by Sersic & Pastoriza (1965) and Feinstein et al. (1990). The nucleus is metal rich (Z — 1 — 2 ZQ) and dusty, as indicated by the high reddening of E(B — V) = 0.7. NGC 7552 is at a distance of 30 Mpc (for Ho = 50 kms^Mpc" 1 ) so that 1" corresponds to 150 pc. The high far-IR luminosity, the optical emission lines, and the WE spectrum suggest strong starburst activity (Kinney et al. 1993). However, in contrast to 30 Dor, relatively few quantitative results on the properties of the starburst are available. One of the reasons is the distance of the galaxy and the correspondingly small angular scale which have made it difficult in the past to resolve the starburst region. Yet a comparison between 30 Dor and NGC 7552 would provide invaluable insight into the physics of starburst due to the large difference in the chemical composition. 30 Dor and NGC 7552 differ in angular scale by a factor of 600. Observations should sample comparable regions for a meaningful comparison. Therefore wide-angle imaging and spectroscopy are required for 30 Dor, whereas data for NGC 7552 must have sub-arcsecond resolution. In this paper we report on such observations and their interpretation. f Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by AURA for NASA under contract NAS5-26555. 283

Leitherer: Nearby and Distant Starbursts

284

420 pc

FIGURE 1. Nuclear region of NGC 7552 imaged with HSTs FOC f/96 at a wavelength of 2200 A. No image restoration has been done. Resolution: 0"05. Several individual starburst knots are visible.

2. Spatial morphology A high-resolution ultraviolet image of the nucleus of NGC 7552 is shown in Figure 1. II

a

The image has a spatial resolution of 0 • 05 and an effective wavelength of 2200 A. At this wavelength, most of the light is due to massive stars with masses above 5 MQ. Some contribution from light scattered by dust grains is expected as well. The starburst was previously thought to be "nuclear", i.e. confined to one nuclear region. It turns out that the nucleus is actually a cluster of individual smaller starburst regions, each having comparable brightness. Saturation effects make flux estimates somewhat uncertain but each of the individual starburst regions appears to be several times as luminous as the 30 Dor complex. In the absence of any color information one can only speculate about the actual distribution of matter but it seems likely that some of the structure visible in Figure 1 is caused by variable dust extinction. The morphology of the NGC 7552 nucleus is strikingly similar to the structure of the starbursts observed with HST in NGC 3690 (Leitherer 1994), He2-10 (Vacca, this meeting) and M83 (Heap, this meeting). If a wave of star formation propagates at the sound speed, a characteristic timescale for the burst in the center of NGC 7552 is 1 - 2 x 107 yr. This is longer than the evolutionary timescale of the most massive stars. Cheng et al. (1992) obtained wide-field ultraviolet images of 30 Dor with the Ultraviolet Imaging Telescope. Their Figure 5 (at a wavelength of 2500 A) covers approximately the same linear size of 30 Dor as does our Figure 1 for NGC 7552 so that the two figures are directly comparable. Strikingly, the central ionizing star cluster of 30 Dor and the central region of NGC 7552 have very similar sizes and morphologies. On the other hand, comparison of the IRAS far-IR luminosities suggests that the rate of massive-star formation is higher in NGC 7552 by about 2 orders of magnitude. Therefore the density of massive stars is correspondingly higher in the central region of NGC 7552.

Leitherer: Nearby and Distant Starbursts

285

3. The ultraviolet as an indicator of massive stars 3.1. Method The ultraviolet is the only wavelength region where spectral lines of hot stars in a starburst are directly observable. The region between 1200 A and 2000 A can be used to investigate starburst regions by synthesizing the continuum and the line profiles due to OB stars. For brevity, I will concentrate here on the two strongest features of SilV A1400 and CIV A1550. These lines are largely formed in O-star winds with outflow velocities in the order of 2000 km s~ . The mass-loss mechanism in hot stars is well understood (Kudritzki et al. 1991). Radiative momentum (L/c) is transformed into kinetic momentum (Aft)) by absorption of photons in the wind. Therefore a relation exists between the stellar luminosity L and the outflow velocity v, and the line profile in general. As there is a stellar mass-luminosity relation, a relation between stellar mass and the SilV A1400 and CIV A1550 velocity is expected. Robert et al. (1993) pursued this idea in a quantitative way. They combined the latest generation of stellar evolutionary models and stellar atmospheres with an WE high-dispersion library of OB- and Wolf-Rayet stars in order to do an evolutionary spectral synthesis in the ultraviolet. The synthetic CIV A1550 is aJways dominated by O stars which produce this line in their stellar winds. Therefore a strong, blueshifted absorption and sometimes a weaker emission component is present. SilV A1400 is not observed in the winds of O mainsequence stars and it is also rather weak in their photospheres, as Si is predominantly Si 4+ . At about 3 Myr the first O supergiants appear. They have higher mass-loss rates and denser winds leading to increased recombination to Si 3+ . Accordingly, blueshifted wind absorption in SilV A1400 is discernible in the synthetic spectrum. Blueshifted SilV A1400 in a starburst spectrum uniquely implies the presence of an evolved stellar population. 3.2. Application to 30 Dor and NGC 7552 30 Dor is an ideal test site for these models. It is close enough that its stellar content can be resolved and studied in detail by spectroscopic means (Parker 1993; Parker & Garmany 1993). Consequently, the stellar mass spectrum can be derived independently and the stellar census can be used to predict the synthetic ultraviolet spectrum, including line profiles and the Lyman discontinuity. Vacca et al. (in preparation) obtained two trailed WE spectra covering an area of 1' by 1' and 3' by 3'. This matches the area studied in Parker's spectroscopic stellar census. The WE data allow one to constrain the stellar population in two independent ways. First, lines of SilV A1400, CIV A1550, and Hell A1640 indicate the age of the burst and the relative contribution of the most massive stars. Second, the absolute continuum fluxes in combination with the observed number of Lyman continuum photons provide a measure of the Lyman discontinuity, and therefore of the number of the hottest O stars. Note that the Lyman photons are not counted via nebular diagnostics. Rather, they are obtained directly from the stellar census and the spectral-type calibration. This eliminates a significant source of uncertainty typically encountered in studies utilizing nebular emission lines. A comparison between the observed WE spectrum and several synthesis models is shown in Figure 2. A stellar population with an upper mass cut-off of Mup > 50 MQ observed 2 to 3 Myr after the starburst is derived. The strong Hell feature and the associated large number of short-lived Wolf-Rayet stars requires the burst duration to be short in comparison with the stellar life-times. Both the continuum and the spectral line fitting give the same result. It is gratifying to note that this agrees with isochrone-fitting techniques (e.g. Heap et al. 1992) as well.

Leitherer: Nearby and Distant Starbursts

286

4 Myr

• 1 2 0 M.

I——I——I——I——i

1300 1400 1500 1600 1700 1800 WAVELENGTH (A) FIGURE 2. Observed model spectra (solid: of 1 to 4 Myr. Right ijj(M) oc M~a with a

3 Myr

.

2 Myr V

1 50 M_

1 Myr I

30

80 M.

1300 1400 1500 1600 1700 1B00 WAVELENGTH (A)

WE spectrum of 30 Dor in a 3' by 3' region (thick line) versus various Z = ZQ\ dotted: Z = 0.25 ZQ). Left panel: Mup = 100 and burst ages panel: T = 2.5 Myr and Mup from 30 MQ to 120 M©. Power-law IMF = 2.35 in all cases.

A comparison between the trailed WE spectrum (1' by 1' area) of 30 Dor and a pointed HST FOS observation of the nucleus of NGC 7552 is shown in Figure 3. The FOS data are part of an HST GO program headed by T. Heckman and will be discussed in full detail by Robert et al. (in preparation). The circular entrance aperture of the FOS has a diameter of 1". The small FOS aperture and the large area swept by the WE allow a direct comparison of the two starburst regions. The linear sizes of 30 Dor and NGC 7552 in the two spectrograph apertures are roughly comparable (although NGC 7552 is still larger by a factor of 3 to 10). The most obvious difference between the 30 Dor and the NGC 7552 spectrum is the strength of the interstellar lines. The nucleus of NGC 7552 is strongly obscured by dust. This is evident from the very red color of the ultraviolet to visual continuum. At the same time absorption lines due to interstellar gas are also very strong in this galaxy (cf. Kinney et al. 1993). Note that the spectral resolution of the FOS is a factor of 3 higher than for the WE. Therefore the unresolved interstellar lines (like Sill A1260 or CII A1335) are narrower in NGC 7552. Strong absorption lines around 1900 A are visible in NGC 7552 but are absent in 30 Dor. These lines are due to Fell/Ill and AMI. They are the signature of a very evolved population of B and A supergiants. Such stars are not present in the central region of 30 Dor, which is clearly a younger starburst region. On the other hand, one should keep in mind that the FOS aperture encompasses a comparatively larger region in NGC 7552. Most likely, regions outside the main starburst knots visible in Figure 1

Leitherer: Nearby and Distant Starbursts

287

1.25 -

1200

1400

1600

1800

2000

Wavelength (angstroms)

FIGURE 3. Top: WE spectrum of the central l' by l' region of 30 Dor. Bottom: HST FOS spectrum of the nucleus of NGC 7552. The continuum of both spectra was normalized to the unity level. The vertical scale refers to NGC 7552. 30 Dor is displaced upward by 0.75 units. The instrumental resolution is 6 A and 2 A for 30 Dor and NGC 7552, respectively. also contribute to the spectrum in Figure 3. These regions may be physically unrelated to the starburst itself. The broad, stellar SilV A1400 is much stronger in NGC 7552 than in 30 Dor. This line is formed in O-supergiant winds. Evolved O supergiants have not yet formed in large numbers in 30 Dor but are numerously present in NGC 7552. This implies a much older age of the starburst in NGC 7552. The radial velocity of the wind-sensitive lines of SilV A1400 and CIV A1550 provides an excellent estimate of the relative number of the most massive stars (cf. Robert et al. 1993). The two lines have blueshifts (in the restframe of NGC 7552) of 800 km s" 1 (SilV) and 1100 km s" 1 . A significant contribution from stars with ZAMS masses above 60 M© is required to account for this blueshift. Therefore the upper ends of the stellar initial mass function in 30 Dor and NGC 7552 are very similar. No strong O-star emission of CIV A1550 and no Wolf-Rayet emission of Hell A1640 is detected in NGC 7552. This does not argue against large numbers of the most massive stars. Rather, it is the natural consequence of a relatively evolved starburst where star formation and stellar death have reached equilibrium (after T > 5-10 Myr). On the other

288

Leitherer: Nearby and Distant Starbursts

hand, despite the evolved nature of the NGC 7552 starburst, high-mass star formation is still proceeding vigorously, as indicated by the strength and blueshift of the SilV A1400 absorption component.

4. Conclusion The stellar populations presently observed in 30 Dor and NGC 7552 exhibit marked differences. These differences result from a combination of different star-formation histories, age effects, and the influence of different aperture sizes. After taking these effects into account, it is found that the upper end of the stellar initial mass function is quite similar in both objects. The chemical composition of 30 Dor and NGC 7552 differs by about a factor of 5. Yet both regions are able to form similar fractions of very massive stars. Clearly, further studies with a larger observational sample are needed to address the question of a metallicity-dependent initial mass function. This paper resulted from two individual projects. The work on 30 Dor was done in collaboration with Bill Vacca, Peter Conti, and Carmelle Robert. NGC 7552 is part of a systematic study of the stellar content of starbursts performed by Tim Heckman, Carmelle Robert, Don Garnett, Anne Kinney, and myself. Support for this work was provided by NASA through grant numbers GO-3591.01-91A and GO-3605.02-91A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. REFERENCES CHENG, K.-P. ET AL. 1992 Ap. J .Lett. 395, L29. DE MARCHI, G., NOTA, A., LEITHERER, C , RAGAZZONI, R. & BARBIERI, C. 1993 Ap. J. 419,

658. FEINSTEIN, C , VEGA, I., MENDEZ, M. & FORTE, J. C. 1990 A & A 239,

90.

S. R., EBBETS, D. & MALUMUTH, E. 1992 In it Science with the Hubble Space Telescope (ed. P. Benvenuti & E. Schreier), p. 347. European Southern Observatory, Garching.

HEAP,

KINNEY, A. L., BOHLIN, R. C ,

CALZETTI, D., PANAGIA, N. & WYSE, R. F. G. 1993 Ap.

J.

S. 86, 5. KUDRITZKI, R. P.,

GABLER, R.,

KUNZE, D.,

PAULDRACH, A. W. A. &; PULS, J. 1991

In

Massive Stars in Starbursts (ed. C. Leitherer, N. Walborn, T. Heckman &: C. Norman), p. 59. Cambridge University Press. LEITHERER, C. 1994 In Reviews in Modern Astronomy 7 In press. PARKER, J. W. 1993 A. J. 106,

560.

PARKER, J. W. & GARMANY, C. D. 1993 A. J. 106,

1471.

ROBERT, C , LEITHERER, C. & HECKMAN, T. M. 1993 Ap. J. 418, SERSIC, J. L. & PASTORIZA, M. 1965 P. A. S. P. 77,

749.

287.

WALBORN, N. R. 1991 In Massive Stars in Starbursts (ed. C. Leitherer, N. Walborn, T. Heckman & C. Norman), p. 145. Cambridge University Press.

WR Stars in the Giant HII Region NGC 4236IIIf By ROSA M. GONZALEZ AND ENRIQUE PEREZ Institutes de Astroffsica de Canarias, Via Lactea, E-38200 La Laguna, Tenerife, Spain. We present long-slit optical and near-infrared spectroscopy of the giant HII region NGC 4236III. We have found broad emission lines at 4686 A attributed to WR stars. We have derived the physical conditions and chemical composition of the nebula.

1. Introduction HII regions are one of the most useful tools to study the properties of massive stars as well as the physical conditions and chemical composition of the interstellar medium. One of the target of the GEFE programme is the giant HII region NGC 4236III located in the outskirts of the SBd galaxy NGC 4236. The object was observed with the 4.2-m WHT telescope in La Palma, using the blue and the red arms of the ISIS spectrograph and an EEV CCD in each arm. The dispersion was 1.4 Apix"1, and the spatial sampling 0.33 arcsec pix" 1 .

2. Results The emission of the region is extended over 15 arcsec. Three different spectra were extracted (A, B, C). These spectra are typical of high-excitation HII regions. In B (Figure la), where most of the continuum emission is concentrated, we have detected a broad emission bump at 4686 A which is attributed to WR stars. The distribution of the emission lines Ha, [OIII], [Oil], [Nil] and [SIII] are quite similar; they show a maximum at 1.5 arcsec to the North of the peak of the continuum distribution. The [Nil] distribution presents an additional peak at the continuum maximum. This excess of the [Nil] emission with respect to [Oil] represents an overabundance of N in the position where the WR stars are detected. The [OIII]/[OII] distribution around the central core shows a double peak separated by 3 arcsec, with a local minimum coincident with the position of the WR stars (Figure lb). Regarding the physical conditions and chemical composition, we have calculated the electron temperature, density and reddening for the three different spectra. The region is in the low density limit, Ne <100 cm" 3 and the electron temperature is almost constant, —11500 K. We have not found any significant reddening variations. The region is characterized by a moderately low abundance (Table 1). As for the WR population, the presence of the NIII and NV blends in the WR bump at 4686 A indicates the presence of WN5-4. The equivalent width of the bump is 28 A and the luminosity 10 3753 erg s" 1 , which implies that 11 WN are present in the cluster. We estimate that about 200 07V ionizing stars are required to account for the H/? luminosity. The ratio L(WR bump)/L(H/3) and the relative number WR/0 are respectively 0.05 and 0.052. t The data presented in this contribution are part of the GEFE collaboration. GEFE, Grupo de Estudios de Formacion Estelar, is an international collaboration of astronomers from Spain, the U.K., France, Germany, Denmark and Italy, formed to take advantage of the international time granted by the Comite Cientifico Internacional at the Observatories in the Canary Islands. 289

290

Gonzalez-Delgado k, Perez: WR Stars in HII Regions

u ([<)III]) A B C

1. 14 1. 16 1. 17

1!e([Oil])

<e([SIII])

<e([SII])

1.21 1.17 1.25

1.21 1.42 4

TABLE 1. tc is in units of 10 K, O/H

4000

5000

cm

O/H

0.23 0.29 0.20

8.21 8.16 8.13

S/H 6.38 6.55 6.43

N/O

He/H

-1.48 -1.40 -1.38

0.086

and S/H abundances are 12+log(X/H), and N / O is log(N/O).

6000

7000

8000

9000

wavelength (A)

o s

-10 FIGURE

-5

0,05 distance (arcsec)

1. (a) Merged spectrum of the cross-section B, where we detected the WR bump, (b) Spatial distribution of the ratios [OIII]/[OII] and [NII]/[OII].

3. Conclusion NGC 4236III is a giant HII region experiencing a recent burst of star formation. At the position of the continuum maximum we detect the presence of WR stars. The estimated age of this zone is between 3 and 4 Myr. At this position the [NII]/[OII] presents a maximum, indicating pollution by the wind from these stars. The ratio [OIII]/[OII], which is sensitive to the ionizing parameter, presents a local minimum, which could be explained by a decrease in the filling factor as a consequence of the WR winds. We gratefully acknowledge partial support from DGICYT Grant PB91-0531(GEFE) and the NATO Grant CRG920198 for collaborative research.

Consequences of High Mass Loss Rates on Wolf-Rayet Populations in Starbursts ByGEORGES MEYNET Geneva Observatory, CH-1290 Sauverny, Switzerland We discuss the effects of a change of the mass loss rates of massive stars on the outputs of population synthesis models of starbursts. We find that models with high mass loss rates well account for the observed ratios of WNL to O-type stars in starburst galaxies.

1. Stellar evolutionary tracks and starburst models The stellar populations as observed after a starburst episode depend on: • The intensity as a function of time of the star formation rate (SFR) during the starburst. The observed properties of the starburst may also depend on the much lower SFR if this prevails between bursts of star formation. However, for very intense bursts, these "underlying" populations can be completely blurred out by the numerous new born stars. • The slope of the initial mass function (IMF) and the lower and upper mass limits of the stars born in the burst. • The time elapsed since the beginning of the burst (which we shall call here the age of the starburst). • Some physical ingredients of the stellar models as, for instance, the mass loss rates by stellar winds and the metallicity. • The frequency of stars in close binary systems. For quite young starbursts however, this effect is certainly quite small, since for the most massive stars, the evolution does not depend so much on the fact that the star has a close companion or not. The mass loss rates are so high that the star will lose, without the aid of a close companion, a great part of its hydrogen envelope before the end of its main sequence evolution. Many of these effects have already been studied in detail (see for instance Arnault et al. 1989; Mas-Hesse & Kunth 1991; Cervino & Mas-Hesse this volume), the effects of a change of mass loss rates for the massive stars have not yet been fully examined. In this work, we take the opportunity of the recent grid of stellar models computed by Meynet et al. (1994), which differs from that of Schaller et al. (1992) by having twice the mass loss rate in the main sequence (MS) and WNL phases, to investigate this effect. Let us recall here that the mass loss by stellar winds is a key physical ingredient to the understanding of the evolution of massive stars (Chiosi & Maeder 1986; Maeder 1991). Various comparisons with the observed luminosities, chemical compositions and number statistics of WR stars (performed in zones of constant SFR) support the models with the enhanced mass loss rates, while the so-called standard mass loss rates (de Jager et al. 1988) lead to difficulties in all these comparisons (Maeder & Meynet 1994). Recently Smith et al. (1994) obtained a general good agreement between observed stellar masses of Wolf-Rayet stars in binary systems and those derived from the enhanced mass loss rate evolutionary tracks. This again gives some support to our models and adopted mass loss rates. Using these new theoretical evolutionary tracks, we shall try to answer the following two questions: 291

292

G. Meynet: Wolf-Rayet Stars in Starbursts

• To which extent are the results of starburst models based on the "standard" evolutionary tracks (those of Schaller et al. 1992; Schaerer et al. 1993a,b; Charbonnel et al. 1993) changed? • Is it possible, with the aid of the high mass loss rate models, which are to be preferred over the standard ones (Maeder & Meynet 1994), to account for the observed ratios of WNL to O-type stars in starburst galaxies? After a brief specification of the ingredients of starburst models in §2, we comment in §3 on the main changes brought about by the enhancement of the mass loss rates. We perform the comparisons with the observations in §4. The conclusions are drawn in §5.

2. The physical ingredients of the starburst models We consider the case of an instantaneous burst of star formation, i.e. we suppose that 20 000 stars are born at the same time, and distributed according to a Salpeter initial mass function (j^ — AM~X, with x = 2) in the mass range from 8 to 120 MQ (it is not useful to consider the lower-mass range since we are interested here in number ratios of high-mass stars). As stellar models we use the recent grids computed by Meynet et al. (1994) for £=0.001, 0.004, 0.008, 0.020 and 0.040 (the high mass loss rate case) and the models by Schaller at al. (1992), Schaerer et al. (1993a,b) and Charbonnel et al. (1993) for the same metallicities ("standard" mass loss rates). The criteria for a star to be considered as a WR, WNL, WNE and WC star are taken as in the papers just quoted. We define the O-type stars as all the H-burning stars having an effective temperature superior to 33 000 K (i.e. logTeff > 4.519). 3. Effects of a change of t h e mass loss rate In Figure 1, the ages at different evolutionary stages are given for stars of various initial masses for both the models computed with the standard mass loss rates and those obtained with an enhanced mass loss rate. The different areas delimited by continuous lines in these figures show the ranges of initial masses and ages corresponding to the O-type stars, and to the different Wolf-Rayet (WR) subtypes. From the comparison of the results obtained with the standard and the high mass loss rates, one can note the following points (see Maeder and Meynet 1994 for a more detailed discussion of the features described below): • Stronger stellar winds slightly decrease the time at which a star of given initial mass becomes a WR star. This effect is more pronounced in the high initial mass range. For the 120 MQ, the age at the entry point into the WNL stage passes from 2.25 to 1.96 106 yr. when the mass loss rates are increased by a factor of 2 during the MS phase. As a consequence, after a starburst, the WR will appear earlier in the case of stronger stellar winds (see Figure 2). • When the mass loss rates are higher, the WR lifetimes as a whole are considerably enhanced. As a consequence the WR-rich phase of a starburst (i.e. the period after the burst corresponding to the appearance of the WR stars) will be longer when the mass loss rates are higher. Since the O-type star phase remains nearly unaltered by a change in mass loss rate, stronger stellar winds produce higher number ratios of WR to O-type stars. • At solar metallicity, in the high initial mass range, the WNL and WNE phases are significantly enhanced in the case of high mass loss rates. In this mass range, the WNE phase may become as long or even longer than the WC phase (note that the WNE phase is quasi non-existent in the standard models). The WC phase is also lengthened in the

G. Meynet: Wolf-Rayet Stars in Starbursts

1.2

1.4 I

I • ' ' I ' ' ' I ' • '

I

Z=0.020 STAN. MASS LOSS+

I

293

1.6 I

I

I

I

I

1.8 I

2

2.2

1

Z=0.020 HIGH MASS LOSS

CD

O

o WNL I

,11 (20 25MO I

1.2

I

1 I

1.4

II

I

40Me 60Me 85Me 120Mo I

I

I

1.6

I

I

I

I

I

I

I

1.8

I

2

I

1 I

I

I I I .

20 25Mo 40Mo 60M e 85MO120M0 , • I . . . I . . . I . . . I . . .

2.2

M ini /M o FIGURE 1. Age as a function of the initial mass (in units of solar masses) when the star leaves the OV MK type, enters into the WNL, WNE and WC/WO stage. The upper line indicates the final age of the stars which have gone through a WR phase. On the left, the models of Schaller et al. (1992) have been used; on the right, the models by Meynet et al. (1994).

high mass loss rate models. This, of course, will change the way different WR subtypes are distributed at a given age after a starburst. • As a last effect, let us mention the fact that an increase in the mass loss rate produces a decrease in the minimum initial mass of single stars which can become WR. In Figure 2, the time evolution of the WR population for an instantaneous burst at time t — 0 is represented for different initial metallicities and for the standard and the high mass loss rate cases. The IMF's slope is x = 2. The effects of metallicity are discussed in Cervino & Mas Hesse (1993, see also this volume) and Meynet (1994). Let us concentrate here on the main effects of a change in mass loss rates. • For a given metallicity, the ratio of WR to O-type stars reaches higher values when the mass loss rate is increased. As an example, one obtains that, at time t = 3.5 106 yr and for Z — 0.008, the relative number of WR to O-type stars, NWR/NO, has a value of 0.08 in the standard case and a value of 0.16 in the enhanced mass loss rate case. • At low metallicity, the appearance of WC and WNE stars is favoured by high mass loss rates. It is interesting to note that at Z = 0.008, when the mass loss rate is enhanced, the WC stars become the dominant WR subtype during most of the WR-rich phase, while the dominant subtype is the WNL type in the standard case. Thus the discovery of a significant number of WC stars in low-metallicity starbursts might be an indication of

G. Meynet: Wolf-Rayet Stars in Starbursts

294

OS

is

+ o

1.2 1 0.8 0.6 0.4 0.2 0 1 0.8 0.6 0.4 0.2 0 .15

Z=0.040 STANDARD MASS LOSS

2 = 0 . 0 4 0 HIGH MASS LOSS

WR|

WR/

u: 2=0.020 STANDARD MASS LOSS WR

7/

2=0.020 HIGH MASS LOSS WR

WC

WN

.?

I 2=0.008 STANDARD MASS LOSS

X

HIGH MASS LOSS "2=0.008 AWR

WNL

" 0 0.04 0.03 0.02 0.01 0

- • I

_L

• I -- • I

2=0.001 STANDARD MASS LOSS

T

Z=0.001 HIGH MASS LOSS VWR

WR=WNL

i

I

i Ji

I i i . I . i i I : : i I i km

6 8 2

6 8

Age (in million years) FIGURE 2. Relative number ratios of WR, WNL, WNE and WC stars to the sum of the numbers of WR and O-type stars, as a function of time for instantaneous coeval starbursts. In each case the mass loss rate and the initial metallicity are specified.

G. Meynet: Wolf-Rayet Stars in Starbursts

295

W

0 INST. BURST o — 1- x=l x=2

—>

•

----

"2

Const. SPR • • *

HIGH MASS LOSS

-3 -1.5

—1

-0.5

0.5

[0/H] FIGURE 3. Relative number of WNL to O-type stars as a function of the relative abundance of oxygen expressed in terms of [O/H]. Filled circles are observed values as given by Vacca & Conti (1992). The open triangle represents 30 Doradus (from Vacca 1991) and open boxes are observations from Vacca as reported by Conti (1994). The dotted lines represent the theoretical values obtained in the case of a constant star formation rate with an IMF slope x — 2. The continuous lines are the predictions for an instantaneous burst with an IMF slope x = 2 and x = 1 (see text).

strong stellar winds. Let us note however that, whatever the metallicity and the mass loss rate, the WNL stars are dominant at the start of the WR-rich phase (see Figure 2).

4. Comparison with the observations The relative number of WNL over O-type stars observed in various starbursts by Vacca & Conti (1992) and Vacca (1991) are given in Figure 3 and compared with the predictions of population synthesis models. These starburst regions may differ not only in metallicity but also in age, the shape and intensity of the starburst and by the other physical parameters briefly commented in §1. Thus one cannot expect that the observed points draw a well-behaved relation. But the observations should be contained between two extreme cases, i.e. the case of a constant star formation rate, which represents an equilibrium situation between the birth and death of massive stars, and the case of an instantaneous burst of star formation which, on the contrary, represents an extreme non-equilibrium situation. It is well known that the more intense and shorter in time the burst is, the higher the values reached by the ratio of W R to O-type stars are. In Figure 3, we have represented the maximum value reached by the number ratio of WNL to O-type stars before a time t = 3.5 106 yr. In the case of a metallicity Z = 0.040, the ratio increases monotonically with time. In that case, we have represented the value reached at a time t = 3 106 yr. The models with an enhanced mass loss rate have been used to draw the theoretical

296

G. Meynet: Wolf-Rayet Stars in Starbursts

lines in Figure 3. From this figure, one can see that most of the observed points are between the constant SFR line and the one obtained with x = 1. Thus, the high mass loss rate models, which reproduce the WR statistics in zones of constant star formation rate quite well (Maeder k Meynet 1994), are also able to account for the relative numbers of WNL to O-type stars observed in starburst galaxies. The difference between the "constant SFR" case and the instantaneous burst case increases when the metallicity decreases (Figure 3). At low Z, only a small range of initial mass stars go through a WR phase, thus the burst of WR formation after star formation is narrower in time. Since in that case, nearly all the WR appear at the same time, the contrast with a stationary situation will be extreme. When the WR arise from a larger initial mass range, as happens at higher Z, the WR burst is more extended in time, making the contrast with a constant star formation situation less pronounced. 5. Conclusion For a given metallicity and a given age of the burst, stronger mass loss rates have as a consequence to increase the number ratio of Wolf-Rayet to O-type stars. Only the high mass loss rate models predict the presence of a significant number of WNE and WC stars at low metallicity. The high mass loss rate models reproduce quite well the statistics of WR stars born in a starburst in the case of a flat IMF. It would be of great interest to obtain more information on the presence or absence of other WR subtypes such as WNE and WC stars in starbursts. In view of the results presented in Figure 2, these observations provide an interesting complement to the already existing data.

REFERENCES D. & SCHILD, H. 1989 Astron. Astrophys. 224, 73. CERVINO, M. & MAS-HESSE, J. M. 1993 Astron. Astrophys. In press. CHARBONNEL, C , MEYNET, G., MAEDER, A., SCHALLER, G. & SCHAERER, D. 1993 Astron. Astrophys. Suppl. 101, 415. CHIOSI, C. & MAEDER, A. 1986 Ann. Rev. Astron. Astrophys. 24, 329. CONTI, P. 1994 Space Science Reviews In press. DE JAGER, C , NIEUWENHUIJZEN, H. &; VAN DER HUCHT, K. A. 1988 Astron. Astrophys. Suppl. 72, 259. MAEDER, A. 1991 Astron. Astrophys. 242, 93. MAEDER, A. & MEYNET, G. 1994 Astron. Astrophys. In press. MAS-HESSE, J. M. & KUNTH, D. 1991 Astron. Astrophys. Suppl. 88, 399. MEYNET, G. 1994 Space Science Reviews In press. MEYNET, G., MAEDER, A., SCHALLER, G., SCHAERER, D. & CHARBONNEL, C. 1994 Astron. Astrophys. Suppl. 103, 97. SCHAERER, D., MEYNET, G., MAEDER, A. &; SCHALLER, G. 1993a Astron. Astrophys. Suppl. 98, 523. SCHAERER, D., CHARBONNEL, C , MEYNET, G., MAEDER, A. & SCHALLER, G. 1993b Astron. Astrophys. Suppl. 102, 339. SCHALLER, G., SCHAERER, D., MEYNET, G. & MAEDER, A. 1992 Astron. Astrophys. Suppl. 96, 269. SMITH, L., MEYNET, G.& MERMILLIOD, J.-C. 1994 Astron. Astrophys. In press. VACCA, W. 1991 Ph.D. Thesis, University of Colorado, Boulder . VACCA, W. & CONTI, P. 1992 Astrophys. J. 401, 543. ARNAULT, PH., KUNTH,

Optical and Ultraviolet Morphology of the Starburst Regions in Wolf-Rayet Galaxies By WILLIAM D. VACCAf Dept. of Astronomy, 601 Campbell Hall, Univ. of California, Berkeley, CA 94720, USA Wolf-Rayet galaxies are a subset of starburst galaxies whose integrated spectra reveal the presence of hundreds to thousands of Wolf-Rayet stars. These galaxies exhibit a number of other properties indicative of a large "starburst" population of young, hot, massive stars. We have obtained optical ground-based and ultraviolet HST images of several Wolf-Rayet galaxies and present examples of the spatial morphologies observed at these wavelengths. Large star-forming regions which appear to be single units in the optical are resolved into numerous compact bright knots in the ultraviolet HST images. These multiple starburst knots are typically less than 100 pc in size and too small and closely spaced to be detected individually in the ground-based optical images. Yet they contain large numbers of hot stars and are typically several times as luminous as 30 Doradus, the giant HII region in the LMC. The intense bursts of star formation in these knots probably began only a few Myr ago and lasted less than about 1 Myr. It is possible that these knots represent proto-globular clusters which were formed as the result of recent galaxy mergers and/or interactions.

1. Introduction Wolf-Rayet (W-R) galaxies are a subset of HII galaxies in whose integrated optical spectra a broad resolved Hell A4686 emission feature has been detected (Conti 1991 and references therein). The width of the feature implies that it has a stellar, rather than nebular, origin. The Hell A4686 line is one of the most prominent emission lines in the optical spectra of Galactic and LMC W-R stars; the presence of 102 — 105 W-R stars in W-R galaxies has been inferred from a comparison of the luminosity and equivalent width of this feature with those of the corresponding line in the spectra of Galactic and LMC W-R stars (Kunth k Sargent 1981; Kunth k Schild 1986; Vacca k Conti 1992). The estimated number of W-R stars in these galaxies is remarkable in view of the relatively small number of these stars found in our Galaxy and other nearby galaxies. The spectra of W-R galaxies also exhibit a number of other properties (such as extremely "blue" continua and strong narrow forbidden and Balmer recombination lines) characteristic of objects containing large populations of young hot stars. Furthermore, the forbidden line intensity ratios indicate that the excitation sources in these objects are primarily stellar (Vacca k Conti 1992). The similarity of the optical spectra to those of clasical HII regions gives rise to their classification as "HII galaxies". Estimates of the numbers of hot (OB) stars in these galaxies are generally in the range of 102 — 106. More remarkable, however, are the estimates of the W-R to O star number ratios in these objects, as these provide some information on the initial mass functions and the history of star formation in these galaxies. Values between about 0.02 and 1.0 can be found in the literature (e.g., Kunth k Sargent 1981; Armus, Heckman, k Miley 1988; Sargent k Filippenko 1991; Vacca k Conti 1992; Conti k Vacca 1994b). The ratio seems to be roughly correlated with the oxygen abundance and optical continuum luminosity of the galaxies (Kunth k Joubert 1985; Kunth k Schild 1986; Sargent k Filippenko 1991; Conti k Vacca 1994b). Given the generally low (less than solar) metallicities in these W-R galaxies, the measured number ratios are significantly higher than what would be t Hubble Fellow 297

298

Vacca: Morphology of Wolf-Rayet Galaxies

predicted from current stellar evolution models under the assumption of a constant (or "steady-state") star formation rate (e.g., Arnault, Kunth, & Schild 1989; Maeder 1991); only models which invoke instantaneous bursts of star formation are able to reproduce the observed number ratios. Because W-R stars are the highly evolved, short-lived (r < 106 yr) descendants of the most massive 0 stars, the large numbers of W-R stars present in these galaxies and the unusually large W-R to 0 star number ratios suggest that the star formation in these galaxies occurred in short bursts (< 106 yr in length), and that the time since the start of each burst must be less than 107 yr (Kunth
Vacca: Morphology of Wolf-Rayet Galaxies

299

from CTIO were obtained with the 0.9-m telescope through broad-band filters and have a scale of about 0'-'5 pixel" 1 . In all figures, North is at the top and East is to the left. 2.1. NGC 1741 An optical image of NGC 1741 taken through the Gunn r filter is shown in Figure la, in which the unusual morphology of the object is apparent. The separation between the two central star-forming regions is about 4", which corresponds to about 960 pc at the distance of this galaxy; a linear "arm" of bright clumps also extends about 12 — 15" to the east of the southern central star-forming region. The UV (2200 A) HST image of the central area is shown in Figure lb, where the two star-forming regions are clearly resolved into numerous bright point-like objects, or knots. These knots each have a luminosity at 2200 A between several and nearly 100 times that of 30 Doradus in the LMC. 2.2. Mrk 1236 Mrk 1236 is a satellite of the larger galaxy NGC 3023. A Gunn r band image of Mrk 1236 and NGC 3023 is presented in Figure 2a; Mrk 1236 is composed of the three easternmost clumps in this image. Again, the unusual morphology of the entire system should be noted. The three regions comprising Mrk 1236 have a total spatial extent of about 8", which corresponds to a separation of about 890 pc. The UV HST image of these three regions is shown in Figure 2b. Once again, the star-forming regions, which appear as single clumps in the optical, are resolved into numerous discrete bright knots. The brightest of these knots has a luminosity at 2200 A that is a few times that of 30 Doradus. While most of the knots appear to be point-like, at least one is spatially extended. 2.3. He 2-10 A V band image of He 2-10 is presented in Figure 3a. (The image is from the data presented by Corbin, Korista, & Vacca 1993.) He 2-10 consists of two star-forming regions, (A) (the central region) and (B) (the eastern region) separated by about 8", or about 350 pc. The spectral analysis presented by Vacca & Conti (1992) indicated that both regions contain numerous O and B stars. However, large numbers of W-R stars are apparently found only in region (A). (The optical spectrum of region (A) obtained by Vacca k Conti 1992 is presented in Conti 1994.) The HST image of He 2-10 is presented in Figure 3b. Here, region (B) is seen to be composed of numerous bright knots, scattered randomly throughout the optical region, similar to what is found in the HST images of NGC 1741 and Mrk 1236. Several point-like objects are located outside of both regions (A) and (B); none of these are detected in the optical image. In addition, the outer sections of region (A) appear to exhibit faint spiral structure, which is also not detected in the optical image. It is not yet clear whether this structure is real or an artifact, resulting from either the ripples in the extended wings of the point spread function or the deconvolution procedure. An enlargement of the UV HST image of the center of region (A), on a different intensity scale from that of Figure 3b, is shown in Figure 3c. Here, the core of region (A) is seen to be composed of 9 — 10 compact knots, each typically ~ 10 pc in diameter and separated from one another by 10 pc, which form an arc about 90 pc long. The entire arc is contained within the central 1.5 — 2" of region (A). This result is quite remarkable and completely unexpected, as optical images of region (A) reveal a smooth brightness profile with an unresolved point-source at the center (Corbin, Korista, & Vacca 1993; Sugai & Taniguchi 1992). Furthermore, the arc-like structure of region (A) is quite unlike the distribution of knots seen in region (B) and is very suggestive of a sequential star formation scenario. The mean luminosity of the knots at 2200 A is 2.5 x 1038 erg s" 1 A" 1 , about 6 times the luminosity of 30 Doradus.

Vacca: Morphology of Wolf-Rayei Galaxies

300

Using a simplified population synthesis model and assuming an evolutionary age of about 5 Myr after an instantaneous burst of star formation, Conti & Vacca (1994a) estimated the masses of the knots to be about 105 - 106 M 0 . The sizes and masses of these knots are typical of those globular clusters and it has been suggested that these knots of 0, B and W-R stars are extremely young globular clusters which may have formed as a result of the interaction between regions (A) and (B). It is interesting to note that the faint spiral arms seen in the outskirts of region (A) seem to start near the two ends of the arc of knots. Thus, the arc might constitute a nuclear "bar" connecting the spiral arms. Further discussion of the HST image of He 2-10 can be found in Conti k Vacca (1994a).

*

,4.

Fig. l(a) Gunn r band image of NGC 1741. The separation between the two central regions is 4"(« 960 pc).

Fig. 1(6) The deconvolved UV (2200 A) image of the two central starburst regions in NGC 1741, taken with the HST FOC.

* V* '

Fig. 2(a) Gunn r band image of Mrk 1236 and NGC 3023. Mrk 1236 consists of the 3 easternmost clumps; the brighter two are separated by about 3"(« 330 pc).

Fig. 2(6) The deconvolved UV image of the 3 starburst regions comprising Mrk 1236, taken with the HST FOC. This image is 11.5"(« 1.3kpc) on a side.

Vacca: Morphology of Wolf-Rayet Galaxies

Fig. 3(a) Ground-based V band image of He 2-10. Region (A) is in the center; region (B) is 8"(« 350 pc) to the east.

301

Fig. 3(6) The deconvolved UV image of He 210, taken with the HST FOC. Note the faint spiral structure in the outer regions of (A) and the knots outside both (A) and (B).

Fig. 3(c) Enlargement of the HST FOC image of the center of region (A) of He 2-10. The arc of knots is about 2"(« 90 pc) long.

3. Conclusions The high spatial resolution U V images of starburst and W-R galaxies obtained by HST reveal that the star-forming regions in these galaxies are actually composed of numerous discrete compact knots of hot stars. The sizes of these knots range from less than about 10 pc (unresolved even by HST) to about 100 pc. They are generally too small and closely spaced to be individually resolved in ground-based images. Typical luminosities of these knots at 2200 A range from a few times to more than 10 times that of 30 Doradus. In the case of He 2-10, the knots have been shown to have masses of 105 - 106 M© (Conti & Vacca 1994a). Based on these results, Conti & Vacca (1994a) have suggested that the star-forming knots seen in W-R and starburst galaxies are extremely young

302

Vacca: Morphology of Wolf-Rayet Galaxies

globular clusters, with ages less than about 107 yr. Given the tendency of such galaxies to have disturbed morphologies, we suggest that the starbursts in many W-R galaxies have been triggered by recent galaxy-galaxy interactions and/or merger events and that perhaps all starbursts triggered in this manner result in the formation of compact knots of hot massive stars; these knots may be proto-globular clusters. The results for He 2-10, along with the recent identification by Holtzmann et al. (1992) and Whitmore et al. (1993) of numerous blue point-like knots in HST images of two other galaxy merger remnants, tend to support the suggestion (Schweizer 1987; Ashman & Zepf 1992) that globular clusters are formed as the result of galaxy mergers. I would like to thank Todd Small for graciously providing me with several optical CCD images of W-R galaxies. I thank Stephane Chariot, Claus Leitherer, and Peter Conti for numerous useful discussions. Support for this work was provided by NASA through grant HF-1026.01-91A, awarded by the Space Telescope Science Institute which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under grant NAS5-26555.

REFERENCES T. M. & MILEY, G. K. 1988 Astrophys. J. Lett. 326, L45. D. & SCHILD, H. 1989 Astron. Astrophys. 224, 73. ASHMAN, K. M. &; ZEPF, S. E. 1992 Astrophys. J. 384, 50. CONTI, P. S. 1991 Astrophys. J. 377, 115. CONTI, P. S. & VACCA, W. D. 1994a Astrophys. J. Lett. In press. CONTI, P. S. & VACCA, W. D. 1994b In The Nearest Active Galaxies (ed. J. Beckman). Madrid. In press. COPETTI, M. V. F., PASTORIZA, M. G. &: DOTTORI, H. A. 1986 Astron. Astrophys. 156, 111. CORBIN, M. R., KORISTA, K. T. & VACCA, W. D. 1993 Astron. J. 105, 1313.

ARMUS,

L.,

HECKMAN,

ARNAULT, PH., KUNTH,

HOLTZMANN, J. A. ET AL. 1992 Astron. J. 103, 691.

H., FRITZE-V. ALVENSLEBEN, U., FRICKE, K. J. &; LOOSE, H.-H. 1992 Astron. Astrophys. 259, L73. KUNTH, D. & JOUBERT, M. 1985 Astron. Astrophys. 142, 411. KUNTH, D. & SARGENT, W. L. W. 1981 Astron. Astrophys. 101, L5. KUNTH, D. & SCHILD, H. 1986 Astron. Astrophys. 169, 71. MAEDER, A. 1991 Astron. Astrophys. 242, 93. MAS-HESSE, J. M. & KUNTH, D. 1991 Astron. Astrophys. Suppl. 88, 399. SARGENT, W. L. W. & FILIPPENKO, A. V. 1991 Astron. J. 102, 107. SCHWEIZER, F. 1987 In Nearly Normal Galaxies (ed. S. M. Faber), p. 18. Springer. KRUGER,

SUGAI, H. & TANIGUCHI, Y. 1992 Astron. J. 103, 1470.

W. D. & CONTI, P. S. 1992 Astrophys. J. 401, 543. WHITMORE, B. C , SCHWEIZER, F., LEITHERER, C , BORNE, K. & J. 106, 1354.

VACCA,

ROBERT,

C. 1993 Astron.

The Starburst Nucleus of M83f By SARA R. HEAP Code 681, NASA/Goddard Space Flight Center Greenbelt MD 20771, USA With its nearly face-on orientation and relatively close distance, the barred spiral galaxy M83=NGC 5236 is one of the best objects for studies of nuclear starbursts. High-resolution images taken with the Planetary Camera on HST combined with ultraviolet spectra obtained by the IUE reveal that the nuclear starburst is actually a collection of over twenty young star clusters, each comparable to the prototypical starburst 30 Doradus in the LMC.

1. Introduction It is known from observations at X-ray, ultraviolet, Ha, infrared, CO, and radio wavelengths, that the nucleus of the barred spiral galaxy M83 is the site of vigorous star formation (Trinchieri, Fabbiano & Paulumbo 1985; Bohlin et al. 1983; Pastoriza 1975; Gallais et al. 1991; Turner, Ho & Beck 1990; Handa et al. 1991). Its infrared luminosity as measured by IRAS is LJR — 4.3 x 109 LQ, about one-sixth that of the starburst galaxy, M82 (Telesco 1988). Figure 1 shows a schematic of the central star-forming region as compiled from ground-based observation (Gallais et a/.1991; Telesco 1988; Turner et al. 1987). The FIGURE 1. M83 schematic. nuclear region has a rather lumpy or "hot-spot" appearance: there is the conventional nucleus (1); and about 8" (160 pc) to the Southwest are some star-forming knots arranged in an arc. In the IR and radio, the nuclear region of M83 looks very different: only Knots 6 and 8 are visible. M83 makes a particularly good test case to investigate the nature of a nuclear starburst because of its relative closeness (4 Mpc) and its near face-on orientation. We have therefore obtained high-resolution images of M83 with the Planetary Camera on HST (Section 2) and have supplemented them with ultraviolet spectra obtained with IUE (Section 3). These observations reveal that the nuclear starburst is actually a collection of about twenty young star clusters with ages from 2 to 8 million yr, and absolute visual magnitudes ranging from —10 to —15. We conclude from these and other HST observations presented here by Vacca & Leitherer that the basic "building blocks" of starbursts are clusters on the order of 30 Doradus in the LMC. This is fortunate because 30 Dor is close enough for extremely detailed examination (c.f. other papers in this volume by Kennicutt; Malumuth & Heap; and Leitherer).

t Based on observations obtained with the NASA/ESA Hubble Space Telescope at the Space Telescope Science Institute, operated by AURA for NASA under contract NAS5-26555. 303

304

Heap: The Starburst Nucleus of M83

2. ifST/Planetary Camera Observations Figure 2 shows the nuclear region of M83 as imaged by the Planetary camera and restored with the Richardson-Lucy algorithm. Along the top are pictures in the wideband U (F336W), V (F555W), and I (F785LP) filters, and also the narrow-band [Nil] filter (F658N), which fortuitously includes emission from redshifted Ho. The V-band image is shown also at the bottom, enlarged by 4x. PC images have a limiting resolution of about 0-1, which means that structures in M83 as small as 2 pc are resolved. The conventional nucleus is the bright source in the Northeast quadrant of the main K-band picture. While it is the dominant source in the / band, it is barely visible in the U band, and it is totally invisible in the Ha image once the contribution of the continuum flux is subtracted out. As can be seen from the color-color plot shown as Figure 3 (left), at least part of the reason that the nucleus is so very red is that its emission is greatly diminished by dust extinction. However, since it is not associated with any Ha emission, it is also possible that it is intrinsically red, and does not contain any young, massive, ionizing stars. To the Southwest and West of the nucleus is the circumnuclear arc, which is now resolved by the PC images into a string of about twenty knots. Each of the knots has a measurable diameter of 0'-'2 (FWHM) or more. As shown in the color-magnitude diagram of Figure 3 (right), their colors and luminosities identify them as young star clusters, each as large and luminous as the central ionizing cluster of 30 Doradus in the Large Magellanic Cloud. These young star clusters are embedded along the rim of the Ha shell and are undoubtedly the ionizing source of the observed Ha emission. Presumably the stars were once the source of the ionized gas itself, which is the aggregate of winds and supernova remnants. Surprisingly, there is also an Ha knot, located 0'-'5 to the East of Cluster 6, that has no counterpart on the wide-band continuum images. Perhaps it is a young supernova remnant - but it is not SN 1968/, which exploded in one of the clusters that comprise Knot 4 (Wood & Andrews 1974). The PC images show the two major dust lanes associated with the spiral bar. Figure 2 includes only the westerly lane, which runs approximately North-South between knots 6 and 8. The PC images also reveal numerous dust capillaries, some of which are punctuated by a young star cluster. Their form is reminiscent of Bok's "elephant trunks" and may be sites of triggered star formation (c.f. Elmegreen 1992).

3. Ultraviolet spectra From IUE The nucleus of M83 is the brightest ultraviolet source of its class and is easily observed with IUE. We obtained far-UV spectra at several positions, as shown in Figure 4. Despite the fact that the IUE aperture (10" x 20") extends over such a large fraction of the nuclear region of M83, there is a clear progression in spectral characteristics, with the composite type running from early-B supergiant in the Southeast, to late-B supergiant to the Northwest. We have compared these observations to synthetic spectra calculated from Schaller et al.'s (1992) evolutionary tracks (Z = ZG) and Fanelli et al's (1988) IUE spectral library. The calculated spectra are similar to those described by Leitherer (this volume) and are not repeated here. Instantaneous star-formation is certainly an appropriate assumption for an individual star cluster in M83, but the IUE aperture includes many clusters, so we have to allow for a mixture of ages. Although our analysis is still in progress, it is already clear that the ultraviolet spectral characteristics confirm the young ages of the

Heap: The Starburst Nucleus of M83

305

FIGURE 2. PC images of the central region of M83. A square-root contrast was used in all images. The North vector is at the 11:30 position.

Heap: The Starburst Nucleus of M83

306

-Ib

on.

ri

. . | .

.HTrnri

rr-r

-14

-12

-10

~~~(

^ . = io5

"I

^cl=1°4 -

-8 -3

-

2

-

1

0

1

.

- 3 - 2 - 1 0

1

2

o

FIGURE 3. Results of Photometry: observed color-color diagram (left) and color- magnitude diagram of star clusters (right). The brighter clusters in both figures are labeled. The jagged line shows the calculated evolutionary track of a cluster from birth ("blue" color) to 10 million yr ("red" color). The cluster tracks are based on Schaller et al.'s (1992) evolutionary tracks (Z = Z©) and Kurucz' (1992) flux distributions, and they assume a Salpeter IMF with an upper mass limit of 120 MQ. Three different cluster masses are shown in the C-M diagram.

star clusters inferred from their unreddened colors and the general strength of the Ha emission. The strength of the SilV A1400 wind feature indicates that quite massive stars (M > 50 MQ) are present. The progression of spectral type from Southeast to Northwest indicates that the cluster ages increase from about 2 million yr in the Southeast clusters to about 8 million yr in the Northwest. The span in ages is consistent with star formation propagating at about the speed of sound. 4. Conclusions In summary, • The nuclear starburst in M83 is actually a collection of young star clusters on the scale of 30 Doradus. • The pattern of ages of the young clusters as suggested by the progression of their ultraviolet characteristics is consistent with propagating star formation. • The nucleus itself is not a star-forming region. • The Ha emission appears to be a super-bubble formed by the winds and supernova shells of massive stars. This report is based on a collaborative study at Goddard by Jarita Holbrook (USRA/ GSFC), Eliot Malumuth (CSC/GSFC), Dara Norman (ACC/GSFC), and myself during the summer of 1993. Support for this work was provided by NASA through the GHRS Investigation Definition Team. We thank Wayne Landsman for the use of his program, CLUSTMOD, which we have used to compute the evolutionary tracks shown in Figure 3.

Heap: The Starburst Nucleus of M83

307

M83 Nuclear Region 1

i

i

t

i

i

i

i

i

i

i

i

i

i/ijA/vN^

0

1200

1400

1600 Wavelength (A)

1800

2000

4. IVE spectrograms of M83. The offset of the entrance aperture is shown at right The zero-point is defined as the centroid of light measured by the IUE FES sensors. It lies approximately midway between the nucleus, source 1, and the circumnuclear source 4. FIGURE

308

Heap: The Starburst Nucleus of M83

REFERENCES BOHLIN, R., CORNETT, R. H., HILL, J. K.,

SMITH, A. M. & STECHER, T. P. 1983 Ap.

J.

Lett. 274, L53. ELMEGREEN, B. 1992 In Star Formation in Stellar Systems (ed. G. Tenorio-Tagle, M. Prieto &; F. Sanchez). Cambridge Univ. Press. FANELLI, M., O'CONNELL, R. W., BURSTEIN, D. & Wu, C.-C. 1988 Ap.J. Suppl. 82, 197. GALLAIS, P., ROUAN, D., LACOMBE, F., TIPHENE, D. &; VAUGLIN, I 1991 A. k A. 243,

309.

HANDA, T., NAKAI, N., SOFUE, Y., HAYASHI, M. & FUJIMOTO, M. 1990 P. A. S. J. 42, 1.

KuRUCZ, R. 1991 Private communication. KEEL, W. C 1984 Ap. J. 282,

75.

PASTORIZA, M. 1975 Ap. & Sp. Sc. 33, SCHALLER,

G.,

SCHAERER,

D.,

173.

MEYNET,

G. &

MAEDER,

TELESCO, C. M. 1988 In Ann. Rev. Ast. & Ap. 26,

A. 1992 A. & A. Supll. 96, 269.

343.

TRINCHIERI, G., FABBIANO, G., PAULUMBO, G. G. C. 1985 Ap. J. 290, TURNER, J., Ho, P. T. P. & BECK, S. C. 1987 Ap. J. 313, WOOD, R. & ANDREWS, P. 1974 M. N. R. A. S. 167,

13.

644.

96.

Spectrophotometry of Haro Starburst Galaxies 1 By SIMON STEEL , NIALL SMITH2, 1 3 LEO METCALFE AND BRIAN McBREEN 1 2 3

University College, Belfield, Dublin 4, Ireland

Regional Technical College, Rossa Ave., Cork, Ireland

Astrophysics Division, Space Science Department, ESTEC, 2200 AG Noordwijk ZH, The Netherlands

Photometric observations on a group of eleven Haro starburst galaxies were carried out with the 1.0-m JKT over several semesters giving details of galaxy morphology and colour. These photometric images were then used in spectroscopic observations with the 2.5-m INT to enable accurate slit placement across the star forming knots. High and low dispersion spectra of the HII star forming regions were obtained, allowing the calculation of element abundance, abundance gradients, ionised gas dynamics, and estimates on burst age and possible starburst cause.

1. Introduction 1.1. The sample The galaxies were chosen from Haro's (1956) list of 44 blue galaxies. They are characterised by having an ultra-violet excess spectrum and emission lines from hot gas. Most of the galaxies showed unusual or chaotic morphologies, with the blue emission emanating from compact, usually central, regions. Our more detailed observation revealed a diverse sample of morphological species, including two spirals, five nuclear ellipticals (Loose 1986), two cigar-shaped irregulars and two clumpy irregulars. Table 1 lists the name, type, absolute visual magnitude, heliocentric distance, actual diameter and integrated colour of each galaxy. A Hubble constant of 75 kms" 1 Mpc"1 is used throughout. 1.2. Observations and data reduction

Photometric observations were carried out on the 1.0-m JKT in February 1985 and January 1988 with 512x320 pixel CCD detectors. The galaxies were observed through B, V and I filters. Contour maps are given in Figure 1. Spectroscopic observations were performed using the 2.5-m INT in February 1992 with the Intermediate Dispersion Spectrograph and 1280x1100 pixel CCD detector. The CCD was windowed in the spatial direction to 400 pixels, equivalent to a slit length of 4 arcmin. A slit width of between 0.8 to 1.5 arcsec was used depending on seeing conditions. Exposure times were 2x30 minutes, some were reduced to 2x15 minutes due to time constraints. Data reduction was performed using NOAO IRAF software. Spectra were flux and wavelength calibrated and corrected for reddening using the H/?/Ha ratio, assuming Case B recombination, and taking into account underlying stellar absorption in the Balmer lines. Detailed description of data analysis and results are to appear in a subsequent paper (Steel et. al. 1994). 309

Steel et al.: Spectrophotometry of Haro Starburst Galaxies

310

H ar 0 1 C —->

1

m

> i— —i

Figure la. V-band CCD images of Haros 1, 3, 20, 21, 22 and 23 showing star forming knots and INT slit positions. Scale bar is 10 arcsec.

Steel et al.: Spectrophotometry of Haro Starburst Galaxies

Figure l b . V-band CCD images of Haros 30, 37, 38, 42 and 43.

311

312

Steel et al.: Spectrophotometry of Haw Starburst Galaxies

Mv

Galaxy

Haro 1 Haro 3 Haro 20 Haro 21 Haro 22 Haro 23 Haro 30 Haro 37 Haro 38 Haro 42 Haro 43

-20.6 -17.7 -17.0 -17.3 -16.4 -16.9 -18.6 -18.3 -15.1 -19.3 -16.3

K(Mpc) 49.1 13.1 24.5 21.0 19.5 17.7 62.8 55.4 12.7 60.0 25.1

d(kpc) 8.8 3.8 5.0 7.9 5.0 6.0 9.1 4.9 2.0

(B-V)

10.5 2.5

0.38 0.44 0.57 0.53 0.43 0.59 0.42 0.55 0.59 0.40 0.55

TABLE 1. The galaxy parameters.

2. Element abundances 2.1. Abundance calculation

Abundances were calculated with ionic flux ratios from Lequeux et al. (1979) using atomic data from Aller (1984). Because of the weakness or absence of the [OIII] A4363A line, electron temperatures were derived from the ratio ([OII]+[OIII])/H/? as described in Shaver (1983). Where the 4363A flux was estimated, the electron temperature calculated was in good agreement with the Shaver equation. Electron densities were obtained from the [SII] A6717/[SII] A6735A ratio (Aller 1984). Values of electron temperature and density are given in Table 2. Line fluxes were measured using the routine SPLOT with a Gaussian fit of the line, the base of which was determined by fitting a polynomial to the underlying continuum. Abundances for oxygen, nitrogen and neon were calculated (Table 2) but because of the absence of SIII lines no accurate sulphur abundance could be obtained (Garnett 1989). Figure 2 shows graphs of the element abundances for the sample of 28 star forming regions and compares them to a set of galactic sources and low-metallicity objects from other surveys (Lequeux et al. 1979; Peimbert et al. 1986; Melnick et al. 1992; Skillman & Kennicutt 1993). A plot of nitrogen against oxygen abundance (Figure 2a) shows a linear relation from chemically young systems such as I Zw 18 up to chemically mature objects in our own Galaxy. The Haro sample forms a broad range through "middleage" to "mature", but contains no severely metal-deficient objects. A plot of neon against oxygen (Figure 2b) only reveals a trend when objects from other samples are included. A proportional increase of neon with oxygen is expected due to the production of both elements occurring in similar types of stars (Walsh & Roy 1989). The ratio of N/O against oxygen abundance (Figure 2c) shows most regions falling within log(N/O) = —1.4 ± 0.3. A high value of N/O suggests the presence of primary nitrogen (Pagel 1987), most noticeable in spirals Haro 1 and 21 but also in cigar irregular Haro 38 and the elliptical Haro 37. Nitrogen is relatively deficient in Haro 22 knot A and Haro 43 knot B, although the signal-to-noise in the latter is poor. 3. Burst age An estimate of starburst age can be obtained from the equivalent width of the H/? emission line. Mas-Hesse & Kunth (1991) and Copetti, Pastoriza k Dottori (1986) calculated the time dependence of W(H/?) in models of star forming regions for various

Steel et al.: Spectrophotometry of Haro Starburst Galaxies

313

Figure 2b. 12+log(Ne/H) vs. 12+log(O/H)

Figure 2a. 12+log(N/H) vs. 12+log(O/H)

12+log(Ne/H)

12+log(N/H) 9.0

8.5 » Haro H galactic « low Z

» Htro a galactic • low Z

7.5 -

8.0

6.5 -

•

#B

7.0 -

5.5

T 7.0

1

r

T

r

T~—i

8.0

6.0

9.0

12+log(O/H)

i

7.1

i

i

8.0 12+log(O/H)

Figure 2c. log(N/O) v«. 12+log(0/H) log(N/O) * Hiro S galactic « low Z

-l.l -1.2 -1.3 -1.4 -1.5 -1.6 -1.7 -1.8

I 7.0

8.0 12+log(O/H)

9.0

9.0

314

Steel et al.: Spectrophotometry of Haro Starburst Galaxies

Galaxy

knot

Haro 1

N A B C D A

Te

Nc

6690 7100 8090 8090 7030 9650 9200 8950 8860 7250 7800 7250 9800 9800 8150 8680 7540 7200 10300 10600 9800 9750 9940 9700 10400 9500 10050 9850

144 140 69 215 67 198 100 100 100 69 100 69 223 264 100 100 100 96 89 61 91 177 127 259 88 71 259 95 -

O/H

N/H

N/O

Ne/H

8.91 7.89 -1.02 8.83 7.71 -1.12 8.67 7.49 -1.18 8.68 7.48 -1.20 _ 8.85 7.66 -1.19 Haro 3 7.61 8.41 7.14 -1.27 8.61 7.04 -1.57 7.85 B 7.53 C 8.49 7.19 -1.30 N Haro 20 8.00 8.52 7.20 -1.32 A 8.17 Haro 21 8.75 7.58 -1.17 . 8.63 7.46 -1.17 B Haro 22 8.41 6.67 7.82 -1.74 A 7.96 8.43 6.91 -1.52 B 8.41 6.78 C -1.63 7.64 Haro 23 8.60 7.35 A -1.25 8.51 7.24 8.04 -1.27 B Haro 30 8.71 7.47 -1.24 8.00 A Haro 37 8.77 7.61 -1.16 N A Haro 38 8.40 7.22 7.73 -1.18 B 8.45 7.39 -1.06 8.40 7.30 7.72 -1.10 C A Haro 42 8.44 6.99 -1.45 7.86 8.42 7.09 -1.33 B C 8.48 7.14 -1.55 8.45 6.90 8.03 -1.34 D A Haro 43 8.43 6.84 7.64 -1.59 B 8.38 6.63 -1.75 7.94 C 8.41 6.83 -1.58 7.74 Solar 8.93 8.00 -0.93 7.87 TABLE 2. Electron temperatures, densities and element abundances in the star forming regions identified in Figure 1. Abundance data are given as 12+log(Z/H) and log(N/O).

IMFs and metallicities. Table 3 shows W(H/9) for each star forming region in the eleven galaxies. Three knots, two in Haro 3 and one in Haro 43, show equivalent widths of greater than 100A, which predicts a starburst age of less than 5 Myr. However, because of the diverse morphologies present in the sample, and consequently the large variation in underlying continuum, a comparison of two line fluxes would be a better indication of age. The excitation parameter [OIII] A4949+5007/H/? shows the proportion of high-mass stars (having greater UV flux) with shorter lifetimes; a large value is indicative of young burst age. Table 2 also lists excitation parameters, with a value greater than 2 suggesting a burst age of less than 4 Myr for the observed metallicities (Copetti et al. 1986).

4. Cause of the starburst Apart from the spiral galaxies, there is no clear answer to the triggering mechanism for the bursts. All systems are isolated and although Haro 3 gives the appearance of an interacting system, this is not supported by redshift differences in the knots. The possibility of stochastic self-propagating star formation (Gerola 1978) whereby one burst, triggered by molecular cloud collisions or supernovae, causes a cascade triggering across the galaxy, could explain the chain of star forming knots in Haros 38, 42 and 43.

Steel et al.: Spectrophotometry of Haro Starburst Galaxies

Galaxy

Knot

EW(A)

Haro 1

N A B C D A B C N A B A B C A B A N A B C A B C D A B C

8.4 13.4 16.5 8.5 24.6 175.1 50.0 102.7 29.7 29.0 79.1 29.4 14.7 67.0 32.4 50.4 69.0 30.6 36.7 6.0 38.6 64.2 80.8 14.2 16.9 123.7 17.2 61.8

Haro 3

Haro 20 Haro 21 Haro 22

Haro 23 Haro 30 Haro 37 Haro 38

Haro 42

Haro 43

315

[oiiq/: 0.45 1.14 0.124 0.91 1.18 6.11 4.57 4.28 4.18 2.35 2.84 6.82 5.35 6.34 2.62 3.52 2.14 1.75 7.11 4.43 6.77 5.08 6.27 3.74 3.40 5.79 4.62 6.31

TABLE 3. H/7 equivalent widths and excitation parameters for the star forming knots.

5. Conclusions The sample encompasses a wide range of morphological species with the common property of an ongoing high rate of star formation. Element abundance places the spirals close to solar values, the ellipticals at Z©/3 in oxygen and Z©/6 in nitrogen, and the irregulars at Z©/3 in oxygen but down to ZQ/20 in nitrogen. Abundance variation within individual galaxies only occurs for the two spirals, where the nuclear metallicities are enhanced, and the large irregular Haro 42, where the central two knots exhibit a higher metallicity than the outer two. The linear alignment of star forming regions in Haro 38, 42 and 43 suggest the possibility of propagating star formation events. Four galaxies - Haro 3, 22, 42, 43 - possess star forming regions under 5 Myr old from W(H/?), but the excitation parameter OIII/H/? suggests that a significant number of the starbursts are younger than 4 Myr.

REFERENCES ALLER, L. H. 1984 Physics of Thermal Gaseous Nebulae. Reidel. C O P E T T I , M. V. F . , PASTORIZA, M. G. & DOTTORI, H. A. 1986 Astron. Astrophys.

G A R N E T T , D. R. 1989 Asttophys.

J. 345, 282.

GEROLA, H. & SEIDEN, P . E. 1978 Astrophys.

J. 223, 129.

HARO, G. 1956 Boh Obs. Tonnntzintla. y Tacubaya 14, 8.

156, 111.

316 LEQUEUX, J.,

Steel et al.: Spectrophoiometry of Haro Starburst Galaxies PEIMBERT, M.,

RAYO, J. F.,

SERRANO, A. & TORRES-PEIMBERT, S.

1979

Astron. Astrophys. 80, 155. LOOSE, H-H. & THUAN, T. X. 1986 Star Forming Dwarf Galaxies and Related Objects (ed. Kunth & Van), p. 465. Frontieres. MAS-HESSE, J. M. &; KUNTH, D. 1991 Astron. Astrophys. Supp. 88, 399. MELNICK, J., HEYDARI-MALAYERI, M. & LEISY, P. 1992 Astron. Astrophys. 253, 16. PAGEL, B. E. J. 1987 Starbursts and Galaxy Evolution (ed. Thuan, Montmerle & Van) p. 227. Frontieres. PEIMBERT, M., PENA, M. & TORRES-PEIMBERT, S. 1986 Astron. Astrophys. 158, 266. SHAVER, P. A., MCGEE, R. X., NEWTON, L. M., DANKS, A. C. & POTTASCH, S. R.

Mon. Not. R. Astr. Soc 204, 53. SKILLMAN, E. D. & KENNICUTT, R. C. 1993 Astrophys. J. 411, 655. STEEL, S. J., SMITH, N. J., METCALFE, L. & MCBREEN, B. 1994 In preparation. WALSH, J. R. & ROY, J-R. 1989 Mon. Not. R. Astr. Soc. 239, 297.

1983

Starbursts in the Irregular galaxy VV 523 By J. HECQUET1, R. AUGARDE2, G. COUPINOT 1 ANDM. AURIERE 1 Observatoire du Pic-du-Midi, U.R.A. 1281 du CNRS, F-65200 Bagneres de Bigorre, France 2

Observatoire de Marseille, 2 place Le Verrier, F-13248 Marseille cedex 4, France

1. Introduction VV 523 (= NGC 3991 - UGC 6933 - IRAS11549+3237) is a typical clumpy irregular galaxy. Its distance is 42.6 Mpc, its absolute blue magnitude is —20.2 and its U — B index —0.44. On CCD frames taken at the Bernard Lyot 2-m telescope of the Observatoire du Pic-du- Midi VV 523 shows a complex structure with three components of bright clumps scattered in a common envelope. The extension of the envelope is 80" x 2" within a surface brightness of 25 mag arcsec"2.

2. Observations Three filters (B, R and /) were used for the CCD observations. The seeing was about 1 arcsecond. A wavelet transform with a scale parameter suited to the size of the starforming cells was used to detect the clumps and to provide their size, location and colour indexes. Figure 1 shows the wavelet transform of the B image. The photometric results are in Table 1. The clumps have a mean size of 200 parsecs and the B — I indexes exhibit a sequence in colour which corresponds to their location in the galaxy. Long-slit spectra with a dispersion of 33 and 260 A mm" 1 were recorded with the Carelec spectrograph at the 1.93-m telescope of the Observatoire de Haute Provence. The spectra show strong emission lines typical of HII regions. Three regions of different intensities are detected and named A, B and C. In Table 1 we give the physical properties of each of the three regions.

3. Interpretation In order to understand the physical conditions of VV 523, we have compared the observational data with the predictions of evolutionary population synthesis models described by Mas-Hesse and Kunth (1991) and Cervirio and Mas-Hesse (1993, hereafter the MHK model). The low values observed for W(Rp) allowed us to choose an intantaneous star formation process rather than a constant formation rate and led to an IMF slope of the order of 2.5. The age of the bursts is between 5.0 and 6.5 Myr. Their luminosities in the different wavelengths lead to an estimate for the total burst region mass of 4 107 M© while the rotational mass of the galaxy is 2.7 1010Mo- Figure 1 shows a plot of the B — I index of the clumps along the galaxy superposed on a plot of the variation of B — I with age in the MHK model. Both plots are fitted so that a linear relation is found between distance and age which suggests a propagating star formation mode as proposed by Kunth et al (1988) to explain their observations of dwarf galaxies. Another explanation is a complex interaction between VV 523 and NGC 3994 or/and NGC 3995, two peculiar galaxies located at the same distance, arranged over an area of only 4' in diameter and which also present vigorous bursts of star formation. 317

318

Nb

Hecquet et a/.: VV 523 Starbursts X

Y

D"

B

B-R

B-I

W(EP) (A)

Region

3 4 5

- 11.0 0.0 6.3

-3.6 0.0 5.8

1.6 1.7 1.6

20.5 17.2 18.3

1.3 0.3 0.6

1.3 0.3 0.6

C C C

6 7 9 14

14.7 22.3 13.1 40.0

11.3 12.5 18.3 44.3

1.7 1.8 1.1 1.1

18.8 18.1 20.9 21.7

0.5 0.3 0.9 1.7

0.6 0.6 0.8 0.9

B B B B

18 20 22 23 24 25

55.0 54.8 66.9 73.6 84.2 83.9

63.4 70.9 76.2 84.1 99.3 104.7

2.5 1.9 1.9 1.8 1.9 2.0

19.1 19.8 18.6 18.8 20.1 19.7

1.0 0.6 0.8 0.8 1.5 0.6

0.9 1.4 1.2 1.1 1.1 0.7

A A A A A

Z/ZQ

(erg s" 1 )

36

3.3 10

40

0.33

20

1.9 10

40

0.28

13

2.1 10 4 0

0.42

A

TABLE 1. VV 523 clumps: Nb, X-Y positions in pixels with respect to the Nb 4, D (size) in arcsecond, B magnitude, B — R and B — I colour indexes; then for the three regions A, B and C the equivalent width of H/3, the absolute total luminosity in the H^ line, and the metallicity. Age in Myr 5.1 I

5.6 i

i

6.1

|

I

1

T

-

1 -

CO .

r •

,

,

i

,

3350

6700

Distance in parsec

FIGURE 1. Left: Extraction of the clumps by a wavelet transform. Right: B — I indexes. Continuous line = MHK model and points with error bars = observations.

REFERENCES CERVINO, M. & MAS-HESSE, J. M. 1993 Astron. Astrophys. In press. KUNTH, D., MAUROGORDATO, S. &; VIGROUX, L. 1988 Astron. Astrophys. 204, 10.

MAS-HESSE, J. M. & KUNTH, D. 1991 Astron. Astrophys. Suppl. 88, 399.

Long-slit spectroscopy of the central regions of starburst galaxies: Henize 2-10 and Markarian 52 ByHAJIME SUGAI1! AND YOSHIAKI TANIGUCHI2{ 1

Department of Astronomy, University of Tokyo, Bunkyo-ku, Tokyo 113, Japan

2

Kiso Observatory, Institute of Astronomy, University of Tokyo, Mitake-mura, Kiso-gun, Nagano 397-01, Japan

Long-slit spectroscopic observations have been made for two starburst galaxies, Henize 2-10 and Markarian 52, in order to discuss physical conditions of the starburst regions in the galactic nuclei. We obtained the spatial variations of physical conditions of the ionized gas, such as Ha surface brightness, electron density, and ionization/excitation conditions. In particular we found an anti-correlation in the spatial variation between the [Olll]/H/? ratio and the [Sll]/Ha and [Ol]/Ha ratios in each galaxy. This anti-correlation suggests that the nuclear starburst region consists of a central massive-star cluster and a single envelope of ionized gas.

1. Which structure do starburst regions have? One of the most important issues on starburst galaxies has been the typical size of the distribution of ionizing stars. Although it is well known that ionized gas in starburst galaxies extends over a few hundred pc to one kpc, the distribution of ionizing stars themselves is still a controversial issue (e.g. Sugai & Taniguchi 1992; Puxley et a/. 1990). There are two models for the structure of starburst regions. In one model, ionizing stars exist in a central cluster with its surrounding ionized nebula (Figure la). This structure is just like that of one giant/supergiant HII region, such as NGC 604 in M33 and 30 Doradus in LMC. In the other model, on the other hand, ionizing stars are distributed over the whole starburst region, with individual ionized regions (Figure lb).

2. Why is this issue important? This issue of the distribution of ionizing stars is important for the following reasons. (1) It is essential to discussions about where and how starburst phenomena begin and evolve. (2) It is also important in terms of the possible connection with active galactic nuclei. Norman k, Scoville (1988) proposed a model for active galactic nuclei in which a massive central star cluster builds up and feeds a central black hole as a result of mass loss during post-main-sequence stellar evolution. In their model, the assumed starting point is a central cluster of ~ 10 pc, which should be observationally demonstrated. This is particularly important for massive starbursts. (3) When we determine the metallicity of the ionized gas, the temperature of ionizing stars, and the ionization parameter, in starburst galaxies from line flux ratios, the obtained values depend on the distribution of ionizing stars. Therefore, for an accurate determination of these parameters, it is necessary to know the distribution of ionizing stars. t Present address: National Astronomical Observatory of Japan, Mitaka, Tokyo 181, Japan % Present address: Astronomical Institute, Tohoku University, Aoba-ku, Sendai, Miyagi 980, Japan 319

320

Sugai L Taniguchi: Long-slit spectroscopy of He 2-10 and Mrk 52

(a) Ionized region

(b) Ionized region

FIGURE 1. Two models for a starburst region, (a) A central cluster of ionizing stars and its surrounding ionized nebula. (6) Ionizing stars distributed over the whole starburst region, with individual ionized regions.

He 2-10

a(1950) 5(1950) Morphology Radial velocity Distance 1" mB

Hi-flux Sl2,25,60,100^im

S6cm

h

m

08 34 07-l -26°14'04" dwarf 450 km s" 1 6 Mpc 29 pc 11.80 mag 0.1 Jy 1 .1, 6.5, 24, 27 Jy 55 mJy

Mrk 52 12 h 23 m 09-3 +00°50'53" SBbc 2150 km s" 1 43 Mpc 210 pc 13.05 mag 3.43 Jy km s" 1 -•, 1.09, 4.58, 5.91 Jy 2.8 mJy

TABLE 1. Table 1. Basic data of He 2-10 and Mrk 52.

3. How do we solve this issue? In order to solve this issue, we investigate the spatial variations of physical conditions of the ionized gas in starburst regions, such as Ha surface brightness, electron density, and ionization/excitation conditions. Then we compare these spatial vatiations with those for a giant/supergiant HII region in order to investigate whether or not they have similarities to each other. The merit of this method is to clarify both the structure of the ionized gas and, at the same time, the distribution of ionizing stars.

4. Long-slit spectroscopic observations In order to investigate the structure of starburst regions, we made long-slit spectroscopic observations for two starburst galaxies, Henize 2-10 and Markarian 52 (hereafter He 2-10 and Mrk 52, respectively). Table 1 shows the basic data of these two starburst galaxies. In the next section, we discuss these galaxies separately because He 2-10 is a dwarf starburst galaxy while Mrk 52 is a spiral galaxy with starburst. We used the University of Hawaii 88-inch telescope with the Faint Object Spectrograph. The spectral resolution was 150 km s" 1 and 120 km s" 1 for the blue and red spectral regions, respectively. The spatial resolution was l"4. The position angle of the slit was set to 90 deg. At this setting the eastern condensation as well as the central condensation of active star-forming regions in He 2-10 were observed simultaneously. In the case of Mrk 52, the slit was almost along the major axis of the galaxy.

Sugai k, Taniguchi: Long-slit spectroscopy of He 2-10 and Mrk 52

321

5. Results and discussion 5.1. He 2-10 We found that the starburst region in He 2-10 has similarities with a supergiant HII region in the following points. (1) We can estimate the total Ha luminosity from the starburst region by assuming spherical symmetry. The estimated Ha luminosity is similar to that of a supergiant HII region in a disk, such as NGC 5461 in M101. (2) We found that the spatial variation of Ha emission measure in He 2-10 is similar, in relative shape, to that of 30 Doradus, a supergiant HII region in LMC (cf. Figure 2). (3) We can express this similarity in terms of the radial variation of the root-mean-square electron density. We can derive the radial variation of the rms electron density, Nrms, from the Ha emission measure (EM) profile, based on a spherically symmetric shell model: K

EM(x) = 2 /

dr

'

where x is the projected radial coordinate in the plane of the sky, r is the true radial coordinate within a starburst region of radius R, and the constant corrects for electrons from helium. The Ha emission measure profile is derived directly from the Ha surface brightness profile. The Abelian integral is solved numerically by dividing the region into shells and integrating inward from the outer shell and the result is shown in Figure 3. We find that the radial variation of the rms electron density of He 2-10 is indistinguishable, in relative shape, from those of giant/supergiant HII regions. (4) We found that the spatial variation of [Om]/H/? ratio is anti-correlated with those of the [Sll]/Ha and [Ol]/Ha ratios (Figure 4). Figure 5 shows this anti-correlation on an excitation diagram. This kind of excitation diagram is useful for investigating the nature of ionizing sources as well as excitation conditions of the emitting regions. Results of model HII regions calculated by Evans & Dopita (1985) in the case of the solar abundances are also shown in the figuref. It is remarkable that the data points of the emission-line ratios of He 2-10 are well aligned on one locus of a constant temperature of ionizing stars. It looks possible to conclude that the temperature of typical ionizing stars for this starburst region is constant at about 41,000 K, at least over 400 pc. However, this is an unlikely conclusion because, if it is true, the starburst phenomena would have taken place very coherently at least over 400 pc. It is more reasonable that the starburst region in He 2-10 is a supergiant HII region. In other words, the ionizing stars are concentrated in a central cluster and the temperature represents that of the central massive-star cluster. The spatial variation of the line ratios is primarily due to the variation of the mean ionization parameter, decreasing outwards. All the above-mentioned facts suggest that the starburst region of this dwarf galaxy is a supergiant HII region with a central cluster of ionizing stars. 5.2. Mrk 52 Although the starburst region of Mrk 52 has a much larger Ha luminosity, it has similar properties to the starburst region of He 2-10 in several points. (1) The values at the center of the Ha emission measure, rms electron density, electron density from the [Sn] line ratio, filling factor, and Ha equivalent width are similar in Mrk 52 to those in He 2-10. (2) If we normalize the radial scale by the full width of half maximum dimension of each starburst region, the radial variations of these values are also similar to those in He 2-10 (Figure 6). (3) Moreover, the spatial variations of emission-line ratios are f Here we should be careful in the comparison between the observations and the model results because the model results are averaged values over the entire Hll regions.

Sugai & Taniguchi: Long-slit spectroscopy of He 2-10 and Mrk 52

322

W

6

w

Brigh

c

d X

4

x

0)

o o

2. Spatial variation of the Ha surface brightness (solid lines) and the Ha/H/9 ratio (dashed lines), which was used to make a correction for the extinction, for He 2-10. West is to the right and east is to the left. FIGURE

X

1000 e E

100 =- V He 2-10

- \\

(0

E

a x

10

:NGC6Q4->-3°. D o r I

0

FIGURE

I I I I I I M

IP--L I I

100 200 Radius (pc)

300

8 X

2 1.5 1 .5 0 .02 r(b) .015 .01 .005

.9 = (c)

.3 .2 .1 0 .15 .1 .05 0

1111

-10

I IIIIIIIII I

-5 0 5 Space (arcsec)

3. (left). Radial variation of the root-mean-square electron density of He 2-10, with those for giant/supergiant Hll regions from Kennicutt (1984).

FIGURE 4. (right). Spatial variation of the emission-line ratios for He 2-10: (a) [Olll] A5007/H/?, (6) [Ol] A6300/Ha, (c) [Nil] A6583/Ha, and (d) [Sll] A6731/Ha. The ratio of [Sll] A6731/Ha is corrected for the effect of deexcitation.

Sugai & Taniguchi: Long-slit spectroscopy of He 2-10 and Mrk 52 i

1

1 9

]LO

1.0

1

i

1

323 1

56,000 K

cm sec" 1

"""--^ io 8 cm sec" 1 ^-^5.000 K to 00 <

O O

-^

/ ' '

-

40,000 K _ / ~~ f

I—

o o <

N

N

0.5

0.0

'

— ^

\

7 1 ^ ^ Q ( H ) = 10 cm sec"

38,500 K

-

"T K

Jr.

\

37,000 —— -0.5 -

\ \

—

y

\ 1

T. = 41,000 K

/ \ ,'

\

7 /

> \

-1.0 —

\

-

\i -1.5 - .

i

-3.0

i '{

-2.5

-2.0

1

1

i

-1.5

-1.0

-0.5

r 0.0

LOG([Oi]A6300/HaA6563)

FIGURE 5. Spatial variation of the [Olll] A5007/H/? and [Ol] A6300/Ha line ratios for He 2-10 (filled circles) and Mrk 52 (open circles). The line ratios obtained by spatially integrating the line fluxes after the correction for the extinction are also shown by crosses. The five dashed curves are the model results calculated by Evans and Dopita (1985) for an Hll region for constant ionizing-star temperatures T, of 37000, 38500, 40000, 45000, and 56000 K, respectively. The curves of a constant ionization parameter are also shown in dot-dash curves.

(a) He 2-10 'E O

E

•^

300 200 100 0

(b) Mrk 52 300 200 100

-MM

0

1500

1500

1000

1000 500

500 I I I

K

oL

o

.06

.06

U_

CO

.04

.04

c

.02 0

K

MM

.02 • M M

012345 Radius (")

0

MM

012345 Radius (")

FIGURE 6. Radial variations of the root-mean-square electron density, the electron density, and the filling factor for (a) He 2-10 and (6) Mrk 52.

324

Sugai & Taniguchi: Long-slit spectroscopy of He 2-10 and Mrk 52

also similar to those for He 2-10, i. e. the spatial variation of the [Om]/H/? ratio is anti-correlated with those of the [Sn]/Ha and [0 i]/Ha ratios. And this can be easily explained by a constant ionizing-star temperature, about 39,000 K, suggesting a central cluster of ionizing stars (Figure 5). From these similarities of the starburst region of Mrk 52 to that of He 2-10, it is suggested that the starburst region of Mrk 52 has also a structure of a central cluster surrounded by a single envelope of ionized gas. 6. Conclusions Using a spherically symmetric shell model for the starburst regions of He 2-10 and Mrk 52, we found that the rms electron density, the electron density, and the filling factor decrease outwards in these starburst regions. We also found an anti-correlation in the spatial variation between the [Om]/H/? ratio and the [Sll]/Ha and [Ol]/Ha ratios in each galaxy. This anti-correlation is reasonably explained by a the nuclear starburst region that consists of a central massive-star cluster and a single envelope of ionized gas, just as is observed in many giant/supergiant HII regions in galactic disks. The spatial variation of the line ratios is primarily due to the variation of the mean ionization parameter, which decreases outwards by a factor of 3-5.

REFERENCES EVANS, I. N. & DOPITA, M. A. 1985 Ap. J. Suppl. 58, 125. KENNICUTT, R. C. 1984 Ap. J. 287, 116. NORMAN, C. & SCOVILLE, N. 1988 Ap. J. 332, 124. PUXLEY, P. J., HAWARDEN, T. G. & MOUNTAIN, C. M. 1990 Ap. J. 364, 77. SUGAI, H. & TANIGUCHI, Y. 1992 A. J. 103, 1470.

Star Formation in Active Galactic Nuclei: the Cases of NGC 5135, NGC 6221 and NGC 7130 1 By H. R. SCHMITT 1 ,2 T. STORCHI-BERGMANN , 3 A. S. WILSON AND J. A. BALDWIN

^epartamento de Astronomia, IF-UFRGS.CP 15051, CEP91501-970, Porto Alegre, RS, Brasil 2

Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA 3

Cerro Tololo Inter-American Observatory, NOAO, Casilla 603, La Serena, Chile

We present long-slit spectroscopy of the composite Seyfert 2 starburst nuclei of the galaxies NGC 5135, NGC 6221 and NGC 7130 (IC 5135). Extended emission is detected in all three galaxies, reaching about 1 kpc from the nuclei. We study the spatial variation of the stellar population and emitting gas properties over the central regions. We compare our observed emission-line ratios with those obtained using composite photoionization models, which include ionization by a power law and hot stars, to find the gaseous abundance and the HII region parameters.

Several Seyfert galaxies exhibit the observational characteristics of vigorous star formation, either around the nucleus or in the galaxy disk, which is evident from extranuclear low excitation optical emission, diffuse non-thermal radio emission and very steep infrared spectra between 25 and 60 //m (Wilson 1988). To investigate the connection of star formation and nuclear activity we have selected three Seyfert 2 galaxies with composite (Seyfert + HII) spectra - NGC 5135, NGC 6221 and NGC 7130. The nuclear spectrum of NGC 6221 presents a low excitation kinematical component and another, which is blueshifted, with high excitation (Pence k Blackman 1984). NGC 7130 also has 2 kinematical components, and the emission line ratios vary from values typical of active nuclei at the nucleus, to values characteristic of HII regions further out (Shields k Filippenko 1990). Ulvestad k Wilson (1989) discovered that the radio emission (6 and 20 cm) of NGC 5135 is extended by 9" (3.6 kpc) in the NE direction, an extension also observed by Haniff, Wilson k Ward (1988) in [OIII] and Ha images. In order to study the spatial variations of the stellar population and gaseous properties in the nuclear region, we have obtained long-slit spectra for these galaxies, with a CCD detector on the Ritchey-Chretien spectrograph of the CTIO 4-m telescope. The spectral range covered was AA3400-7500 A at 8 A resolution. The spectra were extracted binning together two pixels (2"xl''8) at the nucleus and circumnuclear regions.

1. Stellar population and emission lines We investigated the stellar population by measuring the equivalent widths (WA) of the following absorption features: K(CaII, spectral window AA3908-3952 A), G band (AA4282-4318 A) and Mgl (AA 5156-5196 A). We have also measured the ratio between the continuum fluxes at A5870 A and A4020 A. The stellar populations of these galaxies present similar behavior: at the nucleus, the stellar population is mainly composed of young and intermediate age stars, the contribution of which decreases outwards. The continuum ratio A5870 A/A4020 A also presents a gradient consistent with a bluer population towards the center. A comparison with stellar population templates from Bica (1988) indicates that the continua are reddened by E(B — V) ss 0.3 in the three galaxies. 325

326

Schmitt et al: Star Formation in Active Galactic Nuclei 1

. M i l l

IUi: I I I—HH—H-b-H

, r

:

,

I

I

.

,

.

,

1

1

,

1

-

0

5 10 15 0 0.5 1 1.5 0 0.2 0.4 0.6 [OIII]/H/S [NII]/Ha [SII]/Ha FIGURE 1. Models for NGC 5135 (top) and NGC 7130 (bottom). Each line corresponds to the From the emission line ratios [OIII]/H/? and [OIII]/[OII] (corrected for reddening) we found that for NGC 5135 the location with highest excitation is not the nucleus, but at 500 pc NE of it; the excitation decreases further out. For NGC 7130 the excitation is highest at the nucleus and in the inner 1.2 kpc, decreasing for larger distances. For NGC 6221 the excitation is lower than that of the other two galaxies, and also decreases outwards. The emission line ratios [NII]/Ha and [SII]/Ha support this interpretation. 2. Composite models We used CLOUDY (Ferland 1991) to calculate models which incorporate photoionization both by an active nucleus and young stars, since models considering only photoionization of either source alone could not reproduce the observed emission line ratios. We used a gas density of 300 cm" 3 and a grain composition given by Cowie & Songaila (1986). Figure 1 shows the results for NGC 5135 and NGC 7130. The models which better reproduced the observed values were those in which the stars have a temperature of 35000 K. The gas has nitrogen abundance 2 times solar and 1.5 solar for the other elements. For NGC 7130 the corresponding values are 1.7 and 1.2, respectively. Although the [OIII]/H/?x[OII]/[OIII] diagram can be reproduced by models with an ionization parameter \ogU = —2.5 for the stars and ionization parameters for the power law in the range -3.5 < logU < -2.0, not all [NII]/Ha and [SII]/Ha ratios are reproduced. We interpret this as being due to the possible presence of shocks enhancing both ratios.

REFERENCES BlCA, E. 1988 Astron. Astrophys. 195, 76. COWIE, L. L. &; SONGAILA, A. 1986 Ann. Rev, Astron. Astrophys. bf 24, 499. FERLAND, G. 1991 OSU Internal Report N°91-01. HANIFF, C. A., WILSON, A. S. & WARD, M. J. 1988 Astrophys. J. 334, 104. PENCE, W. D. & BLACKMAN, C. P. 1984 Mon. Not. R. Astron. Soc. 207, 9. SHIELDS, J. C. & FILIPPENKO, A. V. 1990 Astron. J. 100, 1034. ULVESTAD, J. S. &: WILSON, A. S. 1989 Astrophys. J. 343, 659. WILSON, A. S. 1988 Astron. Astrophys. 206, 41.

Metallicity Effects on Starburst By M. CERVINO AND J. M. MAS-HESSE Laboratorio de Astrofi'sica Espacial y Fisica Fundamental, POB 50727 E-28080 Madrid, Spain Evolutionary stellar population synthesis models have been performed for several metallicities and two extreme star formation rates, instantaneous (IB) and extended bursts (EB). We discuss the dependence on metallicity of the population of Wolf-Rayet (WR) and Red Supergiant (RSG) stars. We show that both populations become more abundant for higher metallicities. We also show the effects of metallicity on the effective temperature and H/3 equivalent width. These effects are independent of the IMF slope and can account, at least in part, for the higher values of Ten and W(R0) systematically found in low-metallicity star formation episodes. A more complete study can be found in Cerviiio & Mas-Hesse (1994).

1. Wolf-Rayet population In Figure la we compare observational values of WR bump over H/J ratio taken from Kunth k Joubert (1985), Kunth k Schild (1986) and Vacca k Conti (1992) with the model predictions. Ages have been estimated by using the W(R/3) computed for a Salpeter IMF slope (Figure lb). The average WR bump over H/3 intensities fall within the range predicted by our models for an IB regime and cannot be reproduced assuming an extended one. We can therefore reject the possibility of having extended star formation episodes in the majority of the cases, and can also constrain the age of the episodes to a short range between 3 and 5 Myr after the onset of the burst. 2. H/3 equivalent width and effective temperature We show in Figure lb the dependence of W(Rft) and Tefr on metallicity as a function of age, assuming in all the cases an EB regime and a Salpeter IMF slope. From these plots it becomes evident that the metallicity of the gas from which the stars formed strongly influences the properties of the ionizing continuum. The metallicity effect could therefore explain, at least in part, the observational trend towards starbursts with higher Teff and larger VF(H/3) at lower metallicity, without having to invoke a metallicity dependence for the IMF. 3. V — K color index and the RSG population The relative population of RSG decreases in general with metallicity. This effect has important influence on the V — K color index. Figure lc shows how this parameter evolves for a Salpeter IMF slope and IB regime at different metallicities. We have compared observational V — K values taken from Thuan (1983) with the model predictions. Ages have been estimated by using W{Q.0) as indicator. We can see that our model's predictions agree with the measured V — K values for low-metallicity galaxies (triangles), whereas observational values are clearly higher for galaxies at higher metallicities. These results indicate that in these galaxies the contribution of an underlying old stellar population is very important, dominating the infrared part of the spectrum. We have marked with filled symbols in the plots the values corresponding to I Zw 18 (very low metallicity) and NGC 4214 (intermediate metallicity), which will be analyzed in detail in Kunth k Mas-Hesse (1994). A detailed fitting of NGC 4214 327

328

Cervifio & Mas-Hesse: Metallicity Effects

a

1 .1 .01

33

1

B

.1 r .01 r

<

r1

•

1

,.

i • • <

i . . ,

—T5

1 1 1

1

1 i , i , | • i , , | 1

"4

r

zo

/a=3

;

i

.

.

•

,

i

.

i

r

.

.

i

,

1

.

i • 1 1 1 1 1

IB

r f 1 1

2

.1

; *•.

\

ZO/2.5 •

•

I

I

I

—

! i .

I,

, , , {, 3

Fig a ZO/5 \

. . 1 . . .

4 5 t(Myr)

\

w 10

:

Z.75

V --* 4

T

••}

11 • 11

1 [-.

••••4

..•••

1 r .1 r

\

"I

i.

1 | , 1

hH

100

T

EB

•

300

r <*

f 1 1 , , ,

.01

> 1

\

\

\ \

V

V5

3

..........."Ti; 10 15 20

\ \

1

10

Fig b

•

2.75 *^

-

3 -

- 40000

_

-

V*

2 -

35000

'V

H CO

\

;

\

1 -

\

- 30000

_

\

-

\

\

•

V.

0 -

, , , 1 , , , 1 , , , , 1I I I . I

- 25000

-

,1,,"

7.5 8 8.5 9 9.5 12 + log [O/H] FIGURE 1. la: Ratio of WR bump over H/? intensities for three metallicities and to extreme IMF slopes, l b : W(E/3) and Tes for different metallicities and ages and a Salpeter IMF slope. lc: Evolution of the V — K color index for three metallicities compared to data observational (Z0/2O: triangles; Z Q / 5 : squares; Z©/2.5: hexagons).

has indeed shown that the star formation episode alone can not reproduce its observed optical continuum, and that a significant contribution from older stars is needed.

REFERENCES CERVINO, M. & MAS-HESSE J. M. 1994 A. k A. In press. KUNTH, D. & JOUBERT, M. 1985 A. & A. 142, 411. KUNTH, D. & SCHILD, H. 1986 A. & A. 169,

71.

KUNTH, D. &c MAS-HESSE, J. M. 1994 In preparation. THUAN, T. X. 1983 Ap. J. 268,

667.

VACCA, W. D. & CONTI, P. S. 1992 Ap. J. 401,

543.

From 30 Doradus to QSOs ByROBERTO TERLEVICH Royal Greenwich Observatory, Madingley Road, Cambridge, CB3 OEZ, UK I discuss some of the predictions of the Starburst model for AGNs, in particular the relation between observational parameters like the average blue luminosity, the amplitude of variability, the frequency of slow peaks in the light curve, and the time-averaged equivalent width of H/7. The number of slow peaks or SN events in the light curve of low-luminosity AGNs is uniquely related to the nuclear luminosity. An AGN with Ma(,niin) ~ —21.5 t produces 1 slow peak (or SN) per year. This result is independent of the initial mass function, age and/or total mass of the cluster. The time-averaged equivalent width of H/3 is related to the total energy of the SN, almost independently of the initial mass function, age and/or total mass of the cluster and of the assumed cosmology; the observed constancy of the value of the equivalent width of H/? in AGNs is a direct consequence of the universal value of the energy released in a SN explosion. The long term variability of AGNs as a function of their luminosity has a peak at a luminosity similar to the maximum luminosity of cSNR (i.e. MB ~ —20). AGNs with both larger and smaller luminosity than MB ~ —20 should be less variable than those with MB ~ —20. I estimate the optical size of starburst-cores based on recent HST results on the light distribution in 30 Doradus and discuss recent arguments on the energetic of starbursts showing that there is no energetic problem in the Starburst model for AGNs.

1. Introduction In what has now become known as the Starburst model, the observed variability of radio-quiet Active Galactic Nuclei (AGNs), i.e. Seyfert galaxies and most optically selected quasars, is postulated to be produced by the supernova (SN) and compact supernova remnant (cSNR) activity resulting from the evolution of a metal-rich massive stellar cluster, the product of a starburst in the core of an early-type galaxy (Terlevich et aV. 1987, 1992). The multifrequency spectrum of these radio-quiet AGNs can be reproduced by the combined contribution of stars, SNe, cSNRs and dust present in a stellar cluster 10 to 60 Myr old (Terlevich 1990a). Most of the optical/UV (the big UV bump) and bolometric nuclear luminosity is provided by the young stars, while the cSNRs are responsible for the observed nuclear variability and characteristic broad line spectrum. The broad lines observed in type 1 Seyferts and QSO are produced in cSNRs by the interaction of the ejecta of SNe with a high-density circumstellar medium. That broad lines typical of AGNs preferentially detected in the cores of big spheroids are assumed to be related to the high metallicity of the core stars. This high metallicity leads to large mass loss rates that in the high-pressure ISM typical of the Starburst produces high circumstellar densities around the cSNR progenitors. The intrinsic parameters of the Broad Line Region (BLR) can be obtained from theoretical models of cSNRs (Terlevich et ai. 1992), while the observed delays of the response of the broad emission lines to the variations of the continuum are well explained by thermal instabilities during shell formation in cSNRs (Tenorio-Tagle et ai. 1992). The observational data that more strongly support the cSNR origin of the BLR and its variability, are the properties of the "Seyfert 1-like" supernovae SN 1987F and SN 19881 (Filippenko 1989), SN 1983K (Terlevich & Melnick t Unless explicitly mentioned, Ho = 50 kms"1 Mpc"1 and qo = 0.5 329

330

Terlevich: From SO Doradus to QSOs

1988) and SN 1988Z (Stathakis k Sadler 1991). Terlevich and Boyle (1993) have explored the hypothesis that QSOs are the young cores of massive ellipticals forming at z > 2.0. They found that only a small fraction (~ 5%) of the total mass of elliptical galaxies, the core mass, is needed to participate in a burst to explain the observed luminosities and luminosity function of QSO at z > 2.0. I discuss Heckman's (1994) criticisms of this hypothesis in section 7. On the other hand, following the initial burst, models for the subsequent decline in the star formation rate in the core can also provide a good fit to the observed evolution of the QSO luminosity function. Recent work (Hamann k Ferland 1993) supports also the hypothesis that QSOs are associated with a metal-rich environment similar to that found in the cores of elliptical galaxies. I will discuss the following important relations regarding the variability of AGNs in the Starburst model: 1) The supernova rate (or number of peaks of the light curve) and its relation to the blue-band stellar luminosity. 2) The equivalent width of H/? and its relation to the average energy per SN event. 3) The variability and its relation to the core luminosity of low-luminosity AGNs. I briefly review the time dependent process that occur prior to thin shell formation in a rapidly radiating supernova remnant, i.e. as it achieves maximum luminosity. This process, which has a typical time scale of a few weeks and involves energies of about 5 % of the total explosion energy, produces time delays between the continuum and line emission and emission line luminosities with values similar to those observed in low luminosity AGN. In section 6, I present a simple estimate of the core radius of the light distribution of Starburst-cores as inferred from the observed light distribution of nearby low-luminosity starbursts like 30 Dor and NGC 3603.

2. Estimation of the supernova rate in young clusters The variability of AGNs is a key aspect that must be explained by any theory of AGNs. The best information regarding variability in AGNs is provided by a small group of nearby low-luminosity type 1 Seyferts. The optical light curves of these AGNs show two distinct components: an occasional sharp peak variation superimposed onto a long-term recurrent modulation (Lyutyi 1977, 1979; DibaT k Lyutyi 1984; Lyutyi k Oknyanskii 1987; Smith et ai. 1991), called rapid and slow components, respectively (see Smith et ai. 1991 for a schematic picture). The peaks of the rapid component last some weeks, while the cycles of the slow one are about two years long. In the Starburst model, the slow component is related to the long-term behaviour of cSNR, while the rapid component is mostly due to cooling instabilities in the shells of cSNRs and to SN flashes. A double peak variation is expected in the light curve of a SN evolving in a high-density circumstellar medium, the first peak corresponding to the SN flash, and the second one to the time when the cSNR reaches the maximum of its radiative phase (Wheeler et ai. 1980; Shull 1980; Terlevich et ai. 1992). The typical time scale of the radiative phase is,

{V 3 / 4 . 51

(2.1)

where e5i is the energy of the cSNR in units of 10 erg and n-j the circumstellar density in which this cSNR evolves, in units of 107 cm" 3 . During the radiative phase, the luminosity of the remnant can be approximated as,

Terlevich: From SO Doradus to QSOs / , x-11/7

L bol (<)~lO lo L 0 e 7 { 8 nf - M

331

,

(2.2)

X'sg/

and is emitted mainly in the far UV and X-rays. Aretxaga k Terlevich (1994; hereafter AT) have shown that cSNR-type events can account for the slow-component variations in the well sampled light curve of NGC 4151 both in energy and time scale of the process and estimated the number of cSNRs required to explain the luminosity of the nucleus of an AGNs. AT calculated the specific supernova rate L^/Mt and specific blue luminosity ^SN/Mt, for a variety of IMFs using stellar evolutionary models from Maeder (1990). It was found that both parameters decrease with increasing age of the cluster and are similarly dependent on the assumed IMF. A very important consequence of this finding is that the ratio of the supernova rate per unit mass to the blue luminosity per unit mass, i/SN/Mt/L^/Mt = " S N / ^ B ' *-e- ^ne supernova rate per unit stellar blue luminosity, is almost independent of the assumed IMF and basically constant over the whole SN II phase. This is because both quantities, the SN rate and the blue luminosity, depend mainly on the number of stars in a narrow mass range close to the turn-off point, and do not depend on the number of low-mass stars that are the bulk of the mass of the cluster. Figure 1 shows the SN rate as a function of the absolute blue magnitude of the stars of the cluster, and Table 1 gives some of the relevant ratios for a wide range of ages and IMFs. The main result is that a young cluster with age between 10 and 60 Myr and with ~ —21.5 produces 1 SN per year, almost independently of the assumed IMF.

tr

1Wfl("SN = 1 )

"SN/£B

1.0 2.0 3.0 4.0 5.0

-21.7 -21.5 -21.3 -21.3 -21.3

6.0

-21.4

0.15X10" 10 0.17X10- 10 0.21 xlO" 10 0.21xl0~ 1 0 0.21xl0~ 1 0 0.18X10- 10

m\ = 0.3 Mo

JMt for a = 2.35 m\ = 1 M© m 1 = 3 M©

28.70 17.10 12.80 11.00 9.44 8.78

48.80 29.10 21.70 18.60 16.00 14.80

78.00 43.40 30.60 25.10 20.50 18.60

A/SN

19.8 12.1 9.8 8.4 7.6

6.9

1. Description of the cluster evolution: (1) time, in 107 yr; (2) absolute blue luminosity of a cluster with a SN rate of 1 per year; (3) ratio SN rate-stellar blue luminosity, in yr" 1 LB© - 1 ; (4-6) ratio stellar blue luminosity-total initial mass of the cluster for several power-law IMFs of slope a=2.35 and lower mass limits (mi), in solar units; (7) initial mass of the stars undergoing explosion, in solar units. TABLE

3. Variability in very low luminosity AGNs Due to the stochastic nature of the SN activity and the low SN rate in low-luminosity AGNs the nucleus can go through transient quiescent phases when any light from the cSNR is faint and, consequently, no broad lines are detected, while the continuum is still blue due to the presence of hot stars. This is likely to be the case for the deep photometric minima of 1981 and 1984 in NGC 4151 during which the nucleus resembled

332

Terlevich: From 30 Doradus to QSOs

-20

M\ FIGURE 1. Supernova rate vs. stellar blue luminosity for three different IMFs. Dotted line, a = 3 and mi > 0.5 MQ; continuous line,a = 2.35 and mi > 0.3 M©; dashed line,a = 2.35 and mi > 3 M 0 .

a type 2 Seyfert (Antonucci k Cohen 1983; Penston k Perez 1984; Lyutyi et al. 1984; AT). This characteristic of the SN activity gives a simple explanation for the transient stages of type 2 Seyfert observed as LINERS and type 1 Seyfert nuclei. One important and strong prediction of the SB model for AGNs is that the variability amplitude observed in AGNs should reach a maximum for those nuclei with luminosity similar to the maximum luminosity of cSNRs, i.e. MB 20. AGNs with higher luminosity should show less amplitude of variability due to the superposition of cSNR events, while lower luminosity AGNs, i.e. those with MB > -20 (LLAGNs) should also show a smaller variability amplitude produced by old and less luminous cSNRs which evolve slowly in their spectral properties. The SN rate is determined by the nuclear continuum blue luminosity. A Starburst-core with a continuum magnitude of MB = -16.3 should have 1 SN event every 100 years. As the cSNR remains detectable against such a background for about 10 years, one in every 10 LLAGNs with MB 16.3 should show broad Ha emission and 1 in every 100 should be found to have a very bright BLR corresponding to a cSNR less than a year old. Two additional predictions for LLAGNs that stem from the properties of old cSNRs are that, the FWHM of the broad lines in LLAGNs should be smaller than in more luminous AGNs and, furthermore, the FWHM should decrease slowly with time as FWHM a V,fcOcta <~T while its luminosity should decrease as LB O ;« t~^ and the X-ray spectrum of the BLR of LLAGNs should be much softer than that of typical Seyfert type 1. A relation between the FWHM of the BLR and the X-ray spectrum is therefore expected in the sense that those LLAGNs with lower luminosity should have also smaller FWHM and softer X-ray spectra.

4. The energy per supernova In the Starburst model for AGNs, the value of the average energy per SN is not a free parameter but is constrained by the observed equivalent width of Balmer recombination lines, i.e. the ratio of the total flux in the line to the monochromatic flux in the underlying continuum, like H/? ( Wj^). In this model the flux in the broad lines is proportional to the total ionizing flux emitted by the cSNR and therefore proportional to I/SN, while the

Terlevich: From 30 Doradus to QSOs

333

optical continuum is mainly due to the stars. The constancy of the ^ S N / ^ B described

in the previous section implies that the Wj^g depends only on the energy radiated per SN. The emitted H/? luminosity of a cSNR corresponds to ~ 3% of the total luminosity (Terlevich et ai. 1992). On the other hand, the underlying continuum comes mainly from the stars of the cluster with a contribution of ~ 17 c^\ % from the cSNR. Combining Maeder (1990) stellar evolutionary models for the stellar component with Terlevich et ai. (1992) models for cSNRs, AT determined that the W^g in a Starburst-core is,

independent of the SN rate and, therefore, of the total luminosity of the AGN. Furthermore, Wjjfl, like I/SN / ^B > *s a l m o s t independent of the choice of IMF or cluster age; it is, nevertheless, weakly dependent on the adopted bolometric correction for the cSNR. The use of equation (4.3) to determine the value of €51 gives a result that is independent of the adopted cosmology.

Equation (4.4) is of fundamental importance. It indicates that the observed narrow range of values of the W^g in AGNs (Yee 1980; Shuder 1981; Goodrich 1989; Netzer 1990; Binette, Fosbury k. Parker 1993) may be related to a universal value of the total energy per SN. This is a strong result, almost independent of the choice of either cosmology, IMF, or SN bolometric correction. Equation (4.4) therefore provides a strong test of the model. The published values for the universal W^g in AGN ranges from 50 A (Shuder 1981), 63 A for X-ray selected AGNs (Goodrich 1989) to 100 A in type 1 Seyferts (Goodrich 1989; Binette et ai. 1993). The higher values correspond to the Wy^g measured with respect to the so-called non-thermal continuum and represent an over-estimate of the observed Wjja. For an average Wjj/j ~ 75A Eq. (4.4) predicts the average energy per SN to be 2.8 x 1051 ergs" 1 . The best modern determinations of the kinetic energy of type II SNe are: 3 x 1051 erg for SN 1979C (Branch et ai. 1981), and 1.5 x 1051 erg for SN 1987A in the LMC (Woosley 1988). 5. The rapid UV variability and the nature of the lag Terlevich et ai. (1994) have analysed the time-dependent process that occurs prior to thin-shell formation in rapidly radiating cSNR, as maximum luminosity is achieved. They have found that an inherent delay between the photon emission and the time required to increase the density of the cooling gas leads to a lag between the observed continuum burst and the emission line response. This delay is intrinsic to the physical processes investigated and not related to the geometry of the system; there are in fact no light-crossing time arguments involved. The total width of the cold or photoionized region is only about 1013 cm (about 300 light seconds), yet, delays of up to several weeks between continuum and lines emission are generated. Early in the evolution of the cooling wave, large photon fluxes are available but the UV/optical emission lines are very faint due to the high ionization parameter associated with the still relatively low electron density and the low column density of the cold material. As the cooling wave evolves, however, the column density of cold material soon exceeds 1021 cm"2 and its density starts to increase over the post-shock value (4 x 107

334

Terlevich: From 30 Doradus to QSOs

cm 3 ). As a consequence the ionization parameter, initially > 10, begins to decrease and lines like NV 1240 and Hell 1640 start to be emitted copiously. By the time the density of the compressed shell reaches ~ 2 x 109 cm" 3 the ionization parameter has dropped to ~ 0.1 and the electron temperature of the ionised gas is ~ 104 K. At this time lines like Lya and CIV reach their maximum intensity with a delay of several days with respect to the start of the UV burst. CHI shows longer delay in all models. Other lines, like the Balmer series, have even longer delays and the longest predicted delays are for Mgll and Fell that show lags of more than a month. Only one lag-producing flare per cSNR is predicted. This lag-producing flare is the longest and most energetic one. As these luminous flares are produced at the maximum of the light curve, one important prediction is that flares showing lag between continuum and emission lines are expected to occur close to the maximum of the light curve. Conversely, no further luminous and lag-producing flare is expected close to the absolute minima of the light curve. Terlevich et ai. compared the predicted shape and luminosity of the light curve of the emission lines from the cooling wave models with the observations of NGC 5548. For this they compared the models with the best sampled and strongest of the pulses, the second one. Table 2 shows a compilation of the cross-correlation results of the UV and optical campaign of NGC 5548 (Clavel et ai. 1991; Peterson et ai. 1991) together with the results from the models. A quick inspection of Table 2 indicates that the observed lags for all the observed emission lines are similar to the model predictions. Furthermore, the observed and predicted peak luminosities of the emission lines are remarkably similar.

LAG (days) Observed Predicted

Lyal216 NV 1240 SUV 1440 CIV 1550 Hell 1640 CIII 1909 Mgll 2800 H/74861

12 4 12 - 34 8-16 4-10 26 - 32 34 - 72 20

10-15 2-10 12 - 17 10-15 12-20 20 - 30 40 - 50 30 - 40

Peak luminosity Observed

•8° ; Predicted

100. 30. 13. 100. 15. 18. 10. 8.

90. 14. 23. 280. 26. 13. 26. 7.5

2.Observed and predicted LAG in NGC 5548: (1) Emission line identification; (23)Observed and predicted lag in days; (4-5) Observed and predicted peak line luminosity in 1040 ergs" 1 units. . TABLE

A finer degree of comparison between theory and observations can be achieved by studying in detail the light curve of the different emission lines. The observed light curves corresponding to the UV continuum and line emission for the strongest event in NGC 5548 , are shown with open circles in Figure 2. The value of the luminosity corresponds to the excess over the value of the flux at the start of the flare. The origin in the time axis corresponds to JD2447580. Strong emission lines like CIV and Lya have a well defined light curve with small errors, weaker and/or blended lines like Hell, NV or Mgll have errors of comparable amplitude to the observed amplitude of variability.

Terlevich: From 30 Doradus to QSOs

335

UV Continuum

e

oa

o

0 a o

9 Time / 10 days

10

FIGURE 2. Predicted and observed flare light curves in NGC 5548. The open circles indicate the light curve of the continuum and line emission of the strongest flare observed during the 1988 monitoring campaign. The zero in intensity corresponds to the value of the flux at the start of the flare. The broken line shows the results of the spherically symmetric model.

Superposed on the observational results are the results of our spherically symmetric nat2, ne = 108 cm" 3 model. The vertical line indicates the centre of the continuum pulse. The amplitudes of four emission lines, CIV, NV, SilV+OIV and Mgll, have been scaled in amplitude by factors around 2, to help the comparison of the shape of the light curve. The other four emission line light curves have no scaling and the result from the model is directly plotted on top of the observations. The similarity of the predicted and observed light curves is striking. The predicted shapes and delays are in very good agreement with the observations, and the observed ionization dependence of the delay is nicely reproduced by the models. The match of Lya, Hell, CIII and H/? predicted and observed light curves is excellent. For the other lines the predicted peak luminosity is either too strong (CIV,SiIV+OIV,MgII) or too weak (NV) but the ratio of observed to predicted luminosity is at most a factor of three (see CIV in Table 2). This is important, given that the model predicts not just relative fluxes but absolute luminosities in each line. After the first cooling event, dense cold gas will always be present in the outer shell immediately reprocessing the inward-moving ionizing radiation produced by further cooling instabilities; little or no lag occurs between continuum and lines as was observed in the 3rd pulse of the NGC 5548 monitoring campaign. Given the simplicity of the model and the small number of free parameters (basically only one, the circumstellar density), it is remarkable that the simplest cSNR model for the BLR ( spherically symmetric with constant density CSM) is capable of giving a detailed and accurate description of the physical conditions and line ratios (Terlevich et ai. 1992), and also gives an accurate prediction of the temporal behaviour of the broad

336

Terlevich: From 30 Doradus to QSOs

emission lines in NGC 5548. It is even more remarkable when considering that the value of the circumstellar density is not a free parameter but measured in true cSNRs like SN 1988Z, SN 1987F and bright radio supernovae (Filippenko 1989; Stathakis and Sadler 1991; Terlevich et ai. 1994).

6. The size of a Starburst To estimate the optical size of Starburst-cores, I will use recent results on the structure of the two nearest giant bursts of star formation, 30 Dor and NGC 3603. According to Moffat, Seggewiss & Shara (1985) and Campbell et ai. (1992) their core radii are: 30 Dor Re < 0.07 pc NGC 3603 Rc - 0.024 pc

,

{

. >

In a recent work Malamuth & Heap (1994; see also this proceedings) have questioned Campbell et ai. 's determination of the Rc of 30 Dor and suggest that if their innermost data point is disregarded, the luminosity profile can be also fitted by a King-model having a core radius about twice the above value. A very important finding of Malamuth and Heap is strong mass segregation in 30 Dor in the sense that the central region has a higher fraction of massive stars than the outer one. This is reflected in the value of the core mass radius of 0.21 pc, larger than the optical core radius. I therefore adopt a value of O.lpc for the core radius of the luminosity profile of 30 Dor. The core blue luminosities of 30 Dor and NGC 3603 for (B - V) = -0.3 are: Ma(core3o Dor) = —11.4 and MB(coreNGC 3603) == —10.2. Assuming that both clusters have the same M/L ratio, the fact that NGC 3603 is about 3 times less luminous than 30 Dor implies that its mass should be about 3 times smaller. I will assume that these clusters are close to virial equilibrium, i.e. L% a M< a Rc o~2. From Melnick, Tapia and Terlevich (1989) the ionized gas velocity dispersion in NGC 3603 is 16 kms" 1 , and from Melnick et ai. (1987), the gas velocity dispersion for 30 Dor is 22 kms" 1 . If the ionized gas velocity dispersion is equal to the stellar velocity dispersion (Tenorio-Tagle, Murioz Tuiion 1993). The larger luminosity and velocity dispersion in the core of 30 Dor with respect to the core of NGC 3603 implies that the core of 30 Dor should be ~ 6 times bigger than that of NGC 3603, consistent with the observations. The expectation therefore is that a normal young cluster with core luminosity around A/j = —10 and cr = 20 kms" 1 will have a core radius of Rc ~ 0.02 pc. To scale these estimates to the luminosities of Starburst-cores of AGNs, presumably associated with young cores in massive spheroids, I have used the virial theorem and assumed the same M/L ratio for aii young clusters, Rc ~ 0.02pclO-°-4(M>+10> {^\

(6.6)

This equation predicts for a nearby AGNs like NGC 4151 (Mb = -19.5, a = ) &nRc < 1.5 pc, or less than 12m.a.s.;an L* QSO (M4 = -25,
Terlevich: From 30 Doradus to QSOs

337

predicted size is given by: (6.7)

The predicted sizes for the light distribution using (6.7) are: Rc < 0.4 pc or < 3 m.a.s. for NGC4151, Re < 9 pc or < 1 m.a.s. for the L* QSO at redshift 2, and Rc < 150 pc or < 20 m.a.s. for a Mb = -28 QSO at redshift 2. It is interesting to compare these predictions with the core mass distribution of present day ellipticals of Terlevich & Boyle (1993; TB). For the three AGNs listed above, the TB scaling gives sizes about a factor of 6 larger than those given by Eq. (6.7). There are two important points that should be taken into account when comparing the mass distribution of an old stellar cluster with the light distribution of a very young one. First, mass segregation as detected in 30 Dor and most galactic open clusters implies that massive stars preferentially populate the core of the young cluster; therefore, the light distribution will be smaller than the mass distribution. A cluster with a Salpeter IMF and complete segregation of mass will have all the stars more massive than 10 MQ confined to a region of 15% of the mass effective radius. In the case of equipartition of kinetic energy, the massive stars will have smaller velocity dispersion and therefore will be even closer to the centre (Binney & Tremaine 1987). The second aspect is dynamical evolution. Due to the large mass loss in the form of stellar winds and SN ejecta from the central region of the cluster and its short dynamical time scale, the structure will be affected by dynamical evolution in less than a Hubble time. A cluster with a Salpeter IMF will lose about half of its initial mass after one Hubble time, ~ 25% will be lost in the first 100 Myr. According to Richstone & Potter (1982), slow mass loss will produce an adiabatic expansion proportional to the mass loss,

Abrupt mass loss of a small fraction of the cluster mass is twice as effective as the same amount of slow mass loss in puffing up the cluster (Biermann & Shapiro 1979). The combined effect of mass segregation and dynamical evolution can produce a mass distribution that has a core radius after a Hubble time up to 30 times larger than the initial light distribution even in the absence of equipartition. In the more conservative assumption that the mass segregation in the young cluster was only partial, a ratio of about 0.1 between the present day mass distribution Rc and the initial light distribution Rc is expected. This is consistent with the observations of present-day cores of elliptical galaxies and of nearby bursts like 30 Dor and NGC 3603. For these clusters to remain bound through their evolution, the increase in size with time should be accompanied by a corresponding reduction in the velocity dispersion. The possibility of mass segregation raises important questions that are worth investigating. 1) Are massive young clusters in or close to equipartition of energy? 2) Given that it is not possible to determine the size of the spatial distribution nor to measure the velocity dispersion of the low mass stars that make the bulk of the mass of the cluster, is it possible to determine the total mass or the M/L ratio of a cluster with mass segregation? 3) Is the mass segregation responsible for the claims (Rieke k. Lebofsky 1981) of topheavy IMFs in Starbursts?

338

Terlevich: From 30 Doradus to QSOs

7. Discussion of Heckman's criticism In a recent review Heckman (1994) has criticised the Starburst model for AGNs on three grounds: (A) That the angular sizes predicted by the TB model already exceed the upper limits to the angular sizes of high-redshift quasars observed with the Hubble Space Telescope. The "snapshot survey" of high-z QSOs suggests an upper limit to the angular size of 0.1 arcsec. In contrast, TB core sizes for the corresponding luminosities are ~ 0.3 arcsec. As discussed in section 6, if all Starbursts exhibit mass segregation as nearby ones do, the mass distribution should have a larger size than the light distribution. The observed mass segregation also implies that there is kinematical segregation. The massive stars are close to the core because their velocity dispersion is smaller than that of the low mass stars. Furthermore, dynamical evolution induced by stellar mass loss will make the cluster expand up to a factor of 3 during a Hubble time and will reduce the velocity dispersion. (B) The lack of an intrinsic Lyman edge in the spectra of QSOs indicates that the UV continuum is not due to stars. Antonucci, Kinney & Ford (1989) and Koratkar, Kinney & Bohlin (1992) showed that any intrinsic Lyman edge must be less than 15% in amplitude in QSOs. The behaviour of the Lyman edge in the Starburst model has to be properly investigated. There are two aspects to consider, first the observational limits given by work like that of Koratkar et al. (1992) refer only to Lyman edges generated in accretion disks, i.e. edges broadened by the velocity field of the disk. In fact 32% of their sample has sharp Lyman edges at the same redshift as the BLR emission lines, most of these Lyman edges have relatively narrow Lya and metal lines that like CIV 1550 A are expected to be present also in the UV spectra of massive stars. Second, and more important, is that it is difficult to have a good theoretical prediction of the UV continuum of a Starburst with SN remnants. The best work to date is that of Bithell (1991) indicating that the radiation from the remnants may lead to a very small observed rest-frame Lyman break. (C) That the observed variability in AGNs implies that the optical-UV continuum cannot be starlight. Heckman argues that the firm conclusion is that the optical-UV continuum in quasars can not be starlight, as TB implicitly assumed. Rather the model requires the continuum to be powered by the thermalization of supernova kinetic energy rather than by nucleosynthesis. Since the X-rays are also produced in this way in the TB model, and since the IR is simply reprocessed optical and UV light, this is tantamount to saying that most of the bolometric luminosity of quasars must be powered by supernova remnants. This then leads to several severe energetic problems. Heckman estimated the average amount of quasar radiant energy and the implied number of supernovae per elliptical galaxy in the TB model. There have been several recent estimates of the total amount of radiant energy per unit present-day volume element produced by quasars over the history of the universe (Padovani, Burgh, & Edelson 1990; Chokshi & Turner 1992). These estimates are especially useful because they are independent of the cosmology. Heckman has revised those estimates and concluded that their value should be increased by about a factor of two to extend down to B = 25 and out to z = 5. The results for the total amount of quasar radiant energy per unit present-day volume element are 8.3 x 10 53 ergMpc~3 for the revised Padovani et al. calculation and about half this value for Chokshi & Turner. Similar values are found by Heckman for the X-ray background, finally adopting 7 x 10 53 ergMpc~3. The present-day luminosity density in optical light due to elliptical galaxies is about 1.4 x 107 LQ Mpc~3(for AQ = 75 ifcms"1 Mpc~l). Dividing the quasar energy density by

Terlevich: From 30 Doradus to QSOs

339

the elliptical galaxy luminosity density then implies that the amount of quasar radiant energy per elliptical galaxy with luminosity LE is 5.Ox 1O61 [LE / 1O10 LQ] erg. To gel this much energy from thermalizing the kinetic energy of supernovae requires the formation of a total stellar mass of 6.2 x 10 12 MQ[LE / 1O10 Le] for a normal Salpeter IMF extending from O.I to 120 MQand "only" 13% of this value for an IMF with only massive stars. These huge masses can be compared to the present-day masses of the cores of ellipticals. Adopting the parameters in TB the ratio of the mass required by quasar energetic to the present-day mass of a core is then 900 [LE / 1O10 LQ]~03S for a normal Salpeter IMF and l\O[LE / 1O10 I 0 7 " 0 3 5 for a "massive stars only" IMF. From this Heckman concludes that there is a severe energetic problem with the TB model, especially since a "massive stars only" IMF is logically inconsistent with the model (since it is supposed to explain the formation of the cores of ellipticals, explicitly including the low-mass stars that dominate these cores today).

Heckman's estimates have several weak points. Firstly, his "firm conclusion" that the main source of energy cannot be starlight and should therefore be SNe is wrong. As was mentioned in several contributions, (Terlevich 1990a,b; Terlevich et a/. 1992) the main source of energy in a starburst are the stars. SNe and cSNRs provide only a fraction of the bolometric output of the burst, otherwise variability would be too large. The radioquiet AGNs with the largest variability are the Seyfert 1 systems at about Mj 20. In these systems the maximum peak-to-peak variability is ~ 1 mag in the blue band indicating that the maximum contribution of a cSNR at their peak flux is similar to that of the stars (AT). The SN contribution averaged over time is therefore smaller than that of the stars. Furthermore, we have seen in section 4 that to explain the observed WJJ o and the observed variability in nearby AGNs like NGC4151, the energy per SN should be in fact ~ 3 x 1051 erg. Even for this value of the energy per SN the time averaged cSNR contribution in bolometric units is only 5% to 30% of the stellar bolometric luminosity for clusters of 10 to 40 Myr of age. Only for cluster older than 60 Myr and for €$\ = 3 the cSNR contribution to the bolometric luminosity is comparable to that of the stars (Terlevich 1990a,b; AT). Secondly, the estimates of the background light per unit volume produced by QSOs have a large range of values and large uncertainty. Taking Choksi & Turner's (1992) value and applying a bolometric correction Bc = 6 we obtain EBOL/V = 1.2 X 1058 erg Mpc~3. I believe that the bolometric correction of Sanders et ai. (1989) used by Choksi & Turner (1992) is too large (Bc = 16.5). This is because the predicted energy output of the central engine in the unobserved spectral range between 1 Rydberg and 40 Rydberg (~ 0.5 keV) has been included in the estimate of the bolometric correction. This is probably wrong because if at least part of the energy emitted by AGNs at frequencies > 1 Rydberg comes from the reprocessing of the radiation emitted between 1 and 40 Rydbergs, as photoionization models suggest, the inclusion in the estimate of the bolometric correction of both the low-energy reprocessed output and the UV/soft X-ray ionizing input, implies a large overestimate of the bolometric correction. This overestimate is made even larger due to the fact that the assumed QSO emitted spectrum peaks between 1 and 40 Rydberg. Using only the observed multi-frequency spectrum, the bolometric correction should be around 5. A similar argument applies to Heckman's estimate of the cosmic X-ray background. Using a bolometric correction of 50 and taking into account that the fraction of the background due to QSOs is 40% and not 80% as assumed by Heckman (Boyle et al. 1993; Georgantopolous et ai. 1993; Maccacaro et al. 1991), I obtain EBOL/V = 1.3x 1058 erg Mpc~3. The average of both estimates is therefore about a factor of at least 7 below that of Heckman.

340

Terlevich: From 30 Doradus to QSOs

Thirdly, the lower limit of the IMF adopted by Heckman is lower than that adopted by TB producing a factor of ~ 2 discrepancy in the total mass of stars. Finally the fact that there is more than one burst per elliptical in the TB model capable of producing UV radiation introduces a further 1.5 factor correction. All these factors go in the same direction and add up to produce a total systematic shift of > 10 x 7 x 2 x 1.5 = 210. The large factors of 900 and 110 computed by Heckman are reduced to < 4.5 for the complete IMF case and < 1 for the massive stars only IMF. All this illustrates the inherent uncertainties in using this type of indirect approach. Precisely because of these uncertainties TB decided to use instead the observed blue luminosity of present-day ellipticals and QSOs. Simply, we argue that the mass of the metal-rich core of an elliptical galaxy is about 5% of the total mass and has a present day M/L ratio of about 30 in the brightest {MB — —24) nearby ellipticals. The present-day luminosity of the core is therefore 5% of the galaxy luminosity or MB = —20.8. If the core was formed in a instantaneous burst its luminosity when it was 10 Myr old can be estimated from the ratio of the present-day M/L to the M/L ratio of a young cluster (see Table 1). The ratio of M/L is ~ 1100 for mj = 0.5 M©. The young cluster was therefore AMB = 7.6 magnitudes brighter than the present-day core corresponding to MB = —28.4, a value close to the luminosity of the brightest QSOs. The more complex estimate made by Heckman contrast with the simple method used by TB. They showed that the predicted young core luminosity function is an excellent match to the observed luminosity function for QSOs in the redshift range from 2.0 to 2.9. And that the models for the subsequent evolution of the star formation rate in the cores of elliptical galaxies is qualitatively similar to that required to explain the observed evolution of the QSO luminosity function. For a go = 0.5 universe, the rate of evolution predicted by the starburst model comes closer to predicting the observed evolution (L oc t~2) than previous attempts involving supermassive black holes. The TB approach has far smaller uncertainties and strongly supports the idea that the luminosity radiated during the formation of the metal-rich core of the elliptical galaxies is enough to explain the observed luminosity of even the most luminous QSOs, the cSNR activity being responsible for the variability and X-ray and radio emission. 8. Conclusions Several important predictions in the Starburst model are relatively simple to check. As the supernova rate (or number of peaks of the light curve) is related to the blue band stellar light (or minimum of the light curve recorded), independently of the initial mass function, age and/or total mass of the cluster, there should be a tight relationship in low-luminosity AGNs between the number of slow component peaks in the light curve (i.e. the SN rate) and the luminosity at minimum (i.e. the unperturbed cluster luminosity). The variability amplitude observed in AGNs should reach a maximum for those nuclei with luminosities similar to the maximum luminosities of cSNRs, i.e. MB ~ —20. AGNs with higher luminosity should show less amplitude of variability due to the superposition of cSNR events, the same as LLAGNs, i.e. those with MB ^ —20 because they will be representing old and less luminous cSNRs which evolve slowly in their spectral properties. The FWHM of the BLR lines in LLAGNs should be smaller than in more luminous AGNs and should decrease slowly with time; also the X-ray spectra of the BLRs of LLAGNs should be softer than those of typical type 1 Seyferts. The least luminous LLAGNs should therefore also have the narrowest broad lines and softer X-ray spectra. Observations of a well defined sample of LLAGNs should provide enough information to test the predictions of the model. An extended study of the variability of AGNs in the

Terlevich: From SO Doradus to QSOs

341

Starburst model is clearly needed, which includes AGNs with well-sampled light curves and a wide range of luminosities to test the consistency of the predictions. Due to the stochastic nature of the SN activity and the low SN rate in low-luminosity AGNs, transient stages of type 2 Seyferts, as observed in some broad-line LINERS and type 1 Seyfert nuclei, can naturally occur. The time averaged Wg/j is related to the mean energy per SN, independently of the initial mass function, age and/or total mass of the cluster, and of the assumed cosmology. The constancy of W^g in AGNs may be related to a near universal value of the energy in a type II SN explosion of about 3 x 10 51 erg . Time dependent processes during thin-shell formation in a rapidly radiating supernova remnant produce time delays between the continuum and line emission with values similar to those observed in nearby type 1 Seyferts. The predicted delays are shorter for the highionization lines than for the low-ionization ones. The theory also predicts the occurrence, after shell formation, of shorter and less energetic flares with little or no lag between continuum and lines. Based on the latest HST results about the light distribution in 30 Dor I estimated that the light distribution for a Starburst-core can have a FWHM < 1 pc for a nearby AGN like NGC 4151, or less than 6 m.a.s in diameter, while a bright QSO (Mb = - 2 8 ) at redshift 2 would have a FWHM < 150 pc, or less than 40 m.a.s.. Obviously, much more work is needed both in further developing the theory and in verifying the predictions. Important questions that remain to be answered include: • Can the Starburst model explain the observed rapid X-ray variability on time scales down to few hundred seconds? • Can the Starburst model explain the observed properties of radio loud AGNs? I would like to thank my collaborators: Itziar Aretxaga, Brian Boyle, Jose Franco, Jorge Melnick, Michal Rozyczka and Guillermo Tenorio-Tagle with whom most of the work presented here was done. Thanks are due to Elena Terlevich and Brian Boyle, whose comments made this contribution readable.

REFERENCES R., KINNEY, A. & FORD, H. 1989 Astrophys. J. 342, 64. (AT) ANTONUCCI, R. R. J. fc COHEN, R. D. 1983 Astrophys. J. 271, 564. ARETXAGA, I. & TERLEVICH, R. 1994 Mon. Not. R. Astron. Soc. In press. BIERMANN, P. & SHAPIRO, S. L. 1979 Astrophys. J. Lett. 230, L33. BlNETTE, L., FOSBURY, R. A. & PARKER, D. 1993 Pub. Astron. Soc. Pac. 105, 1150. BlNNEY, J. & TREMAINE, S. 1987 Galactic Dynamics. Princeton University Press. BlTHELL, M. 1991 Mon. Not. R. Astron. Soc. 253, 320. ANTONUCCI,

BOYLE, B. J., GRIFFITHS, R. E., SHANKS, T.,

STEWART, G. C. & GEORGANTOPOULOS, I.

1993 Mon. Not. R. Astron. Soc. 260, 49. BRANCH, D., FALK, S. W., MCCALL, M. L., RYBSKI, P., UOMONTO, A. K. & WILLS, B. J.

1981 Astrophys. J. 244, 780. et aJ. 1992 Astron. J. 104, 1721. CHOKSHI, A. & TURNER, E. 1984 Mon. Not. R. Astron. Soc. 259, 421. CHU, Y.-H. & KENNICUTT, R. 1994 Astrophys. J. In press. CLAVEL, J. et aJ. 1991 Astrophys. J. 366, 64. CAMPBELL

DIBAK, E. A. & LYUTYI, V. M. 1984 28, 7. FILIPPENKO,

A. V. 1989 Astron. J. 97, 726.

342

Terlevich: From 30 Doradus to QSOs

GEORGANTOPOULOS, I., STEWART, G. C ,

SHANKS, T., GRIFFITHS, R. E. & BOYLE, B.

J.

1993 Mon. Not. R. Astton. Soc. 262, 619. GOODRICH, R. W. 1989 Astrophys. J. 342, 224. HECKMAN, T. M. 1994 Mass-Transfer-Induced Activity in Galaxies. University of Kentucky. In press. HAMMAN, F. k. FERLAND, G. 1992 Astrophys. J. Lett. 391, L53. KORATKAR, A., KINNEY, A. & BOHLIN, R. 1992 Astrophys. J. 400, 435. LYUTYI, V. M. 1977 Sov. Astron. 21, 655. LYUTYI, V. M. 1979. Sov. Astron. 23, 518. LYUTYI, V. M., OKNYANSKII, V. L. & CHUVAEV, K. K. 1984 Sov. Astron. Lett. 10, 335. MACCACARO, T., DELLA CECA, R., GIOIA, I. M., MORRIS, S. L., STOCHE, J. T. fc WOLTER,

A. 1991 Astrophys. J. 374, 117. A. 1990 Astron. Astrophys. Suppl. 84, 139. MALUMUTH, E. M. & HEAP, S. R. 1994 Astron. J. 107, 1054. MELNICK, J., TAPIA, J. & TERLEVICH R. 1989 Astron. Astrophys. 213, 89. MAEDER,

MELNICK, J.,

MOLES, M.,

TERLEVICH, R. & GARCIA-PELAYO, J. M. 1987

Mon.

Not.

R.

Astron. Soc. 226, 849. MOFFAT, A. F. J., SEGGEWISS, W. & SHARA, M. M. 1985 Astrophys. J. 295, 109. NETZER, H. 1990 Active Galactic Nuclei(ed. T. J.-L. Courvoisier & M. Mayor). Springer-Verlag. PADOVANI, P., BURGH, R. & EDELSON, R. 1990 Astrophys. J. 353, 438. PENSTON, M. V. & PEREZ, E. 1984 Mon. Not. R. Astron. Soc. 211, 33P. PETERSON, B. et aJ. 1991 Astrophys. J. 368 119. RlEKE, G. H. & LEBOFSKY, M. J. 1981 Astrophys. J. 250, 87. RICHSTONE, D. O. & POTTER, M. D. 1982 Astrophys. J. 254, 451. SHUDER, J. M. 1981 Astrophys. J. 244, 12. SHULL, J. M. 1980 Astrophys. J. 237, 769. SMITH, A. G., NAIR, A. D. & CLEMENS, S. D. 1991 In Variability of active galactic nuclei (ed. H. R. Miller & P. J. Wiita). Cambridge University Press. STATHAKIS, R. A. & SADLER E. M. 1991 Mon. Not. R. Astron. Soc. 250, 786. TENORIO-TAGLE, G., TERLEVICH, R., ROZYCZKA, M., FRANCO, J. & MELNICK, J. 1992

In

The Nearest Active Galaxies (ed. J. Beckman, L. Colina &: H. Netzer). CSIC, Madrid. TENORIO-TAGLE, G., MUNOZ TUNON, C. & Cox, D. P. 1993 Astrophys. J. 418, 767. TERLEVICH, E., DIAZ, A. I. & TERLEVICH, R. 1990 Mon. Not. R. Astron. Soc. 242,

271.

R. 1990a In New windows to the Universe (ed. F. Sanchez & M. Vazquez). Cambridge University Press, Cambridge. TERLEVICH, R. 1990b In Windows on Galaxies (ed. G. Fabbiano et a]. ). Kluwer Academic Publishers. TERLEVICH, R., MELNICK, J. & MOLES, M. 1987 In IAU Symp. # I2t: Observational Evidence of Activity in Galaxies (ed. E. Ye. Khachikian, K. J. Fricke & J. Melnick). Reidel, Dordrecht. TERLEVICH, R. & MELNICK, J. 1988 Nature 333, 239. TERLEVICH, R., TENORIO-TAGLE, G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. R. Astron. Soc. 255, 713. TERLEVICH, R. & BOYLE, B. J. 1993 Mon. Not. R. Astron. Soc. 262, 491. TERLEVICH,

TERLEVICH, R., TENORIO-TAGLE, G., ROZYCZKA, M., FRANCO, J. & MELNICK, J. 1994

Mon.

Not. R. Astron. Soc. In press. WHEELER, J. C , MAZUREK, T. J. & SIVARAMAKRISHNAN, A. 1980 Astrophys. J. 237, 781. WOOSLEY, S. E. 1988 In Supernova 1987A in the Large Magellanic Cloud (ed. N. M. Kafatos & A. Michalitsianos), p. 289. Cambridge University Press. YEE, H. K. C. 1980 Astrophys. J. 241, 894.

Distance Indicators to Low-Luminosity AGN By ITZIAR ARETXAGA AND ROBERTO J. TERLEVICH Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, U.K. Variability is one of the most conspicuous properties of AGN. The starburst model postulates that the variability observed in radio-quiet sources is produced by the supernova (SN) and compact supernova remnant (cSNR) activity resulting from the evolution of a metal-rich massive stellar cluster, product of a starburst in the nucleus of an early-type galaxy. In this context, the optical light curves of AGN are reproduced by a random sequence of SN events. The parameters that describe a given light curve are the overall rate of explosions (I'SN), the energy released in each cSNR (esi), and the density of the circumstellar medium in which the remnants evolve (nr). In the case of low-luminosity AGN (MB <; —22 mag) these three parameters are well constrained by the observations: I>SN by the minimum of luminosity and/or by the number of peaks of the light curves; esi by the amplitude and duration of typical oscillations in the light curve and/or by the equivalent width of recombination lines, such as H/?; and 717 by the decay rate of well-isolated peaks (Aretxaga &; Terlevich 1993). The physics involved in the parameters above provides two independent constraints on the distance to a low-luminosity AGN.

1. Firts distance indicator: SN rate versus stellar luminosity The 5-band luminosity arising from a coeval cluster at its SN II explosion phase, 10 to 60 Myr, is mainly due to Main Sequence stars and cSNR. The SN explosion rate is related to the luminosity coming from the stars of the cluster (Lg) by ^

=(2±l)xlO-11yr-1LB0-1)

(1.1)

where the adopted value corresponds to a cluster with a solar neighbourhood Initial Mass Function (IMF), and the range represents the maximum deviations from this value for IMF with slopes between a = 2 and 3 (Aretxaga & Terlevich 1993). The best monitored low-luminosity AGN exhibit spectral transitions between types 1 (broad permitted lines) and 2 (no broad lines) of Seyfert activity. This is the well known case of NGC 4151, but also that of NGC 5548 and some other objects (see Aretxaga & Terlevich, these proceedings). At minimum light no strong broad lines are present and any contribution to the optical luminosity from cSNR should be small. The blue continuum at these minima originates mainly in the hot stars that still remain in the Main Sequence. One way of estimating the VSN of a low-luminosity starburst-powered AGN consists in counting the number of peaks of its optical light curve. The general evolution of cSNR is easy to identify since it involves energies ~ 1051 erg in time scales of a few years. Equation (1.1) can then be used to estimate the absolute luminosity of the photometric minima, which compared with the apparent magnitude gives a distance estimator. 2. Second distance indicator: amplitude of the variability The evolution of cSNR in high-density media has been described by Terlevich et al. (1992). The maximum blue luminosity arising in such a remnant can be estimated by LfNR « 2-9 x 10 9 L BG e 5 i 7 / 8 n7 /4 , where e51 is the energy of the cSNR in 1051 erg units 343

344

Aretxaga & Terlevich: Distance Indicators to Low-Luminosity AGN 1

•

/' o o o

Ho =,•75 = Km s" 1 Mpc y

1 Ho = 100 Km f' Mpc/1'

1 1 Ho = 50-Km s" Mpc"

CO

NGC 5548

/

>8

Ho = 2 5 -Kfti's"1 Mpc" 1

-

CO

_ . . . - - •

'

, . ' • ' ' . . - • ' • '

--T—

^

NGC 4151 i

50

100

150

200

d (Mpc) FIGURE 1. Recession velocities of the sources versus derived distances. The dotted boxes outline the range of values given by the first method, and the dashed ones those by the second method. The solid-line boxes give the range of values compatible with both methods. The dot-dash-dot lines trace the isotropic Hubble law for different Ho values.

and «7 is the circumstellar density in which the remnant evolves in 107 cm" 3 units. A good estimation of £51 can be made through the equivalent width of recombination lines, such as R/3 (Aretxaga & Terlevich 1993), (2.2)

and n-i can be constrained by the decay rate of isolated peaks. Therefore, the peak-to-peak variation of a well-isolated peak in a light curve gives the contrast of a single cSNR against the stellar luminosity. The stellar luminosity can be calculated and, hence, also the distance to the object. 3. The distances to NGC 4151 and NGC 5548 NGC 4151 and NGC 5548 are the best monitored Seyfert 1 galaxies (see Aretxaga k Terlevich 1993 for their 5-band light curves). Figure 1 shows the distances derived to these two nuclei applying the methods decribed above. Through these possible distance indicators, we can deduce that the stellar origin of the variability observed in AGN is compatible with the light curves of the best-monitored Seyfert galaxies only if the Hubble constant is in a range 25 ^ Ho ^ 75 km s" 1 Mpc"1. If one moves the absolute magnitude scale to a value outside this range, either the number of peaks observed in the light curves or the amplitude of the peaks is inconsistent with the minimum luminosity or with the mean equivalent width of H/?. IA acknowledges the Basque Government for finantial support through grant BFI93.009

REFERENCES ARETXAGA,

I. & TERLEVICH, R. 1993 Astrophys. Space Sti. 205, 69.

TERLEVICH, R.J. et al. 1992 Mon. Not. R. Astr. Soc. 255, 713.

Broad- and Narrow-Band Imaging of the CfA Seyfert Sample By A. M. PEREZ GARCIA AND J. M. RODRIGUEZ ESPINOSA Instituto de Astrofisica de Canarias, E-38200 La Laguna, Tenerife, Spain We are carrying out a deep study of the CfA Seyfert sample of Seyfert galaxies in broad-band and narrow-band Ha- Our aim is to perform a complete analysis of the morphological properties of the galaxies hosting the Seyfert nuclei. We will also study the location and number of circumnuclear star-forming regions, and the incidence of interactions and galaxy-wide starbusts, emphasizing the similarities and differences between the type 1 and type 2 objects.

1. Introduction Previous studies of active galaxies have been mainly concerned with the properties of the active nucleus. In fact, the many varieties of AGN recognised today were separated according to their nuclear properties. Yet the importance of the host galaxy in the understanding of the nuclear activity has been pointed out by several authors. For instance, some authors suggest that the host galaxies of Seyfert nuclei are substantially more luminous than similar field galaxies in the far infrared (Rodriguez Espinosa et al. 1987; Edelson et al. 1987; Rieke 1992), implying a connection between galaxy-wide star formation and Seyfert activity. Other authors find differences between the stellar formation rates in the host galaxies of Seyfert 1 and Seyfert 2 (Heckman et al. 1989). It has also been known for some time that Seyfert l's and 2's tend to be hosted by spiral galaxies of different morphological types, Seyfert 2's ocurring more often in latetype spirals than Seyfert l's. This fact is however not taken into account when unified AGN models are described. These models claim that it is just the orientation of the nuclei to the line of sight that accounts for the difference between types 1 and 2 objects (Acosta Pulido 1993; Antonucci k Miller 1985). In an effort to make an exhaustive study of the differences among the Seyfert types, and between these and the starbust objects, we are undertaking a detailed study of the morphology, colours, number of nuclear and non-nuclear ionizing photons and number and location of star-forming regions for the objects in the CfA sample of Seyfert galaxies. We will also look for relations between Seyfert type and optical and far-IR properties, such as the blue to far-IR luminosity ratio or the amount of nuclear versus non-nuclear Ha emission, parameters that will lead to a comprehensive understanding of the true similarities and differences between the two Seyfert types, and of the importance of the host galaxy in the properties of their nuclei.

Object sample We have chosen for this project the CfA Seyfert sample which is a complete subgroup of the velocity sample of the Center for Astrophysics (Huchra and Berg 1987). It consists of 25 Seyfert 1 and 23 Seyfert 2 sources. Selection criteria were the presence of emission lines in their spectra and an apparent visual magnitude brighter than 15.5. The sample 2.

345

346

Perez Garcia & Rodriguez Espinosa: The CfA Seyfert Sample

is uniformly distributed in the sky (according to V/Vm test), as expected for a complete sample free of selection effects. This sample will be observed in a variety of far-IR bands by the ISO satellite as part of its Central Programme of observations, in which the authors have guaranteed time for this project.

3.

Observations We are obtaining deep optical images in broad band (B, V, R and /) and narrow-band (Ha) plus adjacent continuum. For this project we are using the 2.5-m Nordic Optical Telescope at the ORM in La Palma. Good seeing conditions (< 1 ) have been usual during our observations so far. At present (October 93), 40 galaxies out of 48, i.e. 83% of the sample, have been observed in the broad bands. However, only 33% has been observed in Ha and continuum due to the lack of appropiate redshifted filters. We have ordered adequate filters to complete the sample and expect to have them ready for our next run in 1994. The data reduction process is fairly well advanced and we are now analysing the results. 4.

Objectives The main objective of this work is to study differences in the morphologies of the Seyfert 1 and 2 types as well as the incidence by type of companion galaxies, interactions or mergers and galaxy-wide starbursts. To this end we will • study the colours and luminosities of the nucleus, bulge and disc of the host galaxy, • compare the results with normal galaxies of similar Hubble type, • compare the host galaxies of Seyfert l's and Seyfert 2's, • study the stellar population, stellar formation rate, number of ionizing photons, etc, • find possible correlations between nuclear and host galaxy properties, and • study the influence of interactions in triggering nuclear activity and star formation. In addition to the thorough study outlined above, we will end up with a data base that will form the basis of a systematic study of the CfA sample of Seyfert galaxies to be performed with photometric and image data from ISO in a variety of mid- and far-IR bands. We expect to produce a statistically sound and definite study of the differences between the type 1 and 2 Seyfert galaxies and address the unified model approach from the viewpoint of the morphologies of the host galaxies.

REFERENCES ACOSTA

PULIDO, J. A. 1993 Astrophys. Sp. Sc. 205, 195.

ANTONUCCI, R. R. J. & MILLER, J. S. 1985 Ap. J. 297, 621. EDELSON, R. A., MALKAN, M. A. &: RIEKE, G. H. 1987 Ap. J. 321, 233. HECKMAN, T. M., BLITZ, L., WILSON, A. S., ARMUS, L. & MILEY, G. K. 1989 Ap. J. 342,

735. RIEKE, G. H. 1992 A. S. P. Conf. Ser. 31, p. 61. RODRIGUEZ ESPINOSA, J. M., RUDY, R. R. & JONES, B. 1987 Ap. J. 312, 555.

Type Transitions in Starburst-Powered AGN By ITZIAR ARETXAGA AND ROBERTO J. TERLEVICH Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, UK There is mounting evidence that type transitions are a common property of AGN: the broad lines in at least eleven Seyfert galaxies have appeared or disappeared, leading to the reclassification of their nuclei from type 1-1.5 to type 1.8-2 or vice versa. We show that these phenomena find a natural explanation in the starburst model for AGN as transient phases without supernova activity in a 10-60 Myr old metal-rich massive stellar cluster with a low supernova rate (VSN & 3 yr- 1 ).

1. Type transient AGN: casuistry Spectroscopic observations of Seyfert galaxies established early on that the broad permitted lines can experience strong variations in time scales of a few weeks-months. There is a growing number of extreme cases in which the broad components have temporarilly disappeared or become so weak that a reclassification of the objects has been allowed: from Seyfert nuclei of type 1-1.5 to type 1.8-2/LINER. Among them we canfindthe prototypes NGC 4151 and NGC 5548, along with NGC 1566, NGC 3516, NGC 6814, NGC 7603, Mrk 372 and 3C 390.3. In most of these cases, we know that the transitions took place while the nuclei were in deep photometric minima. Conversely, there are some narrow-line objects that have developed prominent broad components while brightening. Among them we can find Mrk 6, Mrk 993 and Mrk 1018. The characteristics of these nuclei are listed in Table 1. Figure 1 shows, as an example, the spectroscopic transition experienced by NGC 4151 during the photometric minimum of April 1984 (Penston & Perez 1984; Lyutyi, Oknyanskii k Chuvaev 1984), which is marked with an arrow in the attached light curve. 2. Variability in starbursts The evolution of metal-rich massive starbursts can mimic many of the observed characteristics of radio-quiet AGN (Terlevich et al. 1992; Terlevich & Boyle 1993, and references therein). The variability observed in these systems is thought to be produced by the supernova (SN) and compact supernova remnant (cSNR) activity of the cluster. cSNRs are the product of the interaction of the ejecta of SNe with the high-density circumstellar medium expelled by the progenitor stars. Detailed hydrodynamical models show that these systems convert most of the kinetic energy of the ejecta into radiation in time scales of a few years. The basic Broad Line Region (BLR) properties can be ascribed to the evolution of cSNRs in a medium with densities n > 107 cm" 3 and metallicities of the order of Z© or higher (Terlevich et al. 1992). The energy and the overall pattern of variability of well-sampled light curves of Seyfert galaxies, such as NGC 4151 and NGC 5548, can be modelled by a sequence of SN events (Aretxaga & Terlevich 1993,1994) and the detailed response of the BLR to continuum variations can also be explained by these phenomena (Terlevich et al. 1994). Even for the most luminous objects, the high-redshift QSOs, clue characteristics such as their variability and luminosity function (Terlevich & Boyle 1993) are naturally explained. 347

348

Aretxaga & Terlevich: Type Transitions in Starburst-Powered AGN

im

70

1500

IMO

FIGURE 1. Left panel: Spectral variation of NGC 4151 in 1984 compared with the normal state of the nucleus (from Penston & Perez 1984). Right panel: light curve of NGC 4151 from 1969 to 1986 (Penston &; Snijders, private communication); the arrow indicates the date of the transition of the nucleus to a Seyfert 1.9 type.

The 5-band luminosity arising from a young stellar cluster at its SN II explosion phase, 10-60 Myr, is mainly due to the contribution of Main Sequence stars and SNe. The SN rate (USN) and the luminosity coming from Main Sequence stars {L*B) are related along the lifetime of this phase by VSN/L*B » 2 x 10

-11

(2.1)

almost independently of the Initial Mass Function and age of the cluster (Aretxaga & Terlevich 1993,1994). ^From this expression, the total mean luminosity of the cluster, due to stars and cSNRs, is related to the SN rate by XB~1O1OI/SJV(C51 + 5)LB0)

(2.2)

51

where e$i is the energy of a cSNR in units of 10 erg. An estimate of €51 can be obtained from the mean equivalent width of recombination lines such as H/J,

40 A

£51

(2.3) 1 + 0.17C51' also independently of the mass and age of the cluster and, therefore, of the SN rate (Aretxaga & Terlevich 1994). Furthermore, the value of €51 obtained from this expression is also independent of the assumed cosmology. The constancy of the mean equivalent with of ftp in AGN (~ 80 A) reflects the universal value of the energy SN explosions have (~ 3 x 10 51 erg). Thus, the luminosity of an AGN is directly related to its SN rate, given the universal value of

Aretxaga & Terlevich: Type Transitions in Starburst-Powered AGN

Object NGC 1566 NGC 3516 NGC 4151 NGC 5548 NGC 6814 NGC 7603 Mrk6 Mrk 372 Mrk 993 Mrk 1018 3C 390.3

AfB(mag) >-18 -19.8 -20.0 -21.1 -16.0 -20.0 -19.9 -20.0 >-21 -18.2 -22.0

Minimum? Yes ?

Yes

Yes ? ?

Yes ?

Yes Yes Yes

349

Reference Alloin et al. 1986 Andrillat k Souffrin 1968 Penston &; Perez 1984; Lyutyi et al. 1984 Iijima et al. 1992 E. Terlevich et al. (priv. comm.) Tohline &; Osterbrock 1976 Khachikian &: Weedman 1971 Gregory et al. 1991 Tran et al. 1992 Cohen et al. 1986 Penston & Perez 1984

TABLE 1. Type-transient Seyferts. Column 1: name of the object. Column 2: MB, mean luminosity of the object (Ho = 50 Km s"1 Mpc"1 ) from Steiner (1981) or Whittle (1992). Column 3: whether or not the object was in a photometric minimum when it lost the broad lines. Column 4: report of the transition.

3. Type transitions in starbursts Young cSNRs originate BLRs like those observed in SN 1987f and SN 1988z, the prototypes of the "Seyfert 1 like SN" (Filippenko 1989; Stathakis & Sadler 1991; Terlevich et al. 1992). However, for low- and medium-luminosity AGN (MB & —22.5 mag), the low SN rates derived (VSN ^ 1 yr" 1 , from eq. 2.1) give non-negligible time scales for states in which no cSNRs could contribute to the existence of broad lines in the spectrum. This is illustrated with the synthetic light curves of Figure 2. Each double peak variation represents a simplified light curve of a single SN. The first sharp peak is produced in the outburst and the second one corresponds to the onset of the radiative phase of the associated cSNR. If we assume that SN explosions are random events within a given mean rate, the light curves of the corresponding clusters are those of Figure 2, in which the arrows indicate the moment of the explosions and the dashed lines the stellar luminosity levels (LB). According to the photoionization models for cSNRs of Terlevich et al. (1992), the equivalent widths of Ha and H/J change in time from WHO/^SN ^ 37.7 A yr, WHQ/VSN & 6.1 A yr for 4tsg to WHC/VSN « 13-1 A yr, WHp/vsN « 1.4 A yr for $tsg, where tsg is the characteristic time of evolution of the cSNR. If we adopt 20 A as the observable limit below which an object is classified as a Seyfert 1.9 (in which case WHQ & 20 A but Wjja ^t 20 A) or a Seyfert 2 (WHO, WH/3 £> 20 A), the transitions take place when the total light emitted by the cluster (solid line in Figure 2) is less than 0.14 and 0.01 mag above the stellar level (dashed line in Figure 2), respectively. The activity level is recovered once the light curve crosses those limits in the opposite direction due to a new cSNR. Using these Monte Carlo simulations we can estimate the time spent by these clusters as Seyfert nuclei of type 1.9 or 2. The values obtained are just upper limits to the time spent in quiescent stages, since our approach ignores the secondary pulses that occur in the evolution of the cSNR due to cooling instabilities. The bottom panel of Figure 3 represents the fraction of time spent as type 1.9 and type 2 Seyferts by clusters with mean B-band luminosities between —16 and —23 mag. The less luminous systems are the ones that experience longer quiescent stages. There is a theoretical upper limit for the luminosity a transient system can have: MB » -22.5 mag. Clusters with luminosities

350

Aretxaga k Terlevich: Type Transitions in Starburst-Powered AGN

0 2. Theoretical light curves of massive stellar clusters undergoing the indicated SN rates (Aretxaga 1993). The arrows in the diagram represent the moments in which the SNe explode and the dashed lines the luminosity level coming from Main Sequence stars. FIGURE

Aretxaga & Terlevich: Type Transitions

NGC 6814

«_NGC 1566

u ft Mrk 1018

in Starburst-Powered

AGN

351

.^.NG£.,3516

ft 3C 390.3 ft NGC 5548 NGb 4151 <_ Mrk 993 NGC 7603

CU

6. 372

o o

Seyfert 1.9 o

Seyfert 2

CM

-16

18

-20

-22

MB (mag) FIGURE 3. Bottom: Fractions of time spent as Seyfert of type 1.9 or type 2 by the stellar clusters of Figure 1. The dashed line describes the transition to type 1.9 stages and the dotted line that to type 2 stages. Top: Number of transient AGN per luminosity interval. The symbols represent the luminosities of these objects. The stars represent the mean luminosities of the nuclei, and the arrows indicate upper limits to these quantities.

exceeding this limit have excessively high SN rates and, therefore, they have a very low probability of showing sporadic quiescent stages. In order to test the above prediction, we have searched in the literature for all the reported transient AGN. These objects are listed in Table 1, and their luminosities are represented in the top panel of Figure 3. The stars represent the mean luminosity of the nuclei and the arrows represent upper limits to these quantities. All the known type-transient AGN have luminosities below the cutoff predicted by the theory. The objects tend to clump in the region around —20 mag, and there is an obvious decline of the number of objects towards higher and lower luminosities. The high-luminosity zone is restricted by the lower probability of occurrence of the transitions, while the lower-luminosity zone is probably unpopulated since these sources are believed to be non-variable and are therefore hardly monitored.

4. Conclusions In the framework of the starburst model for AGN, the type transitions observed in many Seyfert nuclei are expected to occur in systems with MB *S —22.5 mag. In these clusters there is a non-negligible probability that the broad components of the emission

352

Aretxaga & Terlevich: Type Transitions in Starbursi-Powered AGN

lines are undetectable at some epoch. The quiescent stages, with undetectable broad lines, occur at deep photometric minima, while the presence of prominent broad lines is predicted to occur during maxima in the light curve. The transitions are produced due to the stochastic nature of the SN explosions in the cluster, which produces periods of time in which there is a lack of new explosions, while the existing remnants are too old to produce broad lines. The continuum in these epochs is still blue because the light is dominated by the young stars that are still on the Main Sequence. The radio-quiet AGN in which activity transitions have been observed are all below the predicted luminosity cutoff. The only radio galaxy known to have changed its type of activity (3C 390.3) is also below this limit. IA thanks the Basque Government for financial support through grant BFI93.009.

REFERENCES D., PELAT, D., PHILLIPS, M. M., FOSBURY, R. A. E. &; FREEMAN, K. 1986 Astiophys. J. 308, 23. ANDRILLAT, Y. & SOUFFRIN, S. 1968. Astrophys. Lett. 1, 111. ARETXAGA, I. 1993 PhD thesis, Universidad Autonoma de Madrid. ARETXAGA, I. & TERLEVICH, R. 1994 Mon. Not. R. Astron. Soc. Submitted. ARETXAGA, I. & TERLEVICH, R. 1993 Astrophys. Space Sci. 205, 69. ARETXAGA, I. et al. 1994 In The Nature of Compact Objects in AGNs. Cambridge University Press. In press. ARETXAGA, I. et al. 1993 In First Light in the Universe: Stars or QSOs? (ed. Rocca-Volmerange et al.), p. 331. Editions Frontieres. COHEN, R. D., RUDY, R. J., PUETTER, R. C , AKE, T. B. &; FOLTZ, C. B. 1986 Astrophys. J. 311, 135. FILIPPENKO, A. V. 1989 Astron. J. 97, 726. GREGORY, S. A., TIFFT, W. G. & LOCKE, W. J. 1991 Astron. J. 102, 1977. IuiMA, T., RAFANELLI, P. &: BlANCHiNi, A. 1992 Astron. Astrophys. 265, L25. KHACHKIAN, E. YE. &; WEEDMAN, D. W. 1971 Astrophys. J. Lett. 164, L109. LYUTYI, V. M., OKNYANSKII, V. L. &; CHUVAEV, K. K. 1984 Sov. Astron. Lett. 10, 335. PENSTON, M. V. & PEREZ, E. 1984 Mon. Not. R. Astron. Soc. 211, 33P. PETERSON, B. M. et al. 1994 Astrophys. J. In press. STATHAKIS, R. A. & SADLER, E. M. 1991 Mon. Not. R. Astron. Soc. 250, 786. STEINER, J. E. 1981 Astrophys. J. 250, 464. TERLEVICH, R. J. & BOYLE, B. 1993 Mon. Not. R. Astron. Soc. 262, 491. TERLEVICH, R. J. et al. 1994 Mon. Not. R. Astron. Soc. In press. ALLOIN,

TERLEVICH, R. J.,

TENORIO-TAGLE, G.,

FRANCO, J. & MELNICK, J. 1992

Astron. Soc. 255, 713. TOHLINE, J. E. & OSTERBROCK, D. E. 1976 Astrophys. J. Lett. 210, L117. TRAN, H. D., OSTERBROCK, D. E. & MARTEL, A. 1992 Astron. J. 104, 2072. WHITTLE, M. 1992 Astrophys. J. Suppl. 79, 49.

Mon.

Not.

R.

Stellar Ionization of Low-Luminosity Active Galactic Nuclei ByJOSEPH C. SHIELDS Steward Observatory, University of Arizona, Tucson, AZ 85721, USA

Low-Ionization Nuclear Emission-Line Regions (LINERs) are a common constituent of galaxies, and are often regarded as a weak form of Seyfeit activity. LINERs have emission-line luminosities that are similar to those of giant HII regions, however, and recent theoretical work suggests that their nebular properties can be reproduced in many cases with photoionization by normal O stars. In an extension of this scenario, energetic phenomena such as non-thermal radio emission, broad Ha features, and substantial X-ray luminosity seen in some LINERs might be attributable to supernovae. In this review I consider the empirical evidence bearing on an interpretation of LINERs as stellar-powered sources. While stellar phenomena appear capable of matching LINERs of modest luminosity in terms of broad-band energetics, some important differences remain in detailed spectral characteristics, particularly at X-ray energies. A certain amount of anomalous behavior on the part of stars within galaxy nuclei (e.g. in terms of the initial mass function) is required if LINERs result from stellar ionization.

1. Introduction Spectroscopic surveys indicate that ~ 30% or more of bright galaxies contain weak emission-line nuclei that are classified as LINERs (Low-Ionization Nuclear EmissionLine Regions; Heckman 1980b, hereafter H80b). The intensity ratios of low-ionization optical lines relative to recombination features are characteristically higher in LINERs than in HII regions, causing LINERs to be regarded as "active" nuclei subject to unusual energetic processes. The luminosities of LINERs and giant HII regions are comparable, however, with typical Ha luminosity ~ 1040 erg s"1 or somewhat less. This fundamental resemblance provides motivation for examining the possibility that LINERs and HII regions are actually powered by a common energy source, namely a population of hot stars. Under this scenario, the observational distinctions between LINERs and giant HII regions must be accounted for in terms of differences in the associated stellar populations and/or in the properties of the nebular gas. This contribution will review evidence supporting and opposing an interpretation of LINERs as the products of stellar ionization. The conjecture that LINERs are ultimately stellar phenomena is closely related to work by Roberto Terlevich and collaborators that interprets active galactic nuclei (AGN) more generally as anomalous starbursts. The starburst-AGN scenario has been criticized on the basis of the requisite (high) efficiencies of broad-band energy generation (e.g.Heckman 1989), but these arguments are largely obviated when only low-luminosity nuclei are considered. LINERs differ from Seyfert nuclei and QSOs in terms of their optical and X-ray characteristics as well as typical luminosity, and hence it may be physically meaningful to consider LINERs apart from the remainder of active nuclei, as stellar-ionized sources. For purposes of this discussion, LINERs are defined according to criteria suggested by H80b, such that oxygen line intensities are in a proportion of ([ONI] A5007 : [Oil] A3727 : [01] A6300) = (1 : > 1 : > 0.3). A Hubble constant of Ho = 75 km s~l Mpc"1 is assumed throughout this review. 353

354

Shields: Stellar Ionization of Weak AGN

2. LINER ionization mechanisms Several mechanisms are capable of generating plasma with the ionization and thermal conditions necessary for producing LINER-like emission. The general requirement is gas described at least in part by densities n e £ 103 cm" 3 (dictated by the critical densities of the [ON] A3727 and [SII] A6724 line pairs), subject to weak ionization, while maintaining a rather hot electron population responsible for collisional excitation and resultant forbidden-line emission. Shocks with velocities of order ~100 km s" 1 provide one mechanism for generating such plasma (e.g. Shull & McKee 1979). Shocks associated with large-scale winds appear to be responsible for LINER-like emission observed in some starbursting galaxies (Heckman et al. 1990). Photoionization can also generate LINER nebulosity (Ferland & Netzer 1983; Halpern k Steiner 1983). Successful photoionization models require a low ionization parameter (U; ratio of ionizing photon to nucleon density) compared to that describing HII regions and Seyfert nuclei; in addition, the ionizing continuum must be relatively hard in order that the heating per ionization be large enough to produce strong forbidden-line emission. LINER-like emission is seen in the central galaxies of X-ray cooling-flow clusters (Heckman et al. 1989 and references therein). This nebulosity is probably generated by photoionization by the cooling intracluster plasma (Donahue k. Voit 1991) or by energy transfer in turbulent mixing layers at the interface of the hot and warm media (Begelman & Fabian 1990). This review will focus on compact nuclear sources of LINER emission, which constitute the most common manifestation of such nebulosity. This selection avoids most cases in which the nebular gas is powered by starburst winds or cluster cooling flows. The relative role of shocks and photoionization for energizing compact LINERs remains ambiguous, however. A number of LINERs resemble Seyfert nuclei by exhibiting significant X-ray luminosity and a broad component to Ha. Strong evidence exists (e.g. correlated continuum and line variations) implying that Seyferts are primarily photoionized. By extension, it can thus be argued that compact LINERs are similarly dominated by photoionization. This argument is plausible but not robust; for example, photoionization may only dominate the weak broad-line components of LINERs, while the narrow-line emission could be generated largely in shocks. Observational diagnostics of the ionization mechanism in LINERs, such as the [OIII] temperature or ratios of oxygen and sulfur lines, have been confounded by density effects and other ambiguities (e.g.Filippenko & Halpern 1984; Kirhakos & Phillips 1989). It is also unclear whether existing shock models satisfactorily account for the detailed emission spectra of real shocks (see Ho et al. 1993). In my view, the importance of shocks in generating LINERs remains an open question that is in danger of being neglected. The increasing emphasis on photoionization and neglect of shocks stems at least in part from the lack of a well-defined geometrical picture of how shocks are produced and sustained in a galaxy nucleus. With this caveat, the remainder of this discussion will nonetheless emphasize photoionization processes for generating most of the optical emission. As noted above, several observational characteristics suggest continuity in the properties of LINERs and Seyfert nuclei. These features include X-ray emission, which for LINERs is seen in ~ 10 objects to date. X-ray luminosities range from a few times 1040 erg s" 1 (e.g. M81; Petre et al. 1993) to more than 1043 erg s" 1 (Pictor A; Singh, Rao & Vahia 1990); values near the lower end of this range (~ 1041 erg s"1) seem typical. There is some tendency for LINERs detected as X-ray sources to exhibit other signs of Seyfert-like activity, including weak emission of broad Ha or of high-ionization lines (e.g. [NeV] A3426), or relatively strong non-thermal radio emission; in this sense, X-raydetected sources may be more "active" overall than the majority of LINERs. Broad Ha

Shields: Stellar Ionization of Weak AGN

355

emission, which is a hallmark of Seyfert 1 activity, is seen more generally in an important fraction of LINERs (Filippenko k Sargent 1985). Optically identified LINERs tend to be associated with weak, AGN-like, nonthermal radio cores, with 5-GHz luminosities of £ 1027 erg s" 1 Hz"1 (H80b). LINERs as well as Seyferts preferentially occur in earlytype galaxies (H80b). Finally, both LINERs and Seyfert nuclei often show a correlation between forbidden-line velocity width and critical density (e.g. Pelat, Alloin k Fosbury 1981; Filippenko k Halpern 1984). The last correlation implies that these nuclei are structured in such a way that plasma moving at the highest velocities also tends to be described by the highest densities. In light of these similarities, it is understandable that LINERs are often viewed as diminutive forms of Seyfert/QSO activity. Significant continuity also exists between the properties of LINERs and HII nuclei, however. By HII nuclei I refer simply to emission-line nuclei of galaxies that spectroscopically resemble HII regions in other environments. HII nuclei are presumably powered by normal O stars, although these nebulae exhibit some peculiarities relative to typical disk HII regions (Kennicutt, Keel k Blaha 1989, hereafter KKB). LINERs and HII nuclei are similar in Ha luminosity (Heckman 1980a,b; HII nuclei may actually be slightly more luminous on average, cf. Hummel et al. 1990). The nebular gas in both types of nucleus tends to be characteristically denser and probably more highly clumped than in typical HII regions (KKB; Filippenko k Halpern 1984). HII nuclei and (to a greater degree) LINERs both exhibit elevated [SII] and [Nil] line strengths that are not easily interpreted as simply signatures of enhanced heavy-element abundances in otherwise normal HII regions (KKB). Both classes of nuclei also exhibit Ha/radio continuum ratios are an order of magnitude lower than those typical of disk HII regions. For the HII nuclei, the difference is unlikely to be entirely due to extinction of Ha; an extra non-thermal radio component is apparently present in many of these nuclei, as in LINERs. The peculiarities of HII nuclei are open to several interpretations. KKB suggested that essentially all emission-line nuclei contain an AGN that is presumably non-stellar. In low-luminosity sources, the AGN produces LINER-like emission. HII nuclei are then composite sources in which a weak LINER and circumnuclear HII regions are spatially unresolved in the observations; the AGN is thus responsible for any anomalies. The nuclei of galaxies are unusual environments for reasons other than the putative presence of a massive black hole, however, and some of these other aspects may influence the character of the emergent radiation. Galaxy nuclei tend to be metal-rich in comparison to extranuclear sites. Observational evidence exists suggesting that the interstellar dust within galaxy nuclei is unusual, probably in the sense of having larger characteristic grain size (Laor k Draine 1993; Lauer et al. 1993). As the center of a galaxy's potential well, the nucleus is likely to be subject to unusually high interstellar pressures and densities. Pressures that are two orders of magnitude or more higher than for our local interstellar medium have been established observationally for the centers of some early-type galaxies (Thomas et al. 1986) and of the Milky Way (Spergel k Blitz 1992). Gas within the nuclei of galaxies may also be subject generally to an unusual degree of shear and turbulence, as is the case for the Galactic center (Giisten et al. 1987). All of these characteristics may be important for modifying the emission of nebular gas located in such an environment.

3. Stellar ionization of LINERs 3.1. Theoretical picture

If LINERs are ionized by stars, the effective temperature Tejf of the stars must be fairly high. Terlevich k Melnick (1985) proposed that massive stars forming in galaxy nuclei

356

Shields: Stellar Ionization of Weak AGN

would experience enhanced mass loss and stripping due to their high metallicity, and would subsequently evolve into extreme Wolf-Rayet sources that they labeled "Warmers." They used stellar evolution and atmosphere models then available, to argue that the composite spectrum of a population of Warmers and unevolved O stars would resemble a power law through the extreme ultraviolet bandpass. This ionizing spectral energy distribution meets the requirements for photoionization of LINERs and more luminous AGN, thus providing the basis for the starburst-AGN theory as originally framed. The high metallicities required for production of Warmers are apparently not attained outside of galaxy nuclei, making the existence of Warmers difficult to test empirically. In a reanalysis of the Warmer problem, Leitherer, Gruenwald & Schmutz (1992) repeated the calculations of Terlevich & Melnick using updated stellar evolution and atmosphere models. Use of the new models results in a significantly softer ionizing continuum for a stellar population expected in galaxy nuclei. These results cast doubt on whether Warmers are important sources of ionizing radiation within AGN. (In the more recent incarnations of the starburst-AGN model, supernovae actually play a more important role than Warmers in generating the observed emission; e.g. Terlevich et al. 1992.) While the existence of Warmers as originally envisioned is thus in question, in resolving this issue we remain dependent on the stellar modelers. Recent stellar evolution and atmosphere calculations are presumably improvements on older work, yet an enlarged role for Warmers may possibly re-emerge in future predictions as more physics is included in the models. Recent photoionization calculations suggest that the optical emission-line characteristics of many LINERs can actually be reproduced through photoionization by normal O stars (Filippenko & Terlevich 1992; Shields 1992). This scenario requires fairly hot stars (Tefj it 45000 K), and in addition, a reduction in U by an order of magnitude from that characteristic of normal HII regions, for at least a part of the irradiated plasma. The high interstellar densities expected and measured for galaxy nuclei may provide a natural means for reducing U, which will be inversely related to n e for a fixed radiation field. The best match to observed LINER spectra results if the nebular gas is described by a range of ne and U, since the emissivity in different lines will be maximized under different conditions. Emission in high critical-density transitions is enhanced in high-density gas, since collisional quenching of low critical-density cooling transitions (particularly infrared fine-structure lines) leads to a higher nebular temperature and greater excitation. If LINERs are the product of hot stars in dense gas, we can make qualitative sense of several patterns of nuclear emission in terms of a picture in which large galaxy bulges are described by deep central potentials resulting in high central gas pressures and densities. Among these trends is a tendency for low-luminosity nuclear emission in early-type galaxies to appear LINER-like, while late-type galaxies tend to host HII nuclei (H80b). Among early-type galaxies alone, nuclear emission regions in large galaxies are usually LINERs, while small galaxies feature HH-region emission (Phillips ei al. 1986). There is also some tendency for the size of nuclear emission regions to increase from early to late Hubble types (Pogge 1989); a similar pattern is seen in the dimensions of nuclear radio sources (Giuricin et al. 1990; Hummel et al. 1990). In terms of the simple scenario considered here, the variation of nebular size reflects the dependence of Stromgren depth on density, while the same trend in radio dimensions results from increased confinement of the non-thermal plasma in denser environs. Finally, recent observations indicate that the nuclear J, H, K-b&nd (stellar) luminosity is larger in LINERs than in HII nuclei, independent of galaxy luminosity and Hubble type (Giuricin et al. 1993). This trend again may be indicative of deeper central potentials for the LINERs.

Shields: Stellar Ionization of Weak AGN

357

3.2. Direct evidence for hot stars If LINERs are powered by hot stars, we might hope to see signatures of a young population in the stellar continuum of these nuclei. At optical wavelengths, some LINERs do show significant Balmer absorption indicative of intermediate spectral-type (B-A) stars, although this is not so for the majority of sources (e.g. Heckman 1980a). The absence of these features may require that star formation within nuclei be described by an unusually shallow initial mass function (IMF) if LINER emission is stellar-generated, although the quantitative comparisons needed for this statement have not been done. Terlevich, Diaz & Terlevich (1990) have argued that large near-infrared Call absorption equivalent widths observed in Seyfert nuclei are probably indicative of a young red supergiant population. Measured EW(Call) in such objects is inconsistent with simple dilution of a normal galaxy spectrum by a power-law continuum inferred from absorption features at shorter wavelengths. Similar EW(Call) are observed in LINERs (Terlevich et al. 1990); however, little evidence is seen for dilution of any stellar features in most LINERs. EW(Call) in LINERs are consistent with those of normal galaxy nuclei. In early-type galaxies of the sort that typically host LINERs, nuclear EW(Call) also appears to connect fairly smoothly to gradients in this feature observed at larger radii, which probably reflect average stellar metallicity as a function of radius (Delisle & Hardy 1992). The ultraviolet bandpass is likely to be the best place to hunt for evidence of young stars in nuclei, due to the relatively high contrast between young and old populations in this wavelength interval. Ultraviolet spectra from WE exist for over a dozen LINERs, and in a few cases show evidence for hot stars (Goodrich & Keel 1986; Reichert et al. 1993). In general, however, the data quality is limited, and upper bounds to hot star numbers are relatively weak, particularly when possible extinction is considered. Evidence for a power-law continuum through the UV region in LINERs is also somewhat ambiguous (Fosbury et al. 1981; Bruzual, Peimbert & Torres-Peimbert 1982). Elliptical galaxies commonly show a rise in the continuum shortward of ~ 2000 A (Burstein et al. 1988), but this light probably originates in a population of evolved stars that contributes little ionizing radiation (Te/f ,$ 30000 K; Ferguson & Davidsen 1993). 3.3. Hot stars in metal-rich environments The scenario for stellar ionization of LINERs outlined in §3.1 is not without potential problems, including the necessity of early O stars in relatively enriched centers of galaxies. Studies of HII regions in disk environments reveal an inverse relation between metallicity and Tejf of the ionizing stars, where both quantities are inferred from the strengths of nebular lines (e.g. McCall, Rybski & Shields 1985). At metallicities of solar levels or higher expected for galaxy nuclei, ionizing stars with Tejj ,£, 45000 K are unexpected on the basis of this trend. The physical origin of the correlation for HII regions is poorly understood. One possible explanation is that the ionizing radiation emergent from the most massive stars is systematically softened by changes in stellar structure, atmospheres, or winds accompanying increases in metallicity. The latest generation of stellar evolution calculations appears to support this interpretation; see discussions by McGaugh (1991) and Garcia Vargas (this conference). This conclusion is again only as robust as the current stellar models, however. An alternative explanation is that the IMF for star formation responds to metallicity in such a way that fewer high-mass stars form in enriched media (e.g. Shields & Tinsley 1976). The physical basis for this behavior may lie, for example, in the role of heavy elements in cooling pre-stellar gas clouds, which could lead to more efficient fragmentation and collapse, and hence less massive stars, when abundances are high. If variations in the IMF are important in producing the observed metallicity-Te// relation, star-forming

358

Shields: Stellar Ionization of Weak AGN

regions in galaxy nuclei might deviate from this trend if other factors influence the IMF in these locales. Star formation biased to high masses is not an altogether ad hoc proposition for galaxy nuclei. Shearing and turbulent cloud motions could suppress low-mass star formation requiring a quiescent environment, while high densities and violent compression resulting from cloud collisions could facilitate collapse of large masses; this physical scenario is qualitatively consistent with the picture of star formation mediated by magnetic field strength discussed by Shu, Adams & Lizano (1987). Some empirical encouragement for invoking modifications to the IMF in LINERs might be taken from results presented at this meeting by E. Malumuth, which show that the highly compact stellar core of 30 Doradus is described by a relatively flat high-mass IMF compared with standard IMFs or that measured for the remainder of the HII region. If the relevant physical conditions of the 30 Dor core are replicated in galaxy nuclei, preferential formation of high-mass stars would not be surprising. 4. Energetic phenomena and supernovae Normal stars per se do not provide a simple explanation for Seyfert-like signatures broad Ha, nonthermal radio emission, significant X-ray luminosity, high-ionization lines - found in some LINERs. An extra ingredient must be invoked, and one possibility consistent with stellar ionization is the supernovae (SNe) that will result from a population of young stars. This explanation is essentially the same as is employed in the starburstAGN model, although by limiting our view to low-luminosity sources we can again avoid many possible objections and utilize known astronomical entities with less extrapolation into unknown physical regimes. 4.1. Theoretical picture The stellar explosions that empirically come closest to producing the requisite features of LINERs are a class of radio-loud Type II SNe (Weiler et al. 1986). Broad Ha emission is a general characteristic of Type II SNe, but only a small fraction of these events are strong radio sources. The radio emission is interpreted as synchrotron emission generated in an interaction between the expanding SN and dense circumstellar material deposited by a wind from the progenitor star (Chevalier 1981, 1982). Terlevich et al. (1992) have argued that such SNe should be the norm in galaxy nuclei since dense interstellar gas will tend to confine wind ejecta to a dense shell. The supernova ejecta and circumstellar material that have been shocked will both emit thermal X-rays, with the supernova plasma dominating at soft energies, while the circumstellar matter will radiate most of its energy in photons as hard as 100 keV. The cooler supernova gas is expected to dominate the X-ray spectrum, although the detailed proportions depend on density structure in the supernova shell and the wind. Increasing the circumstellar density is expected to increase emission at both radio and X-ray energies, thus providing a theoretical escape if extranuclear SNe that we observe are somewhat underluminous compared to the galaxy nuclei. While higher circumstellar densities may be the norm for evolved stars in galaxy nuclei, the stellar-ionized scenario would obviously be more satisfying if SNe are found that match LINER characteristics without theoretical mediation. 4.2. Empirical comparisons Extranuclear radio-loud supernovae are rare events, and only a small number of examples have sufficient observations to provide useful comparisons with LINER properties. While

Shields: Stellar Ionization of Weak AGN

359

a SN can outshine its host galaxy near maximum light, the duration of this peak in output is expected to be considerably shorter than the interval between SN explosions in LINER nuclei (§4.2.4). As a consequence, the most recent SN seen in a stellar-ionized LINER should, on average, be in an early remnant phase. In the following empirical comparisons I will consequently emphasize late-time observations of modern SNe. The detailed consistency of using SNe to explain high-energy phenomena in LINERs requires additional consideration of the temporal behavior of SNe, as discussed in §4.2.4. 4.2.1. Radio emission The SNe of interest emit non-thermal radio emission with 5-GHz luminosities of up to ~ 1027—1028 erg s" 1 Hz"1 at peak brightness (typically within a few years post-explosion; Weiler et al. 1986,1990), which is in good agreement with LINER luminosities. The radio spectra of the SNe are inverted (a > 0, /„ oc v+a) at low frequencies at early times, with a turnover to a < 0 above a transition frequency that decreases with time. This behavior is interpreted in terms of time-dependent free-free absorption by circumstellar material. The spectral index at late times is typically ~ —0.6; the resulting range in spectral index is generally consistent with spectral slopes found for LINERs (e.g. Heckman, Balick & Crane 1980). A minority of LINERs are considerably more powerful radio sources with extended, double-lobed radio morphologies (e.g. Pictor A; Jones & McAdam 1992 and references therein); it seems highly unlikely that the radio luminosity of these sources can originate in normal stellar phenomena. 4.2.2. Optical emission Near maximum light, the optical spectra of Type II SNe are dominated by broad emission in hydrogen Balmer lines. Three radio-loud Type II SNe have been detected in Ha at much later times: SN1970G (Fesen 1993), SN1979C (Fesen & Matonick 1993), and SN1980K (Fesen & Becker 1990; Leibundgut et al. 1991), observed 22, 12, and 8 years post-explosion, respectively. All three exhibited Ha in emission with full-width near zero intensity in excess of 10000 km s" 1 , which is comparable to Ha widths seen in LINERs. The luminosity of broad Ha for the three SNe was 1037 — 1038 erg s" 1 , with no simple dependence on remnant age. While quantitative estimates of broad Ha luminosity for LINERs are sparse (due to the difficulties of deblending this component) and biased to high-luminosity sources, available numbers tend to be at least somewhat higher than the SNe values. M81, for example, emits at least 5 x 1038 erg s"1 in broad Ha (Filippenko & Sargent 1988), which lies at the lower end of quantitative measurements. Better information on typical LINER Ha luminosities that is appropriate for this comparison will soon be available from the Palomar Observatory Dwarf Seyfert Survey (Filippenko
360

Shields: Stellar Ionizalion of Weak AGN 1

560 and 1700 km s" , respectively) that is likely to be inadequate for explaining the broad features observed in galaxy nuclei (e.g. M81 displays broad Ha FWHM = 2200 km s" 1 ; Filippenko k Sargent 1988). 4.2.3. X-ray emission A relatively new observational development is the detection of these last two SNe SN1978K (Ryder et al. 1993) and SN1986J (Bregman k Pildis 1992) - as X-ray sources approximately 13 and 9 years post-explosion, respectively. Both objects exhibited 0.1 2.5 keV luminosities of ~ 10 40 erg s" 1 , which is comparable to the lower end of detected LINER luminosities. Spectroscopically, however, these SNe differ significantly in Xray properties from the galaxy nuclei. When fit by a power law over a bandpass of 0.1 — 2.5 keV, the spectrum of SN1986J is described by a spectral index of a = —2.1 (although a thermal continuum provides a somewhat better fit; J. Bregman 1993, private communication). SN1978K shows a considerably steeper X-ray continuum, with a = —4.5 (Ryder et al. 1993). Both slopes are significantly steeper than the average for LINERs of a « - 1 . 0 (Mushotzky 1993). 4.2.4. Temporal and statistical consistency The comparisons given above demonstrate that known SNe come close to replicating some general high-energy characteristics of LINERs, although important discrepancies remain in the details. The value of these comparisons is limited by the small number of sources on which they are based. Additional optical and X-ray data for radio-loud supernovae, as well as quantitative measures of broad Ha and X-ray emission in LINERs, would help to define what "typical" behavior is for both types of object. Even if a subset of supernovae is observed to reproduce LINER-like properties in luminosity and spectral detail, good statistics for the LINERs and temporal information for the SNe will be necessary to demonstrate that the duty cycle for observable SN features is consistent with the incidence of high-energy features in LINERs. A simple calculation using approximate numbers illustrates how the latter comparison would proceed. If we suppose that a typical LINER has a (narrow) Ha luminosity of ~ 1040 erg s" 1 , powering this nebula would require ~ 10 3 O stars. Since the lifetimeof an O star is of order 107 yr, in asteady state we would then expect a LINER to produce 1 SN every 104 yr, on average. If 10% of LINERs exhibit a broad Ha component, then a SN must be capable of sustaining a corresponding luminosity in broad Ha for ~ 1000 yr. A duration of this length may exceed that permitted by reasonable thermalization efficiencies; however, a shorter timescale might result from more accurate numbers and consideration of IMF effects in relating ionizing luminosity to SN rate. Analogous calculations could be done for radio and X-ray emission. At late times, SNe generally evolve slowly in their spectral properties. However, observations of a LINER bracketing a nuclear SN explosion in time would be expected to show significant temporal variation. Very limited information is available for variability of sources with the low powers typical of LINERs. Relatively detailed observations exist available for NGC 4278, which brightened by ~ 70% at 5 GHz between 1980 and 1985 to reach a maximum luminosity of ~ 4 x 10 28 erg s" 1 (Wrobel k Heeschen 1991), and subsequently began to decline in power. Comparison of the radio spectra in 1980 and 1985 shows a greater brightening at higher frequencies, which is qualitatively consistent with addition of a SN that is initially optically thick to free-free absorption at low frequency. Detailed modeling would be necessary to see if the quantitative behavior and timescales for the radio variations are consistent with a SN interpretation. At late times, radio luminosity from SNe is usually observed to decay with a power-law dependence on time (typically t~° 7 ; Weiler et al. 1986). More complicated temporal variation in luminosity

Shields: Stellar Ionization of Weak AGN

361

and spectral shape can result from structure in the circumstellar medium (Weiler et al. 1990, 1992). Strong optical variability has been seen in one LINER, NGC 1097, which changed over a few years from having no detectable broad Balmer emission to showing strong, broad features at both H/? and Ha (Storchi-Bergmann ei al. 1993). The broad emission is noteworthy for having a double-peaked profile similar to Balmer features in other AGN that have been interpreted as a product of an accretion disk or jets. Double-peaked profiles have not been observed in the spectra of SNe, although the broad Ha profiles in late-time SN spectra are sometimes notably boxy (e.g. Fesen 1993). NGC 1097 has a compact nonthermal radio core with 5-GHz luminosity of ~ 10 27 erg 1 s" Hz" 1 (Hummel, van der Hulst & Keel 1987), which is typical of LINERs. The galaxy also emits an X-ray luminosity within 0.2-4.0 keV of 1.1 x 10 41 erg s - 1 (Fabbiano, Kim & Trinchieri 1992, based on measurements in 1979). The narrow-line fluxes quoted by Storchi-Bergmann el al. are actually inconsistent with classification of NGC 1097 as a LINER, if we stick to the current definition (§1), and are perhaps in better agreement with a Seyfert 2 designation. However, observations of the same object by Phillips et al. (1984) are fully consistent with a LINER classification. The narrow emission-line equivalent widths in NGC 1097 are small, and discrepancies between the two studies are probably traceable to slight differences in subtraction of the stellar continuum or the use of slightly different measurement apertures. Whether a specific choice of classification in this case is physically meaningful for our purposes is open to question. X-ray variability on timescales of years is well documented for several LINERs (Mushotzky 1993). One LINER, M81, has also exhibited rapid X-ray variability on timescales of minutes (Barr et al. 1985), and such behavior is often taken as evidence in support of the black hole paradigm for AGN. Terlevich et al. (in preparation) have argued that fast shocks in dense, clumpy SN ejecta could give rise to this kind of variability. No variability on short timescales has been reported in the very limited X-ray data for SNe.

5. Summary and future prospects The preceding observational comparison demonstrates reasonable agreement between a number of LINER properties and stellar-powered phenomena in terms of overall energetics, but important differences remain between detailed aspects of the spectral energy distribution for the two source types. A model in which LINERs are powered by stellar phenomena evidently still requires a certain amount of anomalous behavior on the part of stars inhabiting galaxy nuclei. The mismatch of empirical properties and the speculative character of stellar-powered LINERs might both be reduced by additional observational and theoretical work. The kind of observations that are desirable for this purpose are mostly obvious. Supernovae observed to date show characteristic patterns of radio spectral evolution, and it would be of interest to see if similar behavior occurs in LINERs. LINER variability data at all wavelengths would be valuable more generally for understanding the origins of nuclear activity in these sources. Such observations have utility beyond comparison with SNe, and may allow tests of other physical scenarios and paradigms. Improved and more extensive ultraviolet spectroscopy of LINERs using the Hubble Space Telescope would help to place direct constraints on the presence of hot stars or a featureless continuum, and would also provide nebular diagnostics of the continuum shape. Imaging of LINERs with HST would also be of interest for constraining the size of the nuclear source; an example already exists in the case of M81, in which the nuclear H/? emission region is smaller than 0.12 pc (Crane et al. 1993). A size this small does not

362

Shields: Stellar Ionizaiion of Weak AGN

preclude generation of the LINER by stars - 30 Doradus and probably M33 (which does not have a nuclear black hole with mass in excess of 5 x 104 M©; Kormendy k McClure 1993) have stellar cores that are smaller. However, such measurements place potentially interesting constraints on the physical characteristics of a nuclear star-formation environment. Additional X-ray studies of LINERs and SNe would be informative, and data at relatively hard energies (as can be obtained with ASCA) are likely to be the most useful for testing the stellar-ionized LINER scenario. Both classes of object deserve additional X-ray study on their own merits, apart from the current comparison. For SNe, the X-ray behavior provides a probe of the circumstellar environment and blast-wave physics; for LINERs, the relationship to more luminous AGN at X-ray wavelengths needs further study. A potential pitfall for future X-ray studies of LINERs that push to fainter luminosities is confusion with normal X-ray binaries; M33, for example, has a nuclear X-ray luminosity of ~ 1039 erg s" 1 that might best be explained by a small number of stellar binaries (Hernquist et al. 1991). Several avenues also exist for advancing the understanding of LINERs through theoretical work and new approaches to data analysis. Hubble type morphology provides only an approximate measure of more fundamental physical parameters of galaxies. The role of cloud density in influencing nuclear emission characteristics would be better tested by the extent to which direct measures of central potential well depth - such as pressure (inferred from X-ray data) or stellar velocity dispersion - provide sequencing parameters for LINER versus HII nucleus behavior. Comparable work has already been attempted using stellar metallicity as a fundamental parameter (Bonatto, Bica & Alloin 1989). Additional thought should also be given to the structure and geometry of nuclear nebulae. Ionizing stars located in very dense interstellar environments are known in our own Galaxy, and are identified as ultracompact HII regions. These sources are not observed as LINERs, and in fact, are generally undetected optically because they are so heavily embedded. If stars (or accreting black holes) in dense media generate LINERs, the nebular structure must be modified in such a way that coverage is reduced. Violent gas motions within galaxy nuclei might be involved. Within our own Galactic nucleus, the nebular gas within the central parsec is arranged in an apparent ring that is probably photoionized by stars residing within a central cavity (e.g. Jackson et al. 1993); the geometry is far from being a simple Stromgren sphere. Issues of geometry are also important for progress in understanding the role of shocks in generating compact LINERs. The correlation of critical density with velocity width for nebular features in LINERs is providing important information on nuclear structure, although it remains subject to ambiguity. In the accreting black-hole paradigm for AGN, this trend reflects a trend of higher density at smaller radii for clouds orbiting the central collapsed object. Quantitatively this scenario may be imperfect; the observational trend typically implies n oc v8, while Keplerian motion around a central ionizing source implies t; oc r~° 5 so that n oc r~ 4 . If the ionization is dominated by a central source, we would then expect U oc r + 2 - which is contrary to the usual expectation that nebular ionization is dropping with increasing radius (supported observationally by variability studies of Seyfert 1 nuclei). This apparent defect may have a simple explanation if stars are important for defining the nuclear potential so as to make the gradient in orbital velocity shallower (Filippenko k Sargent 1988). A simple mapping between density and velocity also neglects a possible weighting by ionization state and emissivity as a function of radius. Alternatively, understanding of the empirical trend might benefit from consideration of more general scenarios in which cloud compression is accompanied by acceleration.

Shields: Stellar Ionization of Weak AGN

363

I thank the conference organizers for partial financial support allowing me to attend this very stimulating meeting, and A. Filippenko and R. Terlevich for informative comments. My work on AGN is supported by NASA through grant #HF-1052.01-93A from STScI, which is operated by AURA under NASA contract NAS5-26555.

REFERENCES BARR, P., GIOMMI, P., WAMSTEKER, W.,

GILMOZZI, R., & MUSHOTZKY, R. F. 1985 Bull.

Am. Astron. Soc. 17, 608. BEGELMAN, M. C. & FABIAN, A. C. 1990 Mon. Not. R. Astron. Soc. 244, 26P. BONATTO, C , BlCA, E. & ALLOIN, D. 1989 Astron. Astrophys. 226, 23. BREGMAN, J. N. & PlLDlS, R. A. 1992 Astrophys. 3. Lett. 398, L107. BRUZUAL A., G., PEIMBERT, M. & TORRES-PEIMBERT, S. 1982 Astrophys. J. 260, 495. BURSTEDM, D., BERTOLA, F., BUSON, L. M., FABER, S. M. & LAUER, T. R. 1988 Astrophys. J. 328, 440. CHEVALIER, R. A. 1981 Astrophys. J. 251, 259. CHEVALIER, R. A. 1982 Astrophys. J. 259, 302.

P. et al. 1993 Astron. J. 106, 1371. S. & HARDY, E. 1992 Astron. J. 103, 711. DONAHUE, M. & VOIT, G. M. 1991 Astrophys. J. 381, 361. FABBIANO, G., KIM, D.-W. & TRINCHIERI, G. 1992 Astrophys. J. Suppl. 80, 531. FERGUSON, H. C. & DAVIDSEN, A. F. 1993 Astrophys. J. 408, 92. FERLAND, G. J. & NETZER, H. 1983 Astrophys. J. 264, 105. FESEN, R. A. 1993 Astrophys. J. Lett. 413, L109. FESEN, R. A. & BECKER, R. H. 1990 Astrophys. J. 351, 437. FESEN, R. A. & MATONICK, D. M. 1993 Astrophys. J. 407, 110. FILIPPENKO, A. V. & HALPERN, J. P. 1984 Astrophys. J. 285, 458. FILIPPENKO, A. V. & SARGENT, W. L. W. 1985 Astrophys. J. Suppl. 57, 503. FILIPPENKO, A., V. & SARGENT, W. L. W. 1988 Astrophys. J. 324, 134. FILIPPENKO, A. V. & TERLEVICH, R. 1992 Astrophys. J. Lett. 397, L79.

CRANE,

DELISLE,

FOSBURY, R. A. E., SNIJDERS, M. A. J., BOKSENBERG, A. & PENSTON, M. V. 1981

Mon.

Not. R. Astron. Soc. 197, 235. GRIRICIN, G., BERTOTTI, G., MARDIROSSIAN, F. &; MEZZETTI, M. 1990 Mon. Not. R. Astron. Soc. 247, 444. GRJRICIN,

G., et al. 1993 Astron. Astrophys. 275, 390.

GOODRICH, R. W. & KEEL, W. C. 1986 Astrophys. J. 305, 148. GUSTEN, R. GENZEL, R., WRIGHT, M. C. H., JAFFE, D. T., STUTZKI, J. & HARRIS, A. I.

1987 Astrophys. J. 318, 124. J. P. & STEINER, J. E. 1983 Astrophys. J. Lett. 269, L37. HECKMAN, T. M. 1980a Astron. Astrophys. 87, 142. HECKMAN, T. M. 1980b Astron. Astrophys. 87, 152 (H80b). HECKMAN, T. M. 1989 in Massive Stars in Starbursts (ed. C. Leitherer et al. ), p. 289. Cambridge. HECKMAN, T. M., BAUM, S. A., VAN BREUGEL, W. J. M. & MCCARTHY, P. 1989 Astrophys. J. 338, 48. HECKMAN, T. M., ARMUS, L. & MILEY, G. K. 1990 Astrophys. J. Suppl. 74, 833. HECKMAN, T. M., BALICK, B. & CRANE, P. C. 1980 Astron. Astrophys. Suppl. 40, 295. HERNQUIST, L., HUT, P. & KORMENDY, J. 1991 Nuture 354, 376. HALPERN,

364

Shields: Stellar Ionization of Weak AGN

Ho, L. C , FILIPPENKO, A. V. & SARGENT, W. L. W. 1994 in Active Galactic Nuclei Across the Electromagnetic Spectrum (ed. A. Blecha & T. Courvoisier), in piess. Kluwer. HUMMEL, E., VAN DER HULST, J. M. & KEEL, W. C. 1987 Astton. Astrophys. 172, 32. HUMMEL, E., VAN DER HULST, J. M., KENNICUTT, R. C. & KEEL, W. C. 1990 Astron. Astiophys. 236, 333. JACKSON, J. M. et al. 1993 Astrophys. J. 402, 173. JONES, P. A. & MCADAM, W. B. 1992 Astrophys. J. Suppl. 80, 137. KENNICUTT, R. C , KEEL, W. C. & BLAHA, C. A. 1989 Astron. J. 97, 1022 (KKB). KIRHAKOS, S. & PHILLIPS, M. M. 1989 Pub. Astron. Soc. Pacific 101, 949. KORMENDY, J. & MCCLURE, R. D. 1993 Astron. J. 105, 1793. LAOR, A. &; DRAINE, B. T. 1993 Astrophys. J. 402, 425. LAUER, T. R., et al. 1993 Astron. J. 106, 1436. LEIBUNDGUT, B. et al. 1991 Astrophys. J. 372, 531. LEITHERER, C , GRUENWALD, R. & SCHMUTZ, W. 1992 in Physics of Nearby Galaxies (ed. T. X. Thuan, C. Balkowski & J. T. T. Van), p. 257. Editions Frontieres. MCCALL, M. L., RYBSKI, P. M. & SHIELDS, G. A. 1985 Astrophys. J. Suppl. 57, 1. MCGAUGH, S. S. 1991 Astrophys. J. 380, 140. MuSHOTZKY, R. 1993 in The Nearest Active Galaxies (ed. J. Beckman, L. Colina & H. Netzer). CSIC. PELAT, D., ALLOIN, D. & FOSBURY, R. A. E. 1981 Mon. Not. R. Astron. Soc. 195, 787. PETRE, R., MUSHOTZKY, R. F., SERLEMITSOS, P. J., JAHODA, K. & MARSHALL, F. E.

Astrophys. J. 418, 644. PHILLIPS, M. M., PAGEL, B. E. J., EDMUNDS, M. G. & Soc. 210, 701. PHILLIPS, M. M., 1986 Astron. J. 91, 1062. POGGE, R. W. 1989 Astrophys. J. Suppl. 71, 433.

DIAZ,

1993

A. 1984 Mon. Not. R. Astron.

REICHERT, G. A., PUCHNAREWICZ, E. M., FILIPPENKO, A. V., MASON, K. O., BRANDUARDIRAYMONT, G. & Wu, C.-C. 1993 IN The Nearest Active Galaxies (ED. J. BECKMAN, L. COLINA & H. NETZER). CSIC.

S., et al. 1993 Astrophys. J. 416, 167. G. A. & TINSLEY, B. M. 1976 Astrophys. J. 203, 66. SHIELDS, J. C. 1992 Astrophys. J. Lett. 399, L27. SHU, F. H., ADAMS, F. C. & LIZANO, S. 1987 Ann. Rev. Astron. Astrophys. 25, 23. SHULL, J. M. & MCKEE, C. F. 1979 Astrophys. J. 227, 131. SINGH, K. P., RAO, A. R. & VAHIA, M. N. 1990 Mon. Not. R. Astron. Soc. 246, 706. SPERGEL, D. N. & BLITZ, L. 1992 Mature 357, 665. STORCHI-BERGMANN, T., 1993 Astrophys. J. Lett. 410, L l l . TERLEVICH, R., TENORIO-TAGLE, G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. R. Astron. Soc. 255, 713. RYDER,

SHIELDS,

TERLEVICH, E., DIAZ, A. I. & TERLEVICH, R. 1990 Mon. Not. R. Astron. Soc. 242, TERLEVICH,

R. &

MELNICK,

271.

J. 1985 Mon. Not. R. Astron . Soc. 213, 841.

THOMAS, P. A., FABIAN, A. C , ARNAUD, K. A., FORMAN, W. fc JONES, C. 1986, Mon.

R. Astron. Soc. 22, 655. WEILER, K. W., PANAGIA, N. &

SRAMEK,

Not.

R. A. 1990 Astrophys. J. 364, 611.

WEILER, K. W., SRAMEK, R. A., PANAGIA, N., VAN DER HULST, J. M. & SALVATI, M.

1986

Astrophys. J. 301, 790. WEILER, K. W., VAN DYK, S. D., PRINGLE, J. E. & PANAGIA, N. 1992 Astrophys. J. 399, 672. WROBEL, J. M. & HEESCHEN, D. S. 1991 Astron. J. 101, 148.

Line Profiles in Compact Supernova Remnants and Active Galactic Nuclei 1 2 ByROBERTO CID FERNANDES 2 AND ROBERTO TERLEVICH

'institute of Astronomy, Madingley Road, CB3 OHA Cambridge, U.K. 2

Royal Greenwich Observatory, Madingley Road, CB3 OEZ Cambridge, U.K.

This contribution summarises our studies on the emission line profiles from compact Supernova Remnant shells and how they might be related to the broad line profiles in active galaxies. The emphasis is on the theoretical problems associated with radiative transfer effects in spherical and irregularly shaped shells. Line profiles from systems containing many compact remnants are also calculated with the aim of comparing the results to luminous active nuclei, where several remnants are expected to coexist. The observed diversity of profile characteristics in QSOs and the consequences it has to the starburst model are discussed.

1. Introduction Line profiles of any astrophysical object, from stellar atmospheres to the Broad Line Region of active galaxies, provide valuable information on the physical and dynamical conditions which may be used to constrain or even reject theoretical models for such objects. Our goal in this work is to develop models for the emission line profiles in compact Supernova Remnants (cSNR) and Active Galactic Nuclei (AGN). The presence of both cSNR and AGN in the same title can only mean that we are talking about the starburst model for AGN of Roberto Terlevich and collaborators (see Terlevich et al. 1992 as well as Franco's, Plewa's and R. Terlevich's papers in this volume). Indeed, the main idea here is to see how well the starburst model does regarding the broad lines in active galaxies. In a strict sense, however, the starburst model for the Broad Line Region (BLR) is in fact a model for cSNR; whether it does or does not apply to AGN as well is yet to be seen. In this sense, it is perhaps more important to compare its predictions to observed cSNR before making the AGN connection. The difficulty here, obviously, is that bona fide cSNR are much more rare and less well studied than AGN. The comparison with AGN, on the other hand, is a crucial one, since it may add convincing evidence for or against the starburst model.

2. Models of line profiles in cSNR cSNR are associated with SNe whose distinct properties (light curve and spectrum) are thought to be the result of the interaction between the expanding envelope of the exploded star and its circumstellar medium (CSM). When the CSM is dense, UCSM ^ 107 cm"3, most of the ejecta's kinetic energy is thermalized and radiated away in a few years. We shall adopt the model of Terlevich etal. (1992) as our basic working model for a cSNR. The basic structure of the Terlevich etal. model is sketched in Figure la. It consists of two strong shock fronts, the leading and reverse shocks, processing the CSM and the ejecta respectively. After about 200 days of evolution catastrophic cooling takes place behind the leading shock, resulting in the formation of a thin shell of cold dense gas. A similar process produces another shell behind the reverse shock. Both shells are photoionised by the X-rays produced by the shock fronts and the cavity of hot gas between 365

366

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

(a)

CSM (b) n= le7cm-3 Inner Shell , Leading Shock n = lelOcm-3 ^ ^ v = 5000 km/s \ v = 6000 ta/s > Outer Shell Ar=lel3cflf^/ >T^ Hot Cavity^s^JX \ T=Ie8K Nw JK^ n = lel2 cm-3 R e v e r s e ~^/ f \ \ v = 5000km/s \ ^^, .^^^ \ \Ar=lellcm Shock /yC. v=1500kiWs/

^ ^C

/ ^ Eiecta - \ \

f I e 16 cm

FIGURE

I000< y < 10000 km^s \

\

\

1 11

3el6cm

1. (a) Schematic structure of a cSNR. (b) p-z co-ordinate system used in the calculation of line profiles.

them. The emission line spectrum of these shells have characteristic line ratios similar to the ones in AGN. Moreover, their widths (a few thousand kms" 1 ) are also comparable to AGN line widths. The unshocked ejecta and CSM are also photoionised by the remnant's X-rays. Only the former, however, produces broad lines, since the progenitor's wind velocity is small. SN 1987F, the so called "Seyfert 1 impostor" (Filippenko 1989), is possibly a real-life counterpart to Figure la. SN 1988Z (Stathakis & Sadler 1991; Turatto etal. 1993) and SN 1979C (Fransson etal. 1982) are other examples of cSNR. 2.1. Line profiles in spherical shells

An important characteristic of this simple model is that all the cSNR line-emitting regions have the same basic geometry: they are all spherical shells. Even if a particular line is only produced within one such region, due to the spherical symmetry this "sub-region" also has a shell geometry. The physical and dynamical conditions in each shell are however quite different from one another. The basic problem we first have to tackle is therefore calculating the emission line profile from an isolated expanding spherical shell under general physical conditions. This problem is depicted in Figure lb, which introduces the (p, z) co-ordinate system to be used throughout this paper. Besides the geometry, the calculation of line profiles requires quantitative knowledge of the velocity field, the line emissivity and the nature and distribution of the sources of opacity within and external to the line-emitting region. The velocityfield:The ejecta has a v(r) oc r velocity field, and the thin shells move at roughly constant velocities. The internal velocity distribution in the shells is, to a first approximation, thermal at T « 104 K, though turbulent motions may also occur due to their hydrodynamical nature. The line emissivity: Each portion of a shell has a line production rate regulated by the local ionisation state, which is determined by the density, chemical composition and ionising radiation reaching that particular point. The detailed modelling of the emission line spectrum of each zone has to be carried out using a photoionisation code. As a first approach, however, it is more useful to treat the volume emissivity in a parametrized form as a power-law: JL{T) °C T^. The distribution of scatterers and absorbers: The regions exterior to the shell are not empty. The hot cavity between the two thin shells, for instance, contains material which is fully ionised, and may scatter line photons produced elsewhere in the remnant.

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

367

Furthermore, the shells themselves are not necessarily optically thin, which means that line photons produced at a given location may interact with other neighbouring atoms and possibly be destroyed by a collisional process before escaping the region. Radiative transfer effects may therefore be important. In generic terms, the problem may be posed as follows. Consider a spherical shell between ri and ro, with a velocity field v(r) = vo(r/ro)a, emissivity law JLC*") = jo(r/roy and optical depth T£ to line photons from a given transition. (Disregard absorption/scattering external to the shell for the moment.) What is the line profile resulting from such a system? The classical method of solving this problem involves solving the transfer equation along rays of constant impact parameter p to get the intensity of line radiation emerging from the side of the shell nearer to the observer and directed towards him/her (dashed line in Figure lb). This is then integrated in circular rings of area 2-wpdp to obtain the total shell flux at a particular frequency/velocity (e.g. Mihalas 1978). Proper account must be taken of the Doppler shifts resulting from the shell's velocity field. In what follows we summarise the results for some limiting cases of interest here, concentrating on the a = 0 case, applicable to the thin shells associated with the leading and reverse shocks in a cSNR. We define L(x) as the line profile, with x = vz/vo as the adimensional velocity scale, running from —1 (red end) to + 1 (blue end), v2 = v(r) cos 9 being the velocity component in the observer's direction at a given location (r, 6) in the shell. 2.1.1. The optically-thin case

When TL •< 1 , L(x)dx is simply the integral of ji(r)dV over the volume of the shell whose velocity component in the observer's direction lies in the vz = VQX to VQ(X + dx) range. This yields a well know solution (e.g. Robinson 1994):

{

constant

if a = 0

l

The line profile is flat-topped from x = —1 to x — 1 for shells of constant velocity (a = 0), independently of the extent of the shell or of the emissivity law. This is so because the iso-i^ lines in the a = 0 case are lines of constant angle 6, such that cosfl = \i = x, and the volume of the region emitting between x and x + dx is simply proportional to d/z, being therefore independent of a;. For an ejecta-type velocity law (a = 1), the profile has also got a flat-topped core between x — —rj/ro and x = r^/ro, dropping as x@+2 in the line wings. In this case the iso-i^ lines are defined by z — rox, and the region emitting in x —+ x + dx are shell slices of thickness dz = rodx, whose volume is the same anywhere in the — r,- < z < r{ range. Out in the "polar caps" (r,- < \z\ < ro), i.e. in the profile wings, this volume decreases, and so does the line profile. 2.1.2. The optically-thick case We found that when T& ^> 1 the calculation of L(x) depends on the geometrical thickness of the shell. The standard way to find the line profile for an optically-thick expanding shell is the same used for the atmospheres of WR stars (e.g. Castor 1970). It makes use of the Sobolev approximation, which basically replaces the line absorption profile (Av)—the probability of a line photon displaced At/ from the transition's central frequency UQ being absorbed—by a Dirac delta function. This amounts to assuming that the velocity gradient along a given ray p in the shell is much larger than the ID thermal velocity of atoms in the shell, vT = (kT/m)1'2 (i.e. w 10(T/10 4 ) k m s " 1 in the case of

368

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

hydrogen atoms), so that any absorption/scattering takes place locally. However, this can only be true if the shell thickness is large enough so that the values of vz at its inner and outer edges differ by much more than VT- This sets the limit of applicability of the Sobolev approximation. Mathematically, we can express it as Ar/r ^> {VT/VQ)2 « (10/5000)2 = 4 x 10~6. Though this is probably the regime in WR stars, where the line emitting regions are very extended, the outer shell in Figure la, having Ar/r « 3 x 10~6, is too thin to be modelled with the Sobolev method, and needs to be treated with a different approach. The inner shell has Ar/r « 5 x 10~4 and may therefore have its profile calculated using this method. Geometrically thick shells A very general result of the Sobolev method for expanding shells of constant velocity is that the line profile is parabolic, independently of ji,(r) and Ar (e.g. Mihalas 1978). The reason for this is that, since the iso-i^ lines are given by [M — x, photons produced at the top part of the shell (8 = 90°, fi = 0 = x) meet a high velocity gradient while travelling through the shell towards the observer. Since the Doppler shift is much larger than the width of the absorption profile (« VT), absorption of such photons is unlikely to occur as they move away from their birthplace. At the profile extremes (x = ±1 = fi, 9 = 0 or 180°), on the other hand, the velocity gradient is nil in the z-axis. Photons produced inside the shell are thus scattered many times before escaping the shell, increasing their chances of being destroyed and/or escaping in other directions and producing the depression of the line profile as it approaches its wings. Although the Terlevich et al. (1992) model predicts an outer shell too thin to be modelled with the Sobolev method, it may well be the case that the shell is not actually that spherical or that thin. If, for instance, hydrodynamical instabilities distort the shell's surface into an irregularly-shaped contour with a radial extent larger than the shell's original thickness, then we could perhaps model this situation as an effectively thicker shell, whose average density would be smaller than the original one in order to accommodate the same mass. This "reshaping" of the shell is actually seen in the hydro models in Terlevich etal. (1992). The parabolic profile solution may therefore be applicable to cSNR outer shells after all. Geometrically thin shells When Ar/r < (VT/VQ)2 we have little alternative other than to try to solve the full transfer problem by brute force (i.e. numerically). We did this adopting a power-law emissivity function and assuming complete redistribution (i.e. that the absorption and emission profiles are equal). The transfer equation for each frequency x was then solved for many rays with different impact parameters, p, and integrated in 2vpdp rings as explained above. The results were quite puzzling initially. Most profiles had an M-shape, i.e. two maxima at x = ±1 and asteady decline to zero towards x = 0. Runs with Ar/r > (VT/V0)2 did recover the parabolic profiles predicted in the Sobolev approximation, so the code probably worked OK. What then is the reason for this double-peakedness? The interpretation we came up with is based on purely geometrical effects. Unlike the optically-thin case, when T£ ^> 1 not all the volume emitting at a particular x is seen by the observer, but only the very outer skin of the shell where the line optical depth is of the order of unity. Also, unlike geometrically-thick shells, the velocity gradient is not large enough to allow photons from deep within the shell to escape. We may therefore regard the shell as a radiating spherical surface, much like a ping-pong ball. The region in this surface producing photons within x —+ x + dx consists of a ring between fi = x

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

369

and n + dfi, whose projected area as seen from the observer is dAp(x) = 2irpdp, or, since p = rosm6, dAp(x) = 2irrl\(t\d(j. = 27rro|a;|rfa;. The important point to notice here is the |^| term in dAp(x). Basically, this term says that the ring's projected area decreases as n approaches 0. This is indeed in agreement with anyone's ball playing experience: One does not see the top of the ball (9 = 90°) looking at it from 5 = 0°! This rather elementary geometrical fact is actually the reason for the M-shaped ping-pong ball's line profile. If the surface emits as I(fi) then L(x) becomes:

L(x)dx = I(ji = x)dAp{x) = 27r2r^/(x)|z|dz. Therefore, if I(n) = constant, i.e. if the shell's surface emits isotropically, we arrive at L(x) oc \x\, as found in the numerical calculations. It is important to realise which of our assumptions brings about this double-peaked structure in the profile. Clearly, the condition that we see only the shell's surface is what introduces the \fi\ factor in the projected area and consequently the central minimum in the profile. This assumption does break down when the shell's optical depth is small or when Ar increases, but it should hold, at least qualitatively, in the TL > 1, A r / r < (VT/V0)2 regime. The assumption that I(n) is constant depends on the detailed run of JL(I~) within the shell, but it does apply when jx(r) does not vary a lot within a few optical depths inside the shell. When that is not the case, a plane parallel solution for the transfer equation can be used for J(/i) and plugged into the solution for L(x). The assumption of perfect spherical symmetry is also crucial. One has always to be suspicious about such Pythagorisms. The hydro simulations do show deformations of the shell due to instabilities, casting doubt upon the sphericity hypothesis. Could slight distortions of the shell bear significant changes to its line profile with respect to the spherical case? This is what we study next. 2.2. Line profiles in irregular shells All the fast-moving line emitting shells in the cSNR model of Terlevich etal. (1992) are immersed in, and interacting with, a physically complex environment, with strong density, temperature and pressure gradients. Under such hostile conditions, the shells are subjected to all sorts of hydrodynamical instabilities, which may distort its shape and produce significant deviations from the spherical case discussed in the previous sections. As we have just seen, in the limit of optically thick but geometrically-thin shells the hypothesis of sphericity leaves strong imprints in the shell's line profile. It is therefore important to establish what are the effects of such irregularities upon the line profiles of cSNR. Basic physical intuition tells us that small perturbations to the spherical model should not imply big departures from its basic results. As we are about to see, however, that may not be so in this particular problem. As a first approximation to this problem we assume that the irregularities simply redefine the shell contour, preserving its cohesion—much like a drum's surface is perturbed, but not punctured, when hit by the drummer. This probably corresponds to the second stage of development in a real cSNR shell. Initially, an overall spherical shell is formed. Instabilities then promote the "corrugation" of its surface. These "corrugations" or "fingers" possibly evolve to disrupt the shell into a large number of independent cloudlets. The exact shape and location of the corrugations is of course an absolute unknown! We can only hope to study such irregularities by somehow parametrizing our ignorance. As long as there are many of them equally distributed around the shell's contour, their exact locations are unimportant for the computation of line profiles. Their shapes, nevertheless, do affect it, since they determine the projected area of a given part of the shell. For simplicity, let us assume that the shell distortions evolve into "rectangular fingers"

370

Cid Fernandes k Terlevich: Line Profiles in cSNR and AGN

with a radial extent Arj — pjf and angular size Asj = o-jr, where r denotes the average radius of the shell. We confine ourselves to the 2D problem by assuming that the fingers are symmetric around their radial axis, i.e. that their dimension along the axis perpendicular to the p-z plane is A s / as well. Shells where pj > (Tj look like "sea-hedgehogs", whilst those with pj < aj look more like "amoebas" (Figure 2d). A finger located at angle 9 emits at frequency x = \i — cos 9. Its contribution to the line profile is proportional to its projected area as seen by the observer, which is the sum of two components: the "radial wall", with length Arj and the "lateral wall", Asj long. The finger's projected area is therefore Ap(6) = AsjArjsmO + AsjAsfCos9 - rAs/ (pj\/l - x2 + aj\x\\ , where we have replaced \L by x. Since L(x) oc Ap(8), the line profile is the sum of a circular profile due to the radial walls and an M-shaped profile due to the lateral walls. The interpretation is quite simple. The lateral-wall component is just a spherical shell profile, as if no radial perturbations existed. The circular profile of the radial-wall component results from the fact that, due to projection effects, these walls are seen by the observer to their full extent as 9 approaches 90°, whereas they are practically not seen when located at low angles. Thus, a radial instability producing a Rayleigh-Taylor-like finger at 9 = 90° exposes a part of the shell surface which was not visible to the observer in the strictly spherical case. This "new perspective" of the shell's surface completely changes its line profile when pj ^> cry. Radial perturbations of a few percent are enough to turn the spherical shell profile practically "upside-down", as seen in Figure 2, where the numerical results are presented. These results were corroborated numerically for other finger shapes as well. Irregularities of the kind discussed here should bring little change to the line profile in the optically-thin regime, as well as in the optically-thick, geometrically-thick case. 2.3. Which model, then? Having discussed so many possibilities for the line profiles of cSNR shells, we must now discuss which model best describes cSNR. Figure 2 summarises the models discussed above. The optically-thin alternative can be discarded on the grounds that photoionisation modelling predicts the shells to be optically thick to the main transitions, just like in any standard broad line region model (e.g. Ferland et al. 1992). Observationally, few AGN show flat-topped profiles, which also indicates that the shells are optically thick. Geometrically-thin, optically-thick shells, if they exist, are likely to be corrugated enough to produce roundish instead of M-shaped lines. A handful of AGN are known to exhibit double-peaked profiles, which are usually modelled with relativistic accretion disks (e.g. Chen etal. 1989). Although thin, optically thick spherical cSNR shells can produce this sort of profile, there are other configurations involving external obscuration of the shell which yield this same result, so the situation is not clear in this case. Most AGN do not have double-peaked profiles anyway. Parabolic profiles do seem to be observed in cSNR (e.g. 1987F and 1988Z, Chugai 1991; Chugai & Danziger 1994), which points to the geometrically- and optically-thick shell model. As far as AGN are concerned, although some objects could have their broad line profiles reasonably well fitted with a parabola, many definitely could not, as can be seen in a quick browse through Stirpe's (1990) atlas of profiles. Part of the reason for this discrepancy has to do with the fact that the line profile in AGN usually contains the contribution of not one but several cSNR, as discussed in the next section. Even in a single cSNR, the total broad line profile is a mixture of profiles from the leading shell, the reverse shell and the ejecta. The exact proportions of

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

371

p, = 10- 2 at = 1O"1

I ' I I

I I I l| I I I I | I I I I I

FIGURE 2. Emission line profiles for shells of constant velocity and different physical conditions: (a) optically-thin shell, (b) optically-thick, geometrically-thick shell, (c) optically-thick, geometrically-thin shell. The bottom panel (d) shows the effects of slight distortions of the shell surface on its line profile in the optically-thick, geometrically-thin regime. The top figures in (d) show the shell's contour, whereas the bottom ones show the corresponding line profile.

these components evolves as the remnant ages (though one of them probably dominates the line luminosity). Another factor not taken into account in the present analysis are the effects of continuum opacity. The hot cavity, for instance, can have a Thomson optical depth re, of a few tenths. Photons from the far side of the leading shell can therefore be electron-scattered towards other directions, producing a blue asymmetry. As res increases, line broadening could take place due to the high temperature of the

372

Cid Fernandes k Terlevich: Line Profiles in cSNR and AGN

electrons. The inner parts of the ejecta are another source of obscuration. Clearly, there is plenty of scope for improvement of the models summarised here; we are just scratching the surface of a very complex problem.

3. Multi-cSNR systems: QSOs When applying the cSNR model to AGN line profiles we must not forget that they are not single cSNR systems. A massive young cluster, such as the one predicted to be in the core of NGC 5548, for instance, has a SN rate of 1 every 2 years or so (Aretxaga k Terlevich 1993). Since a cSNR lives for TSN « 5-10 years, chances are that we seldom observe an isolated cSNR in this Seyfert, and even more so as we move up in the luminosity scale towards QSOs—the SN rate in an MB — —26 QSO is VSN « 40 yr" 1 . To extend the cSNR line profile calculations to AGN we must somehow allow for their multi-cSNR nature. This can be done adding up many cSNR profiles, each of them with a different age (and perhaps different characteristics). Two things are required for this: (1) the line profile of an individual cSNR of given age and initial conditions (CSM density, kinetic energy, ...) and (2) a prescription for the evolution of cSNR. As to the first point, we consider only the remnant's outer-shell profile (the most luminous in lines as Ha and Hp) and adopt either a flat-topped or a parabolic profile. Regarding point (2), the evolution of a cSNR, and in particular of its leading shock and shell, is quite complex in detail, involving several radiative bursts and shock reflections (Plewa, this volume; Terlevich et al. 1994). It was found, however, that the analytical formulae for the shock evolution derived by Shull (1980) and Wheeler et al. (1980) provide a reasonable approximation to the numerical results (Terlevich et al. 1992). We therefore adopt them as a recipe for the cSNR evolution, with the understanding that we might be missing some detailed features. The analytical solutions of Shull (1980) and Wheeler etal. (1980) provide expressions for the leading shock luminosity (Ls) and velocity (vs) as a function of time: Ls(t) = L,(t,g)(t/tsg)-n'7 and v.{t) = vs(tsg)(t/tsg)-5/7, with v,(tsg) = 4600es{8n7/4 kms" 1 . The time scale t,g corresponds to the start of the radiative phase: tsg = 230eg{ n? ' 4 days, where £51 is the explosion energy in units of 1051 erg and n? is the CSM density in units of 107 cm' 3 . The shell velocity is always very similar to vs, so we already have a prescription for the evolution of the profile width. We have, however, no analytical expression for the evolution of the line luminosity, LL(<), although that could be derived from photoionisation modelling as in Terlevich etal. (1992). We find it more useful at this exploratory stage to adopt a parametrized form, where Li is proportional to the shock luminosity: £L(£) = La(t)eL(t), where ££,(<) plays the role of a "line efficiency" or "bolometric correction" factor for line L. The evolution of the line profile for a multi-cSNR system can now be derived by Monte Carlo simulations. If at time t we have NSN CSNR alive and emitting at line L, the total line profile is simply

-«O=

E

where the Vs are the explosion dates and the line profile LL{V,1 — i,) was split into a normalisation term (actually the weight given to cSNR of a given age in the computation of L%1%) and a dimensionless function f(v,t-U) containing the profile shape—/ depends on the remnant's age indirectly through the shell velocity vs(t — <,). It is clear that as NSN increases the profile V^ will show less and less net variability

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

373

as the relative contribution of newly-born cSNR becomes small compared to that of older remnants. This "steady-state" profile can be derived analytically by transforming the above summation into an integral. The resulting expression is simply V>SN times the integral of Li,(v,t) from t = tsg to tend, the initial and ending times of the shell's line emitting phase. It is more interesting to write this integral in terms of the shell velocity using the relation between vs and t. After some algebra we obtain

where overlined quantities are evaluated at t = tsg. The terms inside the integral sign express the contribution of shells with velocity in the vs —» vs + dvs interval to the total profile at velocity v. The line profile can now be obtained by specifying ei,(<). Figure 3 illustrates the total Ha and Hp profiles obtained using the line efficiencies derived from photoionisation models for times between 1 and 25 t3g's—the code CLOUDY was used for this purpose (Ferland 1990). The conditions are identical to the ones in the cSNR model in Terlevich et al. (1992). Bearing in mind the simplicity of the model, the overall similarity to AGN profiles is quite good. 3.1. The Hp/Ha profile ratio In Terlevich etal. (1992) it was remarked that the evolution of the Balmer lines is such that the ratio Hp/Ha decreases with time. Since the shell velocity also decreases with time, clusters containing several cSNR should exhibit a Hp/Ha profile ratio increasing towards the line wings, where the emission is dominated by young remnants. This is the tendency observed in AGN (e.g. Crenshaw 1986). Figure 3 shows the computed efficiency factors and line profiles, as well as their ratio for both the flat-topped (solid line) and parabolic (dotted line) versions of f(v,vs). The prediction is confirmed: the Hp/' Ha profile ratio increases from « 0.1 in the line core to sa 0.3 in the line wings. Again, the agreement with observations is quite remarkable, especially considering that this is the simplest model possible.

3.2. Statistical properties The profiles above are actually time averaged profiles. A single observation of an AGN may differ from this statistically expected value due to the stochastic nature of SN events. To study what effects this stochasticity might have we performed 1000 simulations of the line profile for each choice of fi(<). a n d f° r 10 values of VSN between 0.1 and 100 yr" 1 , corresponding to absolute magnitudes in the -19 > MB > —27 range approximately. For each individual profile we determined its half width at zero intensity (HWZI), full width at half maximum (FWHM) and inter-percentile velocities (IPV) at 10, 20 and 50%, as well as the ensemble average, median and standard deviation of these quantities. The results are presented in Figure 4. The simulations used a parabolic line profile and the Hp line efficiency shown in Figure 3. The first thing which strikes the eye in Figure 4 is the apparent correlations between the profile width measures and I^SN- Besides that, it is also clear that the spread in profile parameters is smaller the more luminous the AGN is. These two effects are direct consequences of the stochastic nature of multi-cSNR systems. The more cSNR there are, the higher are the chances of picking a very young, fast one, which explains why HWZI and IPV(10%) "correlate" so well with i/SN. The profile's FWHM should be regarded with caution, since it depends more strongly on the assumptions for CL(^)- The smaller dispersions at high luminosity are also a consequence of the larger number of remnants. The optimistic reader will regard these results as good news. Observationally, it is

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN

, I ,, , ,_

:'

—

3--2

_

^

1

.'A'. '' ' / \ '• \' : I \\

file

374

o

u

1.5

cm

V

5 -3

c

.' /

— - ^

- 4 Li , , . , i . , , , i . . . . i . , , 0 5 10 15 20 ,

,

,

,

,

,

•

U

25

-4000-2000

t/t. g

V

, ,| .. ,| ,, ,| ,,, |, 2 0.3

. . I i .

-

I

j

i

i

|

,

,

,

|

.

.

-

UJ

-'

(S

K

ofi

J20.2 -

\

/ • \

a, 0.1 -

-

/ • '

\

:

\J, /

'

, . 1 , . . 1 , , , 1 , , , 1 ., , I , ,

-4000-2000 0 2000 4000 v (km/s)

-

£

^0.5

-

a 0

-4000-2000 V

0 2000 4000 (km/s)

•A1 // VV \\\\. • i , i i r > >^

•

-

•

- 1 - , , ,

0 2000 4000 (km/s)

FIGURE 3. Evolution of ez, for Ha and Hp and the resulting line profiles and profile ratio. Solid and dotted lines correspond to parabolic and flat-topped individual cSNR profiles respectively.

known that the line widths do correlate with the nuclear luminosity (Shuder 1984; Blumenthal etal. 1982; Joly etal. 1985; Padovani k Rafanelli 1988; Stirpe 1991). The question of whether the dispersion of profile properties is smaller in QSOs than in Seyferts is essentially unanswered, though looking at the plots in the above references one does get the impression that this might be true. The pessimistic reader, however, looking at these same plots, will realise that the observed scatter at high luminosity is larger than the one in Figure 4. The predicted one-sigma boxes at high VSN seem to be too small. Why is that? The use of the self-similar solutions for the cSNR evolution certainly contribute to this underestimation of the profile properties, since they smooth out the detailed features seen in the hydro models. Apart from that, the models in Figure 4 implicitly adopted two quite radical assumptions: first, that all cSNR are identical within an AGN, and second, that cSNR are identical from AGN to AGN (i.e. they are galaxy independent). If, for instance, the CSM density is not exactly the same for all cSNR in a given AGN, then their time and velocity scales t3g and vj would automatically differ and the dispersion in the profile width parameters would correspondingly increase. The same would occur if the cSNR properties varied among the AGN population. Elsewhere in this volume, Gary Ferland tells us that high-redshift QSOs may have a metallicity several times larger than that of low luminosity AGN (see also Hamann & Ferland 1993 and Aretxaga et al. 1993). That would certainly produce a systematic difference in the evolution of QSOs' cSNR, in the sense that higher Z increases the efficiency of radiative cooling, accelerating the shell formation phase and ultimately leading to higher shell velocities and broader line profiles. The fact that the simplest model faces the difficulties mentioned above points to

Cid Fernandes & Terlevich: Line Profiles in cSNR and AGN FWHM (km/s)

HWZI (km/s)

IPV 10% (km/s)

375

Average line Profile

sooo • ;

-4000-2000 1

0

2000 4000

vel (km/s)

log(SNR) (yr- )

4. Statistical properties of the profile's FWHM, HWZI and IPV(10%) plotted against fSN- Dots represent individual simulations, boxes mark the ±1
a solution along these lines, i.e. allowing for the possibility that not all cSNR are identical to each other, nor are the cSNR occurring in environments of different properties. These are very reasonable possibilities, and preliminary investigations indicate that they could indeed explain the dispersion of profile parameters at the high luminosity end. We emphasize that besides the hypothesis mentioned above, the total line profile in a multi-cSNR system depends strongly on £/,(*). In fact, CL functions not very different from the ones in Figure 3 can result in logarithmic or power-law line profiles, forms which are often used to describe observed AGN profiles (e.g. Penston etal. 1990) 4. Conclusions The calculation of line profiles for cSNR shells proved to be quite complex and sensitive to the model assumptions. Slight deviations from perfect sphericity may imply significant changes to the resulting profile when the shell thickness is very small but its optical depth is large. Optically-thick shells which are wide enough for the Sobolev approximation to be valid produce a parabolic line profile, which seems to be in rough agreement with observations of cSNR, although there are too few cases studied to reach any definitive conclusions. We are clearly in need of some observational feedback here, and we strongly encourage observers to gather more data on cSNR which could help us understand the nature of their line profiles and check the validity of the starburst-AGN connection. The calculations of line profiles from multi-cSNR systems yielded some quite interesting results, despite the simplicity of the model developed here. Not only do we seem to be able to reproduce line profiles in AGN, but also the observed correlations between profile

376

Cid Fernandes k Terlevich: Line Profiles in cSNR and AGN

width and nuclear luminosity. The tendency of the Hp/Ha profile ratio to increase towards high velocities also finds a natural explanation in the multiple-cSNR model. Despite these successes, the rather wide observed range of profile characteristics in AGN seems to imply that our model is oversimplified, and that we need to consider the effects of SNe of somewhat different initial conditions occurring within an AGN as well as systematic differences between the SNe properties in low and high luminosity systems. We are indebted to Gary Ferland, Enrique Perez and Guillermo Tenorio-Tagle for fruitful discussions. The hospitality of the IAC, where part of this work has been carried out, is also duly acknowledged. RCF acknowledges the Brazilian agency CAPES for financial support through grant 417/90-8 REFERENCES ARETXAGA, I. k. TERLEVICH, R. 1993 Astiophys. Space Sci. 205, 69. ARETXAGA, I., CID FERNANDES, R. & TERLEVICH, R. 1993 "First Light in the Universe: Stars or QSOs?" (ed. Rocca-Volmerange etal. ), p. 331. Editions Frontieres. BLUMENTHAL, G. R., KEEL, W. C. &; MILLER, J. S. 1982 Astrophys. J. 257, 499. CASTOR, J. I. 1970 Mon. Not. R. astion. Soc. 149, 111. CHEN, K., HALPERN, J. P. & FILIPPENKO, A. V. 1989 Astrophys. J. 339, 742. CHUGAI, N. N. 1991 Mon. Not. R. astion. Soc. 250, 513. CHUGAI, N. & DANTZINGER, I. 1994 Mon. Not. R. astion. Soc. In press. CRENSHAW, D. M. 1986 Astiophys. J. Suppl. Set. 62, 821. FERLAND, G. 1990 OSU Astronomy Dep. Internal Report (90-02). FERLAND, G. J. , PETERSON, B. M. , HORNE, K.,

WELSH, W. F. &; NAHAR, S. N.

1992

Astiophys. J. 387, 95. FILIPPENKO, A. V. 1989 Astion. J. 97, 726. FRANSSON, C , BENVENUTI, P., GORDON, C , HEMPE, K., PALUMBO, G. G. C , PANAGIA, N., REIMERS, D. & WAMSTEKER, W. 1984 Astion. Astiophys. 132, 1.

F. & FERLAND, G. 1993 Astrophys. J. 418, 11. M., COLLIN-SOUFFRIN, S., MASNOU, J. L. & NOTTALE, L. 1985 Astion. Astiophys. 152, 282. MlHALAS, D. 1978 Stelleu Atmospheies, Freeman. PADOVANI, P. & RAFANELLI, P. 1988 Astion. Astiophys. 205, 53. PENSTON, M. V., CROFT, S., BASU, D. & FULLER, N. 1990 Mon. Not. R. astron. Soc. 244, 357. ROBINSON, A. 1994 Mon. Not. R. astion. Soc. Submitted. SHUDER, J. M. 1984 Astiophys. J. 280, 491. SHULL, J. 1980 Astiophys. J. 237, 769. STATHAKIS, R. A. & SADLER, E. M. 1991 Mon. Not. R. astion. Soc. 250, 786. STIRPE, G. M. 1990 Astion. Astiophys. Suppl. Sei. 85, 1049. STIRPE, G. M. 1991 Astion. Astiophys. 247, 3. TERLEVICH, R., TENORIO-TAGLE G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. R. astion. Soc. 255, 713. HAMANN, JOLY,

TERLEVICH, R., TENORIO-TAGLE G., FRANCO, J., ROZYCZKA, M. & MELNICK, J. 1994

Mon.

Not. R. astion. Soc. Submitted. TURATTO, M.,

CAPPELLARO, E.,

DANZIGER, I. J.,

BENETTI, S., GOUIFFES, C. & DELLA

VALLE M. 1993 Mon. Not. R. astron. Soc. 262, 128. WHEELER,

J. C ,

MAZUREK,

T. J. &

SIVARAMAKRISHNAN,

A. 1980 Astiophys. J. 237, 781.

Composite Galactic Nuclei By B. BOER 1 2

Laboratory for Space Research, P. 0 . Box 9504, 2300 RA Leiden, The Netherlands Astronomisches Institut, Ruhr-Universitat, Postfach 102148, 44780 Bochum, Germany

We looked at galaxies with line ratios, which put them close to the borderline between starburst galaxies and LINERs/Seyferts in diagnostic diagrams. Comparison of the observed line ratios with line ratios from various models indicates the presence of a composite ionising mechanism: star formation in combination with either an active nucleus or with shocks. This is in agreement with the increasing number of observations of active nuclei with circumnuclear star formation and of starburst nuclei surrounded by shock-ionised shells. The far-infrared properties confirm the important role played by star formation in these galaxies.

1. Introduction We searched the literature for galaxies, which, on the basis of their optical line ratios, could not be uniquely classified as either starburst galaxies or Seyferts/LINERs. The (arbitrary) selection criterion used was that the [NII]/Ha ratio should be within 0.3 from the empirical borderline between starbursts and Seyferts/LINERs in the log([NII] A6583/Ha) vs. log([OIII] A5007/H/?) diagram (Veilleux & Osterbrock 1987). In all cases this also meant that the observed log([SII] A6716+31/Ha) and log([OI] A6300/Ha) ratios were close to their corresponding borderlines. In this way a total of 62 galaxies was selected. All selected galaxies are spiral galaxies, mostly late type: 54% are Sb and 27% Sc galaxies. The percentage of barred galaxies is 39%. We used the photoionisation code CLOUDY to see if the observed line ratios can be explained by a single ionising mechanism, i.e. by a starburst or by an active nucleus. Although the models could explain the observed [NII]/Ha ratios, the observed [SII]/Ha and [0I]/Ha ratios are too high to be explained by a "pure" starburst and too low to be explained by a "pure" active nucleus or by shock models. The observed intermediate line ratios are best explained by a "composite" ionising mechanism: a star formation component with weak low-ionisation lines in combination with an active nucleus and/or shocks, responsible for strong low-ionisation lines. The average of the IRAS FIR-luminosities is log[L(FIR)]=10.01(±0.39), the average of the blue luminosities log[L(i?)]=10.08(±0.25), and of the far-infrared to blue ratios log[I,(FIR)/log(I(5)] = -0.08(±0.34). These are typical values for starburst and Seyfert galaxies. The L(FIK)/L(B) ratio is not correlated with L(B), whereas the correlation with I/(FIR) is practically linear. This indicates, that the level of star formation in the sample is variable, ranging from normal (log[L(FIR)/L(5)] = —0.4) to violent (log[L(FIR)/L(£)]~3). The average values of the far-infrared spectral indices are: a(100,60)=—0.68(±0.75), a(60,25)=-2.21(±0.34) and a(25,12)=-1.74(±0.40). The narrow distribution of a(60,25) and the wide range of a(100,60)-values (from —3 to 1), suggests, that the varying far-infrared component emits predominantly at wavelengths ~60 /tm, corresponding to dust temperatures To ~ 60 K. If the 60 /xm-fiux increases, then the 25-pmfluxincreases by roughly the same factor, while the increases at 100-pm and 12-pm are less. This explains why, with increasing L(FIR)/L(B), a(60,25) remains approximately constant, a(100,60) becomes flatter and a(25,12) becomes steeper. It further explains why 377

378

Boer: Composite Galactic Nuclei

a(100,60) and a(60,25) are uncorrelated, and why a(100,60) and a(25,12) are anticorrelated. The dust temperature of ~60 K confirms, that the variable FIR-component in the galaxies of the sample is due to star formation: an active nucleus would heat the dust to higher temperatures (To ~ 200 K), resulting in FIR-emission at shorter wavelengths (A ~ 20/zm) than is compatible with the observations. There is a tight correlation between the FIR and radio continuum fluxes from star forming regions, which holds for all levels of star formation. Condon & Broderick (1988) found an average value for q=\og[L(FIR)/L(B)] of 2.75(±0.14) for star forming regions, whereas in active nuclei the ratio is less than 2.25. The FIR and radio fluxes of the selected sample follow the indicated correlation, confirming that the FIR-emission in the selected galaxies is dominated by star formation. The average value is 2.57(±0.16). 2. Discussion The number of emission-line galaxies for which spectra with high spatial resolution demonstrate the presence of either a nuclear starburst surrounded by a shock-ionised shell, or an active nucleus with circumnuclear star formation increases steadily (e.g. Heckman et al. 1990; Boer 1993). Furthermore, it seems likely, that all Seyferts contain some level of star formation, with the relative importance of the two components varying from case to case, while evolutionary models of starbursts predict the formation of a starburst wind-driven shock-ionised shell. The observed optical line ratios of both these groups should be in the region of the CGNs in diagnostic diagrams. An important question thereby is: since all galaxies are expected to form a shockionised shell, why have so few galaxies with shells been observed (their number is still only a fraction of the number of bona fide starbursts)? One reason is probably the observational difficulty: spectra with high spatial resolution are needed in order to observe the shell directly, e.g. in the form of a spatial variation of the optical line ratios. With a typical shell radius of ~300 pc, a spatial resolution of better than 100 pc per pixel is necessary to observe the shell directly. With a resolution of 1" per pixel, the presence of the shell could only be seen directly in galaxies closer than 20 Mpc. In more distant galaxies, the shell could show up as blue wings in the emission line profiles. Armus et al. (1989) found, that the line profiles of one third of a representative sample of IR-galaxies show broad blue wings in addition to a narrow core. And Mirabel & Sanders (1988) showed, that the redshifts of FIR-luminous galaxies are systematically blueshifted by ~100 km s" 1 compared to their HI 21-cm velocities. Another reason may be that the shell only contributes significantly to the line emission for a short interval of the starburst duration. The shell is formed after ~10 6 yr and breaks open at the top after ~10 7 yr. The gas can then expand freely, reducing the efficiency of shock ionisation (McCray & Kafatos 1987). The optical line emission will then again be dominated by the starburst. REFERENCES ARMUS, L. et al. 1989 Ap. J. 347,

727.

BoER, B. 1993 In Proceedings of the Kentucky Conference on Mass-transfer Induced Activity in Galaxies (ed. I. Shlosman). CONDON, J. J. & BRODERICK J. J. 1988 A. J. 96, 30. HECKMAN,

T. M. et al. 1990 Ap. J. Suppl. 74, 833.

MCCRAY, R. & KAFATOS, R. 1987 Ap. J. 317,

190.

MIRABEL, I. F. & SANDERS, D. B. 1988 Ap. J. 335, VEILLEUX,

S. &

OSTERBROCK,

104.

D. E. 1987 Ap. J. Suppl. 63, 295.

The Nature and Origin of X-Ray Emission in Active Galaxies By HAGAI NETZER School of Physics and Astronomy and The Wise Observatory, The Beverly and Raymond Sakler Faculty of Exact Sciences, Tel Aviv University Recent observations of X-ray spectra of AGNs can be used to compare the "classical" black-hole type model for such objects with the starburst scenario emphasized in this meeting. The most important aspects of the observations are the commonly observed soft X-ray absorption, the Ka line profile and intensity, the X-ray variability and the hard X-ray cut-off. Modelling the observed X-ray spectrum requires an understanding of the absorption, emission and scattering properties of neutral and ionized gases. Examples from new calculations, including all these components, are shown and compared with the observations. Progress made on the observational and theoretical sides seem to give at least some satisfactory answers to previously open questions of the black-hole model. It remains to be seen whether the starburst model can come up with an equally good explanation.

1. Introduction The purpose of this review is to summarize some well known facts, as well as new Xray observations of various types of AGNs, and to confront them with suggested models. Since X-ray properties involve emission, absorption and scattering it is useful to stop, at each sub-section, and ask the question, What are the physical properties of the emitting, absorbing and scattering material? These are related to the location of the various components, their motion and other properties. They can provide an answer to the central question posed in this meeting, about the relationship between "classical" AGNs, hosting a giant black-hole, and violent star-forming events in nuclei of galaxies. 2. The optical-UV-X-ray spectrum of A G N s 2.1. The multi-wavelength continuum The relationship between the optical, TJV and X-ray continua is a key issue in understanding the AGN phenomenon. In particular, the relative flux in those wavelength ranges may differ from one subclass to the next. This may be due to either real, intrinsic differences between groups or due to absorption and/or obscuration in the line of sight. The IR, optical and UV observations and instrumentation, will not be discussed in this chapter. As for the X-ray, it is important to understand the limitation of various experiments, especially when discussing spectral features such as emission or absorption lines and the shape of the X-ray continuum. Presently there are only two big, operating X-ray missions, ROSAT and ASCA. The ROSATspectral mode (PSPC) is of very limited spectral capability over the range of 0.12.4 keV. Its typical energy resolution (E/AE) is about 2-4 and the detection of emission lines is essentially not possible. The detection of sharp absorption features (edges) is also difficult. Because of its restricted energy range, poor resolution and high sensitivity to the (rather uncertain) amount of neutral column density (Galactic and intrinsic) in the line of sight, it is also not trivial to obtain a good estimate of the overall continuum shape from ROSAT PSPC measurements. The second, more advanced, mission is ASCA 379

380

Netzer: X-ray Emission in AGNs

with two types of detectors and much improved energy resolution (E/AE ~ 50). ASCA SIS operates over the 0.5-10 keV range but partial charge makes it difficult to detect sharp spectral features at the low-energy range (Netzer, Turner & George 1994). Older experiments that were used to study the X-ray spectrum of AGNs include the Einstein IPC, EXOSAT and GINGA. 2.2. The X-ray properties of AGNs The following is a brief description of the characteristics of the X-ray properties of various groups of AGNs, excluding BL-Lac objects. The most important issues are the continuum shape, the absorption features and the strength of the iron Ka line. Quasars: The soft X-ray continuum has been studied with Einstein, EXOSAT and ROSAT. Little is known about the hard X-ray continuum, since most pre-ASCA missions did not have enough sensitivity at those energies. The little information that is available suggests a power-law continuum with a photon index F (Nv ex. hi/~r) in the range 0.7-1.0 (e.g. Shastri et al. 1993 and references therein). The soft (0.2-1 keV) continuum may be steeper. There is very little information on the strength of the Ka line and indication of strong X-ray absorption in a few cases (e.g. Laor et al. 1994 and references therein). Seyfert 1 galaxies: Many Seyfert Is have been observed by Einstein, EXOSAT, GINGA, ROSAT and BBXRT. The hard X-ray continuum is similar to the one observed in quasars and the soft X-ray continuum somewhat steeper. Strong absorption features, centered at around 1 keV, are observed in many low-luminosity sources (e.g. Reichert et al. 1985, 1986). The Ka line is observed in many sources with a typical equivalent width of 100-300 eV. Seyfert 2 galaxies: Observations by GINGA, ROSAT and other missions show a variety of properties in a handful of objects. A power-law continuum is present in some objects with extremely strong (EW of 1000-2000 eV) Ka line. Some Seyfert 2s show a curved X-ray continuum with absorption-like features of various energies. LINERs: There are extremely few observations with good enough statistics to reveal the shape of the X-ray spectrum. The recent study of M81 by Petre et al. (1993) shows a relatively steep power-law spectrum (F ~ 2.2) up to about 10 keV. 2.3. X-ray and unified schemes The above-mentioned X-ray properties fit well with the unified scheme for AGNs (e.g. Antonucci 1993) where obscuration by a central "torus" can explain the difference between the various subgroups. In particular, the "bare" Seyfert galaxies, of type 1 are those where we have a clear line-of-sight to the very center. These sources show best the intrinsic properties of the group. The obscured AGNs (Seyfert 2s) are seen by reflection only. This is consistent with the much larger strength of the Ka line in Seyfert 2s since the scattering medium is directly observed (see below and Marshall et al. 1993). As for the strength of absorption features, it has been suggested (e.g. Turner et al. 1993b) that the obscuring torus become transparent at hard X-ray energies and in those cases one observes the intrinsic central X-ray source, through the torus, above a certain energy. The variable absorption energy is thus a measure of the torus thickness.

3. Soft X-ray absorption 3.1. Absorption in quasars and Seyfert 1 galaxies X-ray absorption has been known to be common in low-luminosity AGNs (e.g. Turner et al. 1993a). Moreover, the recent study of PG quasars by Laor et al. (1994) shows this to be common also in at least some high-luminosity objects.

Netzer: X-ray Emission in AGNs

381

The commonly observed absorption features can be characterized by a low continuum level at ~ 1 keV with recovery on both sides. In principle, the detection of such features is simple if the spectral resolution is high enough. In reality, it is rather complicated in low-resolution, limited energy range experiments such as ROSAT PSPC. In such cases, the suppression of the continuum below 1 keV is easy to detect but the recovery on the high energy side is difficult to deduce since there is not much signal above 2.2 keV. The reality of the suspected absorption is thus deduced from a model fit to the data and its deviation from a simple power-law continuum. Marginal cases, corresponding to a small column density absorber, can be mistaken for a curved (e.g. a broken power-law) continuum. Other experiments, like GINGA, can easily detect the high energy continuum curvature due to large column density absorbers but give no information below ~2 keV. This situation is likely to improve dramatically with the successful operation of ASCA SIS, with its extended 0.5-10 keV range and much improved sensitivity and resolution. 3.1.1. Absorption by neutral gas ("cold absorber") Two explanations for the observed absorption have been proposed. The absorbing material can be neutral ("cold absorber", i.e. low level of ionization). In this case, the depth and strength of absorption features is a function of the absorbing column and metallicity. Model fits to several well-known cases (e.g. NGC 4151, see Weaver et al. 1994) show that the central continuum cannot be fully obscured by the neutral absorber since there is a clear recovery of the continuum below about 0.6 keV. This is the origin of the so called "partial covering" model, in which some 90% of the central source is covered by clouds and the remaining 10% is seen directly (Reichert et al. 1986). 3.1.2. Absorption by ionized gas ("warm absorber") Strong absorption features can also be due to ionized ("warm") absorbers that attenuates different parts of the continuum by different amounts (Halpern 1984; Kallman & Krolik 1984; Yaqoob, Warwick & Pounds 1989). In such cases, the depth of the absorption is maximum at ~ 1 keV and there is a recovery to the intrinsic continuum level below and above this energy. Models like this have been claimed to fit the observed spectrum of several AGNs . The ionized absorber model does not require partial covering, although it is possible to imagine a situation where obscuration by ionized gas is not complete. 3.2. New calculations In a recently published paper (Netzer 1993) I have discussed the various explanations for the observed absorption and gave a more complete analysis of the physical processes and the range of expected conditions. The paper provided an extensive grid of models of various types that can be directly compared with the observations of AGNs. Figure 1 shows some examples computed in that paper, and a new work soon to be submitted (Netzer, Turner & George 1994) illustrates a typical ionized absorber case. The diagram also shows the count rates expected in such a case from ASCA SIS observations. Given the observations of absorbed X-ray spectra, and the deduced column densities (10 2 2 ~ 2 3 cm" 2 ), the important questions to ask, regarding the starburst model, is where in the nucleus can we find such material and why the absorption is seen in some AGNs and not in others.

4. X-ray lines Of all X-ray lines, the iron Ka feature is expected to be among the strongest and has been, so far, the only line detected in the spectra of a large number of AGNs (see however

382

Netzer: X-ray Emission in AGNs

I

' • ' " i

I

i

I I I

I

1

i

I

ASC«

O

O

H

1 I II

I I 1I I

H

1—I

I I II

"I

i

• • •

«.-•

c

ASCA

ZJ

o O

T ir

I

10

1

|

till-

10

Energy (keV)

FIGURE 1. Left panels: Calculated X-ray spectra for a system of clouds each with a column density of 1022'5 era" 2 , solar composition and the ionization parameters in the range 10~ (lowest curve on each diagram, almost completely neutral gas) to 101'9 (top curves, highly ionized material). Spectra corresponding to a general case (including absorption, emission and reflection) are shown on the top and those with pure absorption are shown on the bottom. The right panels show the corresponding simulated spectra after convolution with the ASCA SIS detector. Note that the difference in the ASCA diagrams between those two cases is very small, due to partial charge effects. the possible detection of other lines, in Turner et al. 1991 and references therein). The situation is likely to change soon with the much-improved resolution ASCA observations, and the time is ripe to consider other X-ray emission features. 4.1. The iron Ka line The Ka iron lines have been observed in many AGNs with sufficient S/N but with rather poor resolution. Most of the available data are from the extensive GINGA survey, with only few exceptions. Such observations could not reveal the exact line energy and thus the level of ionization of the emitting gas (a key issue in understanding the physical conditions) was not known. The physical conditions for the formation of the line have been discussed in numerous papers (e.g. George k Fabian 1991, and references therein) but the observed strength (or equivalent width, hereafter EW) is not fully accounted for.

Netzer: X-ray Emission in AGNs

383

There are several possible locations for the line formation: In an accretion disc around the central source (Seyfert Is and quasars), in the Broad Line Region (BLR) of Seyfert Is and Quasars, or in the scattering medium at about the distance of the Narrow Line Region (NLR) in Seyfert 2s. Of the very few pre-ASCA observations with sufficient resolution, the NGC 4151 ones by Weaver et al. (1992) are perhaps the most interesting since they show a combination of a very broad component (from the accretion disc?) and a much narrower core with typical BLR width. Both components are consistent with fluorescence by low ionization (i.e. Fel7 or less ionized) gas. As for Seyfert 2s, the BBXRT observations of NGC 1068 reveal a very interesting multi-component structure for the line and evidence for a very high-temperature gas. The work by Marshall et al. (1993) suggests that the observed Ka complex is made out of three components: a neutral (6.4 keV) line, and two highly ionized species, Fe25 and Fe26. Marshall et al. show that the observations are consistent with a two-component scatterer; one with a temperature of about 200,000 K and the other with a temperature of about 5x 106 K. The lines are very strong, with a combined EW of about 2 keV, consistent with the idea of a clear line of sight to the line formation zone and obscuration of the central X-ray source. There are indications for a large iron abundance and small oxygen abundance in this source. 4.2. New Calculations As explain earlier, absorption by highly ionized material is one likely explanation for the observed spectrum of some AGNs. Such absorbing material must also produce emission lines and other line-like features. This has not been taken into account in almost all old calculations and the effect of the emission on the observed strength of absorption edges was not included in the calculations. The recent paper by Netzer (1993) takes full account of all such processes and calculates a combined spectrum for all X-ray components. Out of the newly calculated features the strongest are found to be: (a) La lines of H-like species, CVI, NVII, OVIII etc. (6) La lines of Hel-like ions, CV, NVI, OVII etc. (c) Iron L lines (Fe 17-22) that are most prominent around 0.9-1.1 keV. (d) Line-like features due to bound-free emission of H-like and He-like ions (these are very sharp due to the relatively low temperature of the photoionized gas). All these and other emission lines are included in the calculations shown in Figure 1. Their expected strength and influence on the ASCA observations can be seen in the top two panels. The important questions regarding the starburst model are: where in the nucleus can we find gas that produces both narrow and broad Ka lines, what is the origin of the large range of ionizations and temperatures observed in NGC 1068 and how to explain the other emission features.

5. X-ray scattering 5.1. Compton scattering at high energies X-ray scattering off neutral gas has been discussed in the seminal work of Lightman & White (1988) and in several papers since. This process is thought to contribute to the X-ray continuum and to flatten the hard X-ray spectrum. The seed X-ray photons originate in a small central source and scatter off the surface of an accretion disc. The scattering is angle dependent and the central source geometry defines most of its properties. Scattering off low-ionization gas is almost negligible at energies below 5 keV due to the high absorption opacity of the gas.

384

Netzer: X-ray Emission in AGNs 5.2. Compton scattering at low energies

If highly ionized material is indeed present in the vicinity of the central X-ray source, it must also scatter the X-ray photons at low energies and affect the observed depth of absorption features. This has not been considered until recently. The new calculations by Netzer (1993) take full account of this process and demonstrate that X-ray absorption is indeed less pronounced. The scattering efficiency depends on the level of ionization and the covering factor of the gas. The models in Figure 1 include this effect for the specified covering factor and column density shown (compare top and bottoms panels). Scattering can be most effective in changing the observed spectrum if the scatterer is highly ionized gas on the surface of a thin accretion disc. The overall spectral appearance, in this case, shows a rise of the spectrum at low (0.5-1 keV) energies above the original power-law continuum.

6. Hard X-ray cut-off and accretion discs 6.1. The new OSSE observations A key issue in understanding the X-ray properties of AGNs is the highest observed X-ray energy. Observations obtained until 1991 were limited to about 35 keV (the GINGA limit) at which energy there was no clear change of slope. The newly available OSSE observations shed new light on this question. It is now evident that the high-energy X-ray cut-off has been observed, at least in several AGNs, with a typical energy of about 100 keV (e.g. Maisack et al. 1993). This extremely important new information can be used to evaluate the total energy in the X-ray range and is also important for understanding the X-ray background (e.g. Zdziarski, Zycki & Krolik 1993). 6.2. The new thermal inverse-Compton model While there is still a debate about the origin and the nature of the hard X-ray component in AGN spectra, there is at least one model that explains, in a natural way, the shape and the cut-off energy for such sources. The model, which is explained in a recent paper by Titarchuk (1994), involves Compton scattering by high-energy thermal electrons, presumably in a hot corona above the central accretion disc. It gives a remarkably good fit to the OSSE observations of NGC 4151 and is likely to produce an equally good fit to other sources. The model predicts plausible values for the temperature and geometry of the hot plasma, and its Compton depth. This has important implications for the physics of AGN accretion discs. Regarding the starburst scenario, one must find an explanation for the new OSSE observations and a different physics to explain the combination of low-energy power-law continuum with a ~100 keV cut-off. 7. X-ray variability Most Seyfert 1 galaxies and many quasars are fast X-ray variables with variability time that is inversely proportional to the X-ray luminosity (e.g. Green et al. 1993). The variability amplitude is very large, a factor of two or more in many sources. The deduced source dimension must be very small, down to about one astronomical unit in some low-luminosity objects. Figure 2 shows data from Netzer, Turner & George (1994) that demonstrate time variation in two Seyfert 1 galaxies. An interesting and important feature is the change of "color" (i.e. continuum shape) associated with the observed variability. The top panels of Figure 2 demonstrate this for the ROSAT PSPC observations. The variation is such that the fainter phase is associated

Netzer: X-ray Emission in AGNs 20

I

I

I

I

I

I

I 1

I

I

I

I

I

1 T

i

t

20

Mkn 766 X

i

385 i

i

T

i

i

r

NGC 4051

15

X

I

!

g 10

15

i

o

o 5 2.5

I

I

I

i

i

i

I

1 I

i

i

i

i

I

i

i

i

i

1 I

I

I

I

I

I

I I

i

10 3

i

I

I

1

I

i

I

(

T

^

I

I

•o

I 1.5

ts. -

X) O

i

.5

I

I

I

i

I

I

105 Time

Time

FIGURE 2. Observed XLi count rates, and the XL2/XL1 color ratio for two Seyfert 1 galaxies (from Netzer, Turner & George 1994), where XLi is the ROSAT band of 0.2-0.6 keV and XL2 the 0.9-1.3 keV band. Note the fast time-scale variation and the inverse correlation between the count rate and XL2/XL1, indicating continuum flattening during high-luminosity phases.

with a steeper continuum, i.e. the variability is more pronounced at higher energies. This seems to be the case also in several other sources not shown here. The luminosity dependence of the variability time-scale and the "color" change of the X-ray continuum, are key issues in comparing black-hole type and starburst type models for AGNs. Any successful model must account for both given the other observed properties. A natural explanation in the black-hole scenario is the larger dimension associated with the more luminous central source. 8. Starburst activity and X-ray properties The various aspects of X-ray properties discussed in this review lead to a long yet incomplete list of questions that must be addressed by any successful model. In my opinion, the most important ones are: (a) The multi-wavelength continuum and the unified scheme: What is the reason for the different energy budgets of the different AGNs? How can the observed difference between the X-ray properties of Seyfert 1 and Seyfert 2 galaxies be explained? (6) X-ray absorption: What is the location of the X-ray absorber? Is the observed column density consistent with the suggested models? Why is absorption strong in some objects and not in others?

386

Netzer: X-ray Emission in AGNs

(c) T h e Fe lines: Where and how is the K a line formed? Is the observed equivalent width consistent with the model? Is the observed range in this property consistent with the model (i.e. what is the origin of the much larger equivalent width in Seyfert 2 galaxies)? What is the origin of the observed Fe25 and Fe26 lines in NGC 1068 and what process is responsible for the hot ( ~ 5 x 106 K) component? (
REFERENCES ANTONUCCI, R. 1993 Ami. Rev. Astr. Ap. 31, 473. GEORGE, I. M. & FABIAN, A. C. 1991 M. N. R. A. S. 249, 352. GREEN, A. R., MCHARDY, I. M. M. & LEHTO, H. J. 1993 M. N. R. A. S. 265, 664. HALPERN, J. P. 1984 Ap. J. 281, 90. LAOR,

A.,

FIORE,

F., ELVIS, M., WILKES, B. J. &; MCDOWELL, J. C. 1994 Ap. J. In press.

LIGHTMAN, A. P. & WHITE, T. R. 1988 Ap. J. 335, 57. MARSHALL, F. E. et al. 1993 Ap. J. 405, 168.

MAISACK, M. et al. 1993 Ap. J. Lett. 407, L61. NETZER, H. 1993 Ap. J. 411, 594. NETZER,

H., TURNER, J. & GEORGE, I. M. 1994 Ap. J. Submitted.

PETRE, R., MUSHOTZKY, R. F., SERLEMITSOS, P. J, JAHODA, K. & MARSHALL, F. E. 1993

Ap. J. 418, 644. REICHERT, G. A., MUSHOTZKY, R. F., PETRE, R. &; HOLT, S. S. 1985 Ap. J. 296, 69. REICHERT, G. A., MUSHOTZKY, R. F. & HOLT, S. S. 1986 Ap. J. 303, 87. SHASTRI, P., WILKES, B. J., ELVIS, M. & MCDOWELL, J. 1993 Ap. J. 410, 29.

TlTARCHUK, L. 1994 Ap. J. Submitted. TURNER, T. J., WEAVER, K. A., MUSHOTZKY, R. F., HOLT, S. S. &; MADEJSKI, G. M. 1991

Ap. J. 381, 85. T. J., GEORGE, I. M. & MUSHOTZKY, R. F. 1993a Ap. J. 412, 72. TURNER, T. J., URRY, C. M. & MUSHOTZKY, R. F. 1993b Ap. J. 418, 653. TURNER,

WEAVER, K. A. et al. 1992 Ap. J. Lett. 401, Lll. WEAVER,

K. A. et al. 1994 Ap. J. In press.

YAQOOB, T., WARWICK, R. S. & POUNDS, K. A. 1989 M. JV. R. A. S. 236, 153. ZDZIARSKI,

A.,

ZYCKI,

P. T. & KROLIK, J. 1993 Ap. J.Lett. 414, L81.

Starbursts and Compact Supernova Remnants By J.FRANCO 1 , S.J.ARTHUR 1 1 2

AND W. MILLER2

Institute) de Astronomi'a UNAM, Apaitado Postal 70-264, 04510 Mexico D.F., Mexico

Department of Astronomy, University of Wisconsin - Madison, 475 N. Charter St., Madison, WI 53706, USA

In this paper we discuss the differences between normal supernova remnants (SNRs) evolving in comparatively low-density ambient media (no 104 cm" 3 ) environment that one would expect to find in starburst regions of galaxies. For normal SNRs, radiative losses do not start to become important until time scales of the order of 104 yr, after the onset of thin shell formation. For compact SNRs, however, the evolution proceeds at a much quicker pace, with radiative losses due to free-free emission, given the high temperatures (> 107 K), being important because of the high densities. The onset of thin shell formation in this case occurs over time scales of the order of years, and most of the radiation is emitted in X-rays and the UV. We argue that the compact supernova activity associated with starburst regions in the centers of galaxies gives rise to most of the typical properties of the Broad Line Regions of active galactic nuclei.

1. Introduction Starbursts are usually traced by their bright photoionized regions and large cluster luminosities, in either optical or IR wavelengths. The mass spectrum of the resulting stellar groups is difficult to derive but a large fraction of massive stars is usually implied. Stars with initial masses above ~ 8 M© have strong UV radiation fields and significant mass loss during their whole evolution, and they are also the progenitors of Type II and Ib supernovae (SNe). This pre-supernova activity is not restricted to any particular evolutionary phase and provides a large flux of photoionizing radiation as well as large amounts of heavy elements and mechanical energy into the interstellar medium. Thus, the collective action of stars from starbursts can generate the complex phenomena discussed in this conference, but each individual star creates its own structure in the local surrounding medium. Direct evidence for the actual influence of single massive stars includes a series of different features ranging from optical shells (both in the Galaxy and in the Magellanic Clouds) to interstellar structures around young pulsars (e. g. Braun, Goss & Lyne 1989; Le Coarer ei o/. 1993; Rosado, Le Coarer & Georgelin 1993). Thus, the density distribution in the neighborhood of Type II supernova explosions is expected to be fairly complex (see recent review by Franco 1994), and the interaction between the ejecta and the surrounding medium does not follow a simple evolutionary track. This interaction and the resulting phenomena have a number of transient properties and the evolution is largely determined by the average density of the ambient medium. At low densities, as in the case of SNe exploding in the galactic interstellar medium, the evolution proceeds at a slow pace and the thin shell formation time is long, i. e. several times 104 yr for densities around 1 cm" 3 (e. g. Franco et al. 1994). Remnants evolving in a denser medium, however, have higher cooling rates and reach higher luminosities over much shorter time scales. This is the case for progenitor stars embedded in dense circumstellar or interstellar structures, such as HII regions or wind-driven shells, and represents the most logical environment for the SN activity in starbursts. Given the high interstellar densities and pressures expected in the nuclear regions of galaxies, the 387

388

Franco et al: Starbursts and Compact Supernova Remnants

resulting phenomena can be extreme in the case of nuclear starbursts. Here we review some of the main features of SNR evolution in these dense interstellar conditions and discuss the effects of thin shell formation in their surrounding ambient medium. One of the most striking aspects is the resemblance of these features with the activity observed in AGN. The spectra of SN and "normal" galactic remnants do not resemble the spectra of AGN, but recent observations show that the optical spectrum of at least some luminous SNe exploding in dense regions have a striking resemblance to the spectrum of broad line Seyfert galaxies (e. g. Filippenko 1989). Similarly, the flares of some Seyfert galaxies have the luminosity, lifetime and spectral signatures of Type II SNe (Terlevich &; Melnick 1985). The fundamental difference between "Seyfert-like" (or "Seyfert impostors") and normal Type II SNe can be understood if the former are associated with shocks expanding into regions of high circumstellar gas densities.

2. Thin shell formation and compact supernova remnants 2.1. "Normal" supernova remnants The expansion of an SNR at low densities follows several well known stages. These stages, usually termed "free-expansion", "Sedov track" and the "momentum conserving stage", represent the main features in the life of a "normal" SNR and they have been studied by a number of analytical and numerical means (see review by Franco 1994). Ejecta thermalization occurs via a reverse shock and leads to the quasi-adiabatic Sedov track. After the ejecta have been thermalized, the hot remnant expands supersonically doing p — V work in a quasi-adiabatic mode. Most studies of SNR evolution do not consider the thermalization stage and simply assume that the SN energy has been already thermalized and is driving a blast wave into the surrounding medium. Radiative cooling is unimportant during this quasi-adiabatic stage because, even when the expanding gas reaches very high temperatures (of order 107 K or higher), the densities within the structure are very low and the cooling time is long. A recent discussion of the relevant time scales for remnants evolving at low densities is given in Franco et al. (1994), and here we just summarize some of the main results. The cooling time of a gas element with temperature T in the optically thin case is 'cool — o

^ T

— I

1C

,

I

rj

2Aon \ 16K / 2Aon where C is a constant, n is the gas number density in the postshock flow, m p is the mass per gas particle, and A(T) ~ AoT'3 is the cooling function. The exponent /? depends on both the heavy-element abundances and temperature range. For a plasma with cosmic abundances and temperatures below 3 x 107 K, line cooling from heavy elements becomes dominant and /? ~ —0.5 (e. g. Kahn 1976; Sutherland & Dopita 1993). Given that the temperature is proportional to the square of the shock velocity, the cooling time for the case of equilibrium line cooling goes as '

(2 2)

'

where vs is the shock speed, A = Aok~l/2mp' , and po = mpn0 is the ambient mass density. For an SNR with an initial energy Eo and an age ts on the Sedov track, the radius grows as Rs(t) = (2.Q2Eo/poY^t2* , and the cooling time decreases rapidly as a — 9/5

function of the evolutionary time, as ts

.

Franco et al.: Starbursts and Compact Supernova Remnants

389

When the remnant evolutionary time, ttot, becomes larger than 1> the gas elements can adjust to the temperature drop and the cooling process operates quasi-isobarically at about the post-shock pressure. In the other case, when tCooiAsound < li radiative cooling occurs faster than any pressure readjustment and the process becomes quasi-isochoric at about the post-shock density. Hence, a large pressure imbalance develops in the flow, and the cool gas is compressed by a series of new additional shocks. This "catastrophic cooling" regime (Falle 1975, 1981) appears during thin shell formation but the instabilities continue to operate during the rest of the evolution. These cooling instabilities operate in the post-shock region (e. g. Vishniac 1983) and produce variations in the emission. The onset of thin shell formation, then, marks the epoch of strong emission, preludes the appearance of instabilities, and represents a distinct moment in SNR evolution. To derive the onset of thin shell formation one has to find the minimum value of ttotFor an ambient density distribution p = Br~w, given by O3/20-5

t

U)^

u

f

5-w

The evolution of ttot for three values of w are shown in Figure 1. For the constant-density case (w = 0) the onset of thin shell formation occurs at time

where E$i is the energy in units of 1051 erg. This time marks the beginning of the strong cooling epoch, while the maximum luminosity is achieved at about 1.65tfsf, sharply decreasing by ~ 1.8tsf, indicating that a large fraction of the energy (about 30% ; see Tenorio-Tagle et al. 1990) has been radiated away by ~ 1.8tsf. Thus, the gas is compressed into a thin and cool shell in a time scale of about 0.8tsf and the average luminosity during the formation of this shell is

^ l ^ x l O ' ^ V 7 L0. 3

(2.5)

This approximation is not valid for densities below 0.1 cm" , and the peak luminosity is ~ 2 times this value. The average radiative flux from the remnant during shell formation is simply (F) = (L)/(47ri?g) and scales as E^14nl°^7. The average temperature of the emitting gas is close to the post-shock temperature and scales as E% «O • These dependences show that both the gas temperature and the outgoing flux at the strongly radiating epoch increase with increasing ambient density. As a consequence, the spectra become harder and the photoionizing flux increases as the ambient density increases. Thus, SNRs evolving in high-density media undergo a rapid radiative cooling, generating a large ionizing flux and thereby creating a significant photoionized region ahead of the

Franco ei al.: Siarbursts and Compact Supernova Remnants

390 a)

b)

u == 0

3 -

3 -/ t

tot/ -a.min

.V_* //

2 -

/ /

\

t

-

/

'•s.min

/ / /

1 /

1 -

/

0

1

1

2

1

1

3

4

0

1

2

3

4

S' s.min

c)

t

A

10

Hot' Ls.min

0

1

2

3

4

s' a.min FIGURE

1. Evolution of itot, normalized by <s,min, for three sample cases.

shock front. Such a consequence should have an obvious impact in regions with vigorous star formation, because at densities of the order of 104 cm" 3 the average luminosity of a single remnant reaches values above 107 L©, which is similar to the luminosities of massive clusters. 2.2. Compact supernova remnants The previous formulation is not applicable when the strong cooling regime occurs before ejecta thermalization is completed. This is the case for ambient densities above ~ 105 cm" 3 , and one loses the simple formalism used in the previous section as there are no analytical solutions for the shock structure at the shell formation epoch. Remnants evolving in these high densities exhibit very high luminosities as early as several months after the explosion. Indeed, thin shell formation can be completed within months after the strong cooling epoch begins (e. g. Terlevich ei al. 1992, 1994; Tenorio-Tagle et al. 1993).During this phase the ejecta end up radiating away their thermalized kinetic energy as soon as they traverse the reverse shock. Efficient cooling at early phases can dispose of large fractions of the SN energy, in effect eliminating the quasi-adiabatic Sedov track. Unlike stable adiabatic shocks, radiative shocks suffer from cooling instabilities that are likely to lead to time variability of the emergent flux. Because the ambient

Franco et ah: Starbursts and Compact Supernova Remnants

391

densities are so large, an unsteady ionizing flux will result in photoionized regions that have drastically variable size and emission measures. Given the conditions briefly stated above, SNRs evolving at high densities can be expected to have a dramatic impact on their surrounding environment with distinct observable features. But the question of their existence must first be answered, thus we mention two new kinds of peculiar SNe which show evidence for evolution at a very high density (~ 10 7 cm" 3 ) circumstellar media: the luminous Type II radio SNe and the so-called Seyfert-like or Seyfert I impostors (see the review by Terlevich 1994). As mentioned above, the time scale on which strong radiative cooling acts to form a thin dense shell in these compact SNRs is very short, and it commences a relatively short time after the initial SN explosion. One then has to be careful to differentiate the SN event itself, which is bright because the photosphere grows as the stellar envelope is expelled after the explosion, from a very young and compact remnant, which radiates its energy away in a short time after the explosion. The SN event is characterized by the photospheric emission in a series of absorption lines, whereas the compact remnant is photoionized by its own cooling region and is characterized by the emission lines from the photoionized gas. It is worth noting some of the details of the cooling region and the time scales involved in a compact remnant, and we do so very briefly below. The interaction of the ejecta with the dense circumstellar medium generates two shocks: an outward shock sweeping through the circumstellar gas and a reverse shock back into the ejecta. The outgoing and reverse shocks have effective velocities ranging between ~ 104 to ~ 10 3 km s" 1 , and raise the gas temperatures to values between ~ 109 and ~ 10 7 K. With these large post-shock temperatures, the plasma is almost fully ionized and the cooling is dominated by free-free emission; the cooling function is then given by A(T) ~ AoT 1 / 2 . The resulting cooling times are proportional to T 1 / 2 , and thus scale linearly with the shock velocity tcoo, ~ 0.2v 8 /n 7 yr, 8

1

(2.6) 7

3

where i>8 = v s /10 cm s" , the pre-shock density is nj = no/10 cm" , and we have assumed C = 1 in Eq. 2.1. The shell formation time cannot be obtained in a simple closed form as before, but can be estimated with a simple approximation (see Franco et al. 1994). Defining the ratio of the cooling column density to the column density of hot gas behind the shock front

(fro)- 1 *!^ 1 ,

(2-7)

where £ is a geometrical factor that takes into account the expansion of the shocked gas (£ = 1 for planar shocks, and decreases below 0.3 for spherical shocks), and Ri$ — Rs/l016 cm is the shock radius. Figure 2 shows the resulting column densities for 3 test cases and free-free cooling. The onset of thin shell formation occurs at F = 1 and for a constant velocity shock is given by ttt = 0.26-|- yr.

(2.8)

The shock velocity is obviously not constant, but the cooling times are so short that they allow only small changes in shock speed during a cooling time, and one can safely use this approximation. The adequate value of £ from numberical models turns out to be

392

Franco et ai: Siarbursts and Compact Supernova Remnants p- +0.5

no=0.1

N,

|,-..o

;

o

.

-\ \ "\

-

' . .

.

:

•

to

LOG (column densi

X 20

•

/

* ' • • • .

' ....

" • ' • •

18

i

17

2x10" time (years) FIGURE

3x10"

2. Evolution of the cooling and swept-up column densities for three cases with constant ambient density.

larger than ~ 0.1. Thus, for densities of nj = 1, thin shell formation on both the forward and reverse shocks occurs only a few years after the explosion. 2.2.1. Temperature structure All flow variables change rapidly with time under these conditions. Given that radiative losses dominate over the temperature drop due toflowexpansion during the strongly radiative epoch, one can approximate the temperature structure using the radiative cooling rate as the sole loss term. In this case, defining e as the thermal energy density, the rate of temperature change behind the shock front is dT/dt ~ 2e/(3fc). Further assuming an isochoric cooling regime (a very good approximation in this case where the cooling time is much shorter than the dynamical time) with constant flow velocity, the temperature gradient is then 16n A ,

0 0 (2.9) dr 9kv, where Ts is the post-shock temperature. The resulting temperature distribution for va > 1.5 follows as

8noAor (2.10) T(r) ~ T.\ 1 K, ~ 1.4 x 9kv.T;"\ L where ri 5 is the location of a cooling parcel in units of 1015 cm measured downstream from the shock front (i. e. inwards for the outgoing shock and outwards for the reverse

Franco et al.: Starbursts and Compact Supernova Remnants

393

shock). This distribution is valid up to the point when line cooling begins to dominate as we have used a cooling law based purely on free-free emission. Assuming equilibrium cooling conditions, this point corresponds to the location where the temperature reaches a value of about 3 x 107 K and is given by r(T = 3 x 107) =

9kvsT}/2

cm. (2.11)

8n0A0

2.2.2. The emergent spectrum and line ratios The ionizing spectrum emitted by the cooling region can be easily obtained for the free-free emission case. Given the high post-shock temperatures, most of the emission comes out in X-rays. The free-free emission per unit volume and unit frequency is (e. g. Lang 1980) /„ ~ 3.7 x 10- 38 n 2 T- 1/2 ln(4.7 x 10 lo r/j/)exp(-/ii//fcT) erg cm" 3 s" 1 Hz" 1 , (2.12) where the Gaunt factor has been approximated by the logarithmic term and the variables are in cgs units. The resulting flux passing through each shock front is then -l

where AR is the size of the cooling region and the temperature gradient is given above. This integral has no analytical solution, but can be easily estimated numerically providing a useful estimate of the emergent ionizing spectrum. The resulting emission compares well with the UV/X-ray emission obtained with more detailed, high-resolution hydrodynamical models (Franco et al. 1993b). Given the wide range of gas temperatures, emission above one Rydberg mimics a power-law spectrum and the resulting photoionized region has line ratios similar to those found in the Broad Line Regions of active galaxies (Terlevich et al. 1992). Also, the X-ray emission from the compact remnant develops features, such as the Fe K a line, that resemble those observed in Seyfert galaxies (Franco et al. 1993b). Figure 3 illustrates the X-ray spectrum of a compact SNR evolving in a medium with no = 107 cm" 3 at the age of 3.1 yr (see Franco et al. 1993 for a model description). Another interesting property is related to the time taken by the cool shell to form. After ts{ the events proceed on very short timescales, the luminosity increases rapidly and the large photon flux is able to self-ionize the cooling gas. As a consequence, some time after the UV emission increases to its peak value, the relative ratios of the photoionized species vary and some lines reach their peak luminosity. Thus, a well-defined lag between the variations of the UV continuum and the emission lines is expected (Terlevich et al. 1994). Again, such a lag is similar to the one observed in well-monitored AGN. 3. Discussion One of the most interesting characteristics of compact SNRs evolving at high densities is that they reach high luminosities a short time after the explosion, when the two existing shocks are moving at speeds of several 103 km s""1. Given these large shock velocities, most of the energy radiated by these remnants is emitted in the extreme UV and X-ray regions of the spectrum. Their impact on the surrounding medium is similar to that of a whole stellar cluster, but their photoionizing spectra are quite different, resembling a

Franco et al.: Starbursts and Compact Supernova Remnants

394

t = 3.1 years

log(E [eV])

3. Spectrum, binned in 10 eV bins, of a compact SNR moving in a constant ambient density medium with n0 = 107 cm"3 at 3.1 yr after explosion.

FIGURE

"power-law" spectrum similar to that required to generate the line ratios observed in the Broad Line Regions of AGN. Such a connection between SNRs and AGN is better understood after the discovery of a new class of supernovae with broad emission lines, the Seyfert-like SNe, resembling Seyfert galaxies (Filippenko 1989). These compact remnants are also expected to occur in galaxies undergoing a nuclear burst of star formation, where the nuclear activity is a direct consequence of the evolution of a massive and young stellar cluster. Thus, one can trace a logical line between nuclear starbursts, compact SNRs, and AGN activity. In the starburst scenario for AGN, these compact SNRs are the Broad Line Regions (see Terlevich et al. 1992, 1994). The idea that starbursts and supernovae could be powering the activity observed in AGN is not new, and has been discussed for several decades (e. g. Shklovskii 1960; Field 1964). As stated above, the properties of compact SNRs (e. g. luminosities, emission line ratios and lags) are very similar to those considered typical of the BLRs. Furthermore, the "natural" variability of cooling shocks adds a relevant ingredient to this scheme because it provides a low-amplitude variability with time scales of months, probably even weeks. Another source of variability comes from the expected properties of the ejecta: the ejecta are not smoothly distributed plasma but rather a collection of high-density condensations or fragments (see Franco et al. 1993a). The cooling time for fragments is shorter than for the inter-fragment medium, and the fragment cooling will add a modulation to the X-ray emission. The relevant time scales in this case are not fixed by a given physical process but depend on the fragment distribution within the ejecta. Low-luminosity highly variable Seyfert Type 1 nuclei probably host about one

Franco et a/.: Starbursis and Compact Supernova Remnants

395

compact SNR at any one time, whereas the less variable and most luminous QSOs may contain several tens of coexisting compact SNRs. We would like to thank our very good friends and collaborators Don Cox, Jorge Melnick, Michal Rozyczka, Guillermo Tenorio-Tagle, and Roberto Terlevich for many years of exciting collaboration. JF was partially supported by DGAPA-UNAM through the grant IN103991 and a CRAY R&D grant. WM was supported by the NASA grants NAG5-629 and NAGW-2532, and thanks the Instituto de Astronomia-UNAM for their hospitality. Part of the work reported here was obtained with the CRAY/YMP of the Supercomputing Center-UNAM.

REFERENCES R., GOSS, W. M. & LYNE, A. G. 1989 Asttophys. J. 340, 355. CIOFFI, D. F., MCKEE, C. F. & BERTSHINGER, E. 1988 Astrophys. J. 334, 252. FALLE, S. A. E. G. 1975 Mon. Not. R. Astron. Soc. 172, 55. FALLE, S. A. E. G. 1981 Mon. Not. R. Astron. Soc. 195, 1011. FIELD, G. B. 1964 Astrophys. J. 140, 1434. FILIPPENKO, A. 1989 Astron. J. 97, 726. FRANCO, J. 1994 Rev. Mex. Astron. Astrophys. In press. FRANCO, J., FERRARA, A., ROZYCZKA, M., TENORIO-TAGLE, G. & Cox, D. 1993a Astrophys. J. 407, 100. BRAUN,

FRANCO, J., MILLER, W.,

COX, D. P., TERLEVICH, R.,

ROZYCZKA, M. &

TENORIO-TAGLE,

G. 1993 Rev. Mex. Astron. Astrophys. 27, 133. FRANCO,

J.,

MILLER,

W.,

ARTHUR,

S. J.,

TENORIO-TAGLE,

G.

& TERLEVICH,

R.

1994

Astrophys. J. Submitted. FRANCO, J., TENORIO-TAGLE, G., BODENHEIMER, P. & ROZYCZKA, M. 1991 Publ. Astron. Soc. Pac. 103, 803. KAHN, F. D. 1976 Astron. Astrophys. 50, 145. LANG, K. R. 1980 Astrophysical formulae. Springer. LE COARER, E. et al. 1993 Astr. Astrophys. In press. ROSADO, M., LE COARER, E. &; GEORGELIN, Y. P. 1993 Astr. Astrophys. In press. SHKLOVSKII, S. I. 1960 Sov. Astron. 4, 885. SUTHERLAND, R. S. &; DOPITA, M. A. 1993 Astrophys. J. Suppl. 88, 253. TENORIO-TAGLE,

G.,

BODENHEIMER,

P.,

FRANCO, J. & ROZYCZKA, M. 1990

Mon.

Not.

R.

FRANCO, J. & MELNICK, J. 1993

In

Astron. Soc. 244, 563. T E N O R I O - T A G L E , TERLEVICH, R.

J.,

ROZYCZKA, M.,

Star Formation, Galaxies and the Interstellar Medium (ed. J. Franco, F. Ferrini &; G. Tenorio-Tagle), p. 153. Cambridge Univ. Press. TERLEVICH, R. 1994 Preprint.

R. & MELNICK, J. 1985 Mon. Not. R. Astron. Soc. 213, 841. R., TENORIO-TAGLE, G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. R. Astron. Soc. 255, 713. TERLEVICH, R., TENORIO-TAGLE, G., FRANCO, J. & ROZYCZKA, M. 1994 Mon. Not. R. Astron. Soc. In press. VlSHNIAC, E. T. 1983 Astrophys. J. 274, 152.

TERLEVICH,

TERLEVICH,

Broad-Band and Line Emission from Fast Radiative Shocks in Dense Media By TOMASZ PLEWA Warsaw University Observatory, Al. Ujazdowskie 4, 00-478, Warsaw, Poland The evolution of fast, radiative shocks in a high-density medium is discussed. Approximate broad-band light curves of the shocked gas are calculated, and the emitted spectra are used as the input spectra for photoionization models. The results are in good agreement with parameters characteristic of Active Galactic Nuclei.

1. Introduction Over a decade ago Chevalier k Imamura (1982) showed that radiative, steady shocks are subject to an oscillatory instability. This result was confirmed on the basis of nonlinear hydrodynamical analysis for nonstationary shocks (Gaetz, Edgar k Chevalier 1988) as well as for steady radiative shocks (Innes, Giddings k Falle 1987). Both groups found unstable behavior of shocks faster than ~ 130 km s" 1 . In a series of papers Terlevic.h and collaborators (see e.g. Terlevich et al. 1992) developed the starburst model for Active Galactic Nuclei (AGNs). In this model AGNs are powered by compact, dense supernova remnants (cSNRs), and the bulk of the radiation is emitted by the supernova shock wave evolving in a dense medium (n = 107 cm" 3 ). Terlevich et al. (1992) calculated 1-D and 2-D hydrodynamical models of cSNR evolution and demonstrated that it was possible to recover observed characteristics of the Broad Line Region of AGNs. Using very simple assumptions, Terlevich et al. (1994a) successfully explained observed differences between times of maximum continuum and line emission taking into account the dependence on the ionization parameter. While rapid X-ray variability still remains to be explained in this model, preliminary investigations (Terlevich et al. 1994b) showed that interactions between dense, fast-moving clumps of the gas could be partially responsible for observed variations of the high-energy emission. Detailed criticism of the Starburst model was presented by Heckman (1991) and Filippenko (1992). Apart from AGNs, compact supernova remnants were detected in nearby galaxies. On the basis of spectroscopic studies, such objects as SN 1980K (Fesen k Becker 1990) or SN 1988Z (Stathakis k Sadler 1991; Turatto et al. 1993) are interpreted as supernova remnants evolving in a medium of density as high as 107cm~3. Therefore, the application of our theoretical models to more ordinary objects is straightforward. It should be noted that all theoretical investigations mentioned above apply to slow shocks (v, < 300 km s" 1 ) and low-density media (n = 1 cm" 3 ). Therefore, they are not applicable to conditions expected for cSNRs or for central regions of AGNs. The lack of appropriate models was also noted by Leibundgut et al. (1991) in their study of SN 1986J and SN 1980K. This observation gave motivation to our work on the hydrodynamical evolution of fast shocks evolving in a dense environment. 2.

T h e model A detailed description of the physical assumptions and numerical method employed is given in previous papers in the series (Plewa k Rozycka 1992; Plewa 1993). Here we will briefly present some details needed for a proper understanding of the model. 396

Plewa: Fast Radiative Shocks

reflecting wall (outer supernova shell)

shock front

397

inflowing gas

shocked gas

FIGURE

1. Schematic representation of the initial conditions for shock evolution.

The hydrodynarnical evolution is modeled using the Piecewise Parabolic Method (PPM) of Colella & Woodward (1984) in 1-D plane-parallel geometry. The high resolution of the PPM scheme is enhanced by concentrating grid points in regions of steep gradients. The grid motion is governed by an adaptive grid algorithm developed by Dorfi & Drury (1987). Our resolution function contains gradients of density, internal energy, and emission rate (see Dorfi & Drury 1987 for details). We also added the term proportional to density allowing for better resolution in dense regions. At the left boundary of the grid we imposed a reflecting boundary condition while a constant-density gas (71 = 107 cm"3) of temperature (Tam\, — 104 K) is introduced into the right side. The velocity of the inflowing gas varies with time as V{n ~ < ~ T to mimic the expansion of the outer supernova shell (Figure 1). This set of initial conditions leads to the formation of a strong shock at the left boundary. The initial shock speed for constant gamma-law gas with 7 = 5/3 is given by v, — | v t n . Radiative losses are allowed for in the post-shock region, and the shock moves to the right through the as yet unshocked gas. Energy losses are calculated implicitly under equilibrium conditions for an optically thin medium of solar composition. Our cooling function was calculated using the CLOUDY 84.06 code (Ferland 1993). The same code was used for the calculation of photoionization models. Spectra of the post-shock region were calculated as the sum of continuum and line emission. Continuum emission consists of f-f emission with an approximate Gaunt factor calculated after Hummer (1988), while b-f and 2-photon emissions were calculated using formulae given by Mewe, Lemen k, van den Oord (1986). The contribution from over 2600 lines was calculated using line emissivities taken after Stern, Wang k. Bowyer (1978) and Mewe (1992) (see also Mewe, Gronenschild k van den Oord 1985). In order to estimate the relative distribution of energy in the spectrum we defined a set of three filters corresponding to wavelengths A < 100 A (high energy, HE), 100 < A < 900 A (medium energy, ME), and above 900 A (low energy, LE). The filter is defined as a part of the total energy emitted per given wavelength interval at a given temperature (Figure 2).

Plewa: Fast Radiative Shocks

398

E

I CO

6

7 log 10 T[K]

FIGURE

2. Definition of the broad-band filters. HE, ME, and LE filter transmissions are drawn by thick, thin and dotted lines, respectively.

3. Results 3.1. Evolutionary phases The evolution starts with a nearly adiabatic shock expansion as cooling is ineffective at temperatures over 108K (see Raymond, Cox k Smith 1976 for a typical cooling function). As both the temperature of the radiating gas and the velocity of the inflowing gas decrease with time, the shock slows down while the cooling efficiency grows rapidly with temperature approaching ~ 2 x 107K. At temperatures lower than ~ 106 K the cooling time scale suddenly drops down. This phenomenon is known as catastrophic cooling and causes the formation of a thin transition region called a cooling wave, which radiate most of the gas energy in the ME band. In our previous study (Plewa 1993) we showed that any flow perturbations are strongly amplified by rapid cooling (cooling waves) making post-shock flow highly discontinuous. Strong pressure gradients develop between the cooling waves leading to the formation of a very dense, thin shell. Finally, the main shock stops in its motion with respect to the left boundary of the grid surface and begins to recede. For an initial shock velocity v, = 6000 km s" 1 the maximum extension of the post-shock region is Rmax = 1.2 x 1016cm and the total flux of radiation is Fmax = 1.5 x 108 erg cm" 2 s" 1 . At this time the radius of the outer supernova shell is equal to Rsh w 3 x 1016cm (Terlevich et al. 1992, Figure 8a) and the corresponding total shock luminosity is roughly equal to L, sa 1.7 x 1042 ergs" 1 . Figure 3 represents the evolution of the model shock in the form of a sequence of vertically shifted density plots in logarithmic scale. All plots are equidistant in time with At — 0.15 yr. The left boundary corresponds to the outer surface of the supernova shell and the ambient gas enters the grid from the right-hand side. The thin shell is visible as a density spike moving slowly to the right. 3.2. Broad-band emission Figure 4 represents the broad-band emission behavior of the post-shock region. The flux maximum at time t ss 3.2 yr corresponds to the moment of thin shell formation. Afterwards the flux slowly decreases at first and then increases again reaching the second maximum at time t w 4.8 yr which marks the moment of main shock reflection from the shell. Subsequent evolution is completely dominated by oscillations of the main shock in

Plewa: Fast Radiative Shocks

0

0.2

0.4

0.6

0.8

399

1

FIGURE 3. Density evolution for initial shock velocity v, = 6000 km s ' (see text for details). 1

FIGURE 4. Variations in time of the global parameters of the model shock for an original shock velocity vs = 6000 km s""1. (right scale) position and velocity of the main shock (dashed and thick dotted lines, respectively); velocity of inflowing gas: thick dotted; (left scale) total, HE, ME, and LE luminosity of the post-shock region (thick, medium, thin and thin dotted lines, respectively). Luminosity scale: 1.52 x 10s erg cm" 2 s" 1 ; distance scale: 1.18 x 1016 cm; velocity scale: 6000 km s" 1 .

front of the shell with time scale corresponding to the fundamental mode of the oscillatory instability. The mean period of the oscillations is roughly equal to Posc = 12 d. The fast variability of the emission is visible after the time of shell formation and reflects the discontinuous character of the flow in the region close to the shell. 3.3. Emitted spectrum Figure 5 shows the spectrum emitted by the shock at time t = 3.14yr, which corresponds to the moment of thin shell formation. The total fiux,Fv, has a power-law shape with a mean spectral index of a = -0.63 with high-energy cut-off around lOkeV. The highenergy end of the spectrum is produced by the gas residing behind the shell. This gas was processed by the shock moving with a velocity nearly equal to initial velocity (v, ss 6000krn s"1) at the beginning of the evolution, and the temperature of this gas

400

, , , , • ? , , ,

Plewa: Fast Radiative Shocks

9 ,

r

4 -2

-2

l

'

r 1

0

2

I ' ' ' 11 L'

-2

0

2

' ' I ' ' ' I " ' ' 1

4 -2

5. The spectrum of the 6000 km s" 1 shock at time t = 3.14 yr. Left panel: total emission; middle panel: net transmitted continuum; right panel: solid line: sum of the reflected continuum and half of the ionizing energy flux; dashed: reflected component only. The energy flux vFv (erg cm" 2 s"1) is plotted on a logarithmic scale. The bottom and top scales are given as the logarithm of energy in Rydberg and keV, respectively. FIGURE

corresponds to the cut-off energy. Therefore, the spectrum of a shock with an initial velocity equal to the initial velocity of a supernova blast wave (vs ss 2 x 10 4 km s" 1 ) would have an energy cut-off at around 100 keV since the post-shock temperature scales as the square of the shock velocity. It has to be noted that our model shock spectrum very closely resembles the thermal Comptonization component used by Zdziarski, Zycki k, Krolik (1993a) in calculations of model X-ray AGN spectra as well as high energy spectrum of NGC 4151 observed by Ginga. and OSSE Zdziarki, Lightman k MaciolekNiedzwiecki (1993b). The emitted shock spectra were used as input spectra to CLOUDY code. Photoionization models were calculated assuming plane-parallel geometry with a density distribution within the shell obtained from hydrodynamical simulations. The results are presented in Figure 6 with values typical for AGNs (Kwan k Krolik 1981) marked on the vertical axes by dots. In general, the agreement between the observations and our theoretical predictions is very good. The predicted column densities are of the order 10 2 3 cm~ 2 and the maximum gas density is roughly equal to 1011 cm" 3 . The model fails to explain the low Lya/H/? ratio but this is a common problem for photoionization models of AGNs. 3.4.

The nature of the time lag

It has been proposed by Terlevich et al. (1994a; hereafter TTRFM) that the time lag between the maximum level of continuum radiation and intensity of emission lines occurs due to the finite time of thin shell formation. Using a simple structure of the shell and an approximate dependence of the continuum variation in time, they obtained qualitatively correct time lags for several emission lines as well as a proper dependence of the time lag on ionization level. In Table 1 we present a comparison between our results and those obtained by TTRFM with the data taken from their Table 1. In our study we used the shell structure and the shape of ionizing flux obtained from detailed hydrodynamical simulation and we used them as input data to the CLOUDY code. The agreement between the two independent studies is good and confirms the findings of TTRFM. We must stress that the present

Plewa: Fast Radiative Shocks N(H)

Hel)l5876/Hp

-"

I''

1

1'

f \ i ll

t

V

4

6

0 *

V 1111

2

NVM240/Lya

\i

H

N -\ 0

C IV JU549 / Lya

MgIU2798/Hp

1 " 1

i

Ha/

Lya / HP

i|ii

r - ' ,_

.s

401

-

1 1 1 , 1 1111

8

10

10

12

12

FIGURE 6. Temporal variation of selected line ratios and parameters derived from photoionization models. Line ratios typical for AGNs are marked by dots on the vertical axes. All values are shown on a logarithmic scale; time is given in years on the horizontal axis.

Time lag [d]

Line H/J Mgll CIII] CIV NV

A4861 A2798 A1909 A1549 A1240

Observed 20 34-72 26-32 8-16 4

TTRFM 30-40 40-50 20-30 10-15 2-10

This paper 17.9 40.5 25.2 3.3 5.1

TABLE 1. Time lags for selected emission lines

study is entirely free from somewhat arbitrary assumptions concerning shell structure and the temporal behavior of the ionizing flux.

4. S u m m a r y The evolution of fast, radiative shocks in high dense media was studied using a highresolution hydrodynamical method. The spectra of the shock have a power-law shape Fv ~ v~a with mean spectral index a ~ 0.6 — 1.0 in the energy range 0.01 — 100 keV. The spectrum has a high-energy cut-off around E « 100keV. This cut-off could be regarded as a natural limit imposed by the maximum velocities attainable by the supernova

402

Plewa: Fast Radiative Shocks 4

1

shocks (t, « 2 x 10 km s" ). The time lag between maximum continuum emission and maximum line emission has a physical nature corresponding to the finite formation time of the thin shell where the lines are produced. Using shock spectra as input data to photoionization models we obtained the line ratios closely resembling those of AGNs. I thank Michal Rozyczka for his steady support and many valuable discussions, and Gary Ferland for making his code available to me. This work was supported by the Committee of Scientific Research through the grant 2-1213-91-01.

REFERENCES R. A. & IMAMURA, J. N. 1982 Astrophys. J. 261, 543. COLELLA, P. & WOODWARD, P. R. 1984 J. Comput. Phys. 59, 264. DORFI, E. A. & DRURY, L. O'C. 1987 J. Comput. Phys. 69, 175. FERLAND, G. J. 1993 HAZY a brief introduction to CLOUDY. University of Kentucky Department of Physics and Astronomy Internal Report. FESEN, R. A. k BECKER, R. H. 1990 Astrophys. J. 351, 437. FlLlPPENKO, A. V. 1992 In Physics of Active Galactic Nuclei (ed. W. J. Duschl & S. J. Wagner). Springer, Berlin. GAETZ, T. J., EDGAR, R. J. & CHEVALIER, R. A. 1988 Astrophys. J. 329, 927. HECKMAN, T. 1991 In STSci Symp. Massive Stars in Starbursts (ed. C. Leitherer, N. R. Walborn, T. M. Heckman fc C. A. Norman), p. 289. Cambridge Univ. Press, Cambridge. HUMMER, D. G. 1988 Astrophys. J. 327, 477. INNES, D. E., GIDDINGS, J. R. & FALLE, S. A. E. G. 1987 Mon. Not. R. Astron. Soc. 226, 67. KWAN, J. &; KROLIK, J. H. 1981 Astrophys. J. 250, 478.

CHEVALIER,

LEIBUNDGUT, B., KIRSHNER, R. P., PINTO, P. A., RUPEN, M. P., SMITH, R. C , GUNN, J.

E., AND SCHNEIDER, D. P. 1991 Astrophys. J. 372, 531.

R., GRONENSCHILD, E. H. B. M. & VAN DEN OORD, G. H. J. 1985 J Astron. Astrophys. Suppl. 62, 197. MEWE, R., LEMEN, J. R., AND VAN DEN OORD, G. H. J. 1986 Astron. Astrophys. Suppl. 65, 511. MEWE, R. 1992 (private communication). PLEWA, T. & ROZYCZKA, M. 1992 Ada. Astron. 42, 295. PLEWA, T. 1993 Acta. Astron. 43, 235. PLEWA, T. 1994 Mon Not. R. Astron. Soc. Submitted. RAYMOND, J., Cox, D. P. & SMITH, B. W. 1976 Astrophys. J. 204, 290. STATHAKIS, R. A. &; SADLER, E. M. 1991 Mon. Not. R. Astron. Soc. 250, 786. STERN, R., WANG, E. & BOWYER, S. 1978 Astrophys. J. Suppl. 37, 195. TERLEVICH, R., TENORIO-TAGLE, G., FRANCO, J. & MELNICK, J. 1992 Mon. Not. R. Astron. Soc. 255, 713. TERLEVICH, R., TENORIO-TAGLE, G., ROZYCZKA, M., FRANCO, J. & MELNICK, J. 1994a Mon. Not. R. Astron. Soc. In press.

MEWE,

TERLEVICH, R., TENORIO-TAGLE, G., CID-FERNANDES, R., FRANCO, J. & ROZYCZKA, M.

1994b In preparation. TURATTO, M., CAPPELLARO, E., DANZIGER, I. J., BENETTI, S., GOUIFFES, C. & DELLA VALLE, M. 1993 Mon. Not. R. Astron. Soc. 262, 128.

A. A., ZYCKI, P. T. & KROLIK, J. H. 1993a , Astrophys. J. Lett, 414, L81. A. A., LIGHTMAN, A. P. & MACIOLEK-NIEDZWIECKI, A. 1993b, Astrophys.J. Lett. 414, L93.

ZDZIARSKI, ZDZIARSKI,

Study of the Stellar Populations in AGN By M. SEROTE-ROOS 1 ' 2 , C. BOISSON 1 AND M. JOLY 1 2

, Observatoire de Meudon, 92195 Meudon, France Centro de Astrofisica da Universidade do Porto, Rua do Campo Alegre 823, 4100 Porto, Portugal

We present here preliminary results of a study of the stellar population in AGN. Our aim is to quantify the stellar population within the nuclear regions by means of spectroscopic observations and to determine whether the central activity influences the stellar population or vice versa. The results will have general relevance to understanding the evolution of galaxies and the energy generation within the nucleus.

1. Introduction We have observed 30 galaxies, of different levels of activity, using long-slit spectroscopy at the CFHT, in the range 5000-10000 A (including Mgl, Nal, TiO, CN and CallT). In order to determine the composite stellar population of a galaxy, it is necessary to obtain spectroscopy of several different wavelength regions, including a number of different absorption lines. Without such information it is not possible to disentangle the effects of abundance variation, luminosity class and stellar type. We are able to detect radial gradients in the stellar distribution (if any), as well as the possible dilution in the nucleus of the stellar component by a featureless component. We shall then establish at what wavelengths and to what degree the stellar population is responsible for the observed activity. 2. Results and Conclusions Preliminary results for NGC 3516 (type 1 Seyfert, SBO/a), Mkn 620 (type 2 Seyfert, S(B)a) and NGC 3379 (non-active galaxy, E0) are presented. For the two active galaxies, EW(CaT + TiO) decreases towards the nucleus. This can be interpreted as a dilution of the nuclear stellar component by an additional spectral component. Under this hypothesis, the fraction of stellar light has been estimated for the two nuclei. In Figures 1 (a) and (b), we plot the nucleus and the bulge of NGC 3516 and Mkn 620, on an arbitrary scale. The bottom spectrum is the difference spectrum after scaling. It appears that the residual continuum is featureless. The existence of dilution by a "nonthermal" continuum could therefore be inferred. In Figure 2, we plot the nucleus and bulge of NGC 3379, where the nuclear stellar population obviously differs from the bulge one. We observe a stellar population gradient. This shows that a stellar population gradient may still exist in the AGN, but is somewhat hidden by the emission line spectrum. A more detailed analysis including a maximum of absorption lines is unavoidably necessary to build a full population synthesis that could answer the problem of the origin of the featureless component. This work is in progress. M.S-R is supported by Junta Nacional de Investigagao Cientifica e Tecnologica under grant number BD/2093/92-RM. 403

404

Serote-Roos et a/.: Stellar Populations in AGN

5000J0

TXXDOJO 7500X) Wavelength ( A ) FIGURE 1. Nucleus and bulge of (a) NGC 3516 (the bottom spectrum is the difference between these two spectra), and (b) Mkn 620.

500OO

75000 Wovelength (A) FIGURE

2. Nucleus and bulge of NGC 3779.

txx»D

Bidimensional Spectroscopy of Seyfert Galaxies: Offset BLR in NGC 3227 By S. ARRIBAS AND E. MEDIAVILLA Instituto de Astron'sica de Canarias, E-38200 La Laguna, Tenerife, Spain. Recently, we have found evidence for an offset active nucleus in the Seyfert 1 galaxy NGC 3227 (Mediavilla & Arribas 1993). In fact, it was found that the BLR is offset with respect to the kinematical centre derived from the ionized gas. This kinematical centre has a heliocentric radial velocity which is in good agreement with previous CO and HI determinations for the systemic velocity of the galaxy, suggesting that the BLR is not at the galactic mass centroid. In addition, some kinematical distortions are observed around the BLR which are probably related to the activity. Here we will give some additional details on thefitof the observational data to a simple model, and the proper location of the kinematical centre.

1. Introduction: bidimensional spectroscopy with optical fibres Bidimensional spectroscopy with optical fibres is a new technique which allows twodimensional mapping of spectral features. Using this technique we are carrying out a programme to study the circumnuclear environment of Seyfert galaxies. To this end, we have developed several optical fibre systems for their use with the telescopes at the Observatorio del Roque de los Muchachos, on La Palma. In particular, the HEXAFLEX system (Arribas, Mediavilla & Rasilla 1991; hereafter Paper I) was conceived for the Nasmyth focus of the 4.2-m, William Herschel Telescope (WHT). The 2d-ISlS works at the auxiliary focus of the WHT, and HEXAFLEX-II with the Nordic Optical Telescope (NOT). From the astronomical point of view we have analyzed three galaxies so far. NGC 4151 (Mediavilla, Arribas & Rasilla, 1992; hereafter Paper II), NGC 5728 (Arribas & Mediavilla 1993), and NGC 3227 (Mediavilla & Arribas 1993). In the context of this meeting the results obtained for NGC 3227 may be of interest interest and these are the main subject of this communication. However, we shall first briefly comment on the basis of the experimental technique. The aim of the bidimensional spectroscopy with optical fibres is to perform simultaneous spectroscopy of the different regions of an extended object. The kernel of this technique consists in using specially designed optical fibre bundles. One end of the bundle forms a two-dimensional array offibres,and is connected to the telescope focal plane. The other end of the bundle has the fibres aligned at the entrance of the spectrograph. Thus, when the telescope points to an extended object each fibre transmits the light coming from a region of the object, and all the spectra (one per region in the object) are simultaneously recorded. This implies important advantages over other techniques such as long-slit, Fabry-Perot, or filter images, as all the information (spatial and spectral) is simultaneously recorded. For details, see papers I and II. 2. The case of N G C 3227 NGC 3227 is a Seyfert 1 galaxy with two well-defined arms, one of which ends in the elliptical companion NGC 3226. It has been classified as SABa pec by de Vaucouleurs et al. (1991) and as Sb(s) (tides) by Sandage and Tammann (1987). We observed this galaxy in May 1989 using the HEXAFLEX system (see Paper I) connected to the Nasmyth focus of the 4.2-m (WHT). The spectral range was 6300-7000 A, the spectral 405

Arribas k Mediavilla: Offset BLR in NGC 3227

406

O

o-

-5 0 La (arcsec) FIGURE

-5

1. Spatial distribution of thefibresat the telescope focal plane of the 4.2-m William Herschel Telescope.

resolution about 3 A, the spatial sampling (mean distance between two adjacent fibres) 1.2 arcsec, and the spatial element (fibre core diameter) 0.9 arcsec. The bundle had 61 fibres forming a hexagonal array covering about 13 arcsec in diameter (see Figure 1). With this configuration we took three consecutive exposures of 1800 sec each. The indvidual radial velocities were determined from simultaneous Gaussian fitting to the Ha, [Nil] AA6548.6584, and the [SII] AA6716,6731 emission lines, restricting wavelength differences, the intensity ratio for the N lines, and considering the same width for all the lines. The velocity uncertainty is estimated to be about 15 km s" 1 for most of the spectra. The velocity field is presented in Figure 2a. With the exception of the NW region this velocity field looks like a rotational pattern around the black dot (kinematical centre), which is situated in the region of largest velocity variation, which in turn is between the two poles. Its amplitude is about 200 km s" 1 which is typical of a spiral galaxy. However, some distortions are also apparent, probably due to movements radially and perpendicularly to the disk. To illustrate this, we have adjusted a simple model to the data. This model assumes a plane disk and axial symmetry; allows rotation, radially as well as perpendicularly to the disk movements, and can be formally represented by the equation:

v{p, ) =

Z(R) cos i

cos(<£ - 6(R)) sin i(R) (1)

where Vsya is the systemic velocity, Q,(R) the strength of the velocity field projection

Arribas k Mediavilla: Offset BLR in NGC 3227

:

5O V 01

o

0-

3

'• U / ^ ^ 1

: :

^

6

407

\

J /k

)

•

\/ 1 ^

-5'• 1 • • • • 1 • « • • 1 • ' '• i • • ' • 1 • • •

5

0 Aa (arcsec)

5

0 Aa (arcsec)

-5

FIGURE 2. (a) Velocity field derived from the Ha, [Nil] and [SII] emission lines. The dot indicates the position of the kinematical centre and the star the position of the BLR. (b) Model adjusted to the data (see text).

on to the galactic plane, and Z(R) the corresponding perpendicular component, p and <j> are projected polar coordinates. In Figure 2b we represent the model adjusted to the data. The r.m.s. of the differences with the data is in the range 15-25 kms" 1 , depending on the distance to the kinematical centre. The heliocentric radial velocity of the kinematical centre is 1136 kms" 1 . This value is in good agreement with previous systemic velocity determinations derived from CO and HI measurements (1156 km s" 1 , Braine k Combes 1992; 1146 km s" 1 , Mirabel 1982; 1138 km s" 1 , Huchtmeier k Bohnenstegel 1975; 1106 km s" 1 , Heckman, Balick k Sullivan 1978). Note that the average of all these values is 1136.5 ± 20 km s" 1 , or 1147 ± 9 km s- 1 neglecting the Heckman, Balick k Sullivan value. In any case, there is good agreement with our heliocentric radial velocity for the kinematical centre. This suggests that the kinematical centre derived from the ionized gas represents the galactic mass centroid. However, when we situated the position of the BLR over the velocity field (this can be done unequivocally, thanks to the observational technique used), we found it to be shifted by about 3 arcsec NW from the kinematical centre (see Figure 2a). The BLR is near the position where the largest deviations from the simple kinematical model are found. Thus, the BLR (and its enviroment) is blueshifted with respect to the kinematical centre. However, the BLR is coincident with the continuum maximum, the line intensity maxima of the Ha, [Nil], and [SII] emission lines, and the radio peak according to the work by Argyle k Eldridge (1990). Thus, NGC 3227 reminds the cases of NGC 2110 (Wilson et al. 1985) and NGC 5548 (Wilson et al. 1989), where displacements between the kinematical centre and the continuum maxima were found. In the case of NGC 3227 it is the BLR itself (in addition

408

Arribas k Mediavilla: Offset BLR in NGC 3227

1300 -

-2

0

2

6

distance (arcsec) FIGURE 3. Velocity curves crossing the kinematical centre along PA ~ 140 deg (continuous line), and crossing the BLR along PA ~ 50 deg (dashed line). The circle and the star mark the positions of the kinematical centre and the BLR, respectively. The straight line is at the systemic velocity of the galaxy derived from HI and CO (see text).

to the continuum maximum and the line emission maxima) which is found to be shifted with respect to the kinematical centre. It any event, even if doubts about the exact location of the kinematical centre of NGC 3227 exist, the BLR can hardly be associated with the kinematical centre. To illustrate this we have represented in Figure 3 the changes in heliocentric velocity along the best suited major kinematical axis (determined by using the kinematical model) in both cases: a) across the kinematical centre determined by us (PA ~ 140 deg), and b) across the BLR (PA ~ 50 deg). Note that the situation of the BLR is very asymmetric along the velocity curve which crosses it. On the other hand the velocity curve corresponding to the kinematical centre is symmetric, shows greater variations in amplitude and is very well centred with respect to the systemic velocity of the galaxy. In 1992 we observed the region around the BLR in more detail with a system similar to HEXAFLEX for the ISIS spectrograph of the 4.2-m (WHT). The signal-to-noise ratio was lower (which implies a smaller spatial coverage), but the spectral resolution was higher (about 2 A) and we observed the E/3 and [OIII] AA4959,5007 emission lines (in addition to the Ha, [Nil] and [SII] lines). We found good agreement between the velocity field derived from these data and the one obtained from the 1989 ones. The relevant point with respect to these higher spectral resolution spectra is that, in some cases, we found double-peaked emission lines or clear traces of them. This suggested Gaussian decomposition of the [OIII] profiles into two components, which could be discriminated in terms of their mean parameters. Thus, one component was narrower (FWHM ~ 3 ± 1.7 A) and was generally redshifted, having a typical galactic behaviour in terms of line width and central wavelegth. The other was broader (FWHM ~ 9 ± 3.5 A) and was generally blueshifted. This component may indicate outflow and

Arribas k Mediavilla: Offset BLR in NGC 3227

409

its spatial coincidence with the vicinity of the BLR suggests that it is connected to the activity. 3. Conclusions and final comments The main conclusions of this communication are as follows: 1) The distortion appearing in the velocity field of NGC 3227 may be related to the activity. 2) Using a simple kinematical model, we have found that the BLR can hardly be associated with the kinematical centre, which is likely to be situated about 3 arcsec SE from it. In addition, the coincidence between the heliocentric radial velocity of this kinematical centre and the systemic velocity determinations derived from CO and HI measurements suggests that the kinematical centre derived from the ionized gas represents the galactic mass centroid. It is worth commenting that we do not find a secondary emission at the kinematical centre position in the optical continuum map, though our data do not offer the best possibilities for finding it. Thus, our future work on this galaxy will be focussed on the detection of this secondary emission (if it indeed exists), and on the determination of the velocity field from a stellar feature (if possible, with the current observational possilities).

REFERENCES ARGYLE, R.W. & ELDRIDGE, P. 1990 M. N. R. A. S. 243, 504. ARRIBAS, S. & MEDIAVILLA, E. 1993 Ap. J. 410, 552. ARRIBAS, S., MEDIAVILLA, E. & RASILLA, J. L. 1991 Ap. J. 369, 260. (Paper I) DE VAUCOULEURS, G. ET AL. 1991 Third reference Catalogue of Bright Galaxies. Springer. BRAINE, J. k COMBES, F. 1992 A. A. 264, 433. HECKMAN, T. M., BALICK, B. & SULLIVAN, W. T. 1978 Ap. J. 224, 745.

HUCHTMEIER, W. K. & BOHNENSTENGEL, H.-D. 1975 A. A. 44, 479. MEDIAVILLA, E., ARRIBAS, S. & RASILLA, J. L. 1992 Ap.J. 396, 517. (Paper II) MEDIAVILLA, E. & ARRIBAS, S. 1993 Nature 365, 420. MIRABEL, I. F. 1982 Ap. J. 260, 75.

A. & TAMMANN, G. A. 1981 A Revised Shampley-Ames Catalog of Bright Galaxies. Carnegie Inst.

SANDAGE,

WILSON, A. S. & BALDWIN, J. A. 1985 Ap. J. 289, 124. WILSON, A. S., Wu, X., HECKMAN, T. M., BALDWIN, J. A., & BALICK, B. 1989 Ap. J. 339,

729.

ROSAT Detection of the Most Rapidly Varying Seyert Galaxy By Th. BOLLER AND J. TRUMPER Max-Planck-Institut fur Extraterrestrische Physik, D-85748 Garching bei Munchen, Federal Republic of Germany X-ray variability in the 0.1 — 2.4 keV ROSAT energy band with a doubling time scale of 800 s and a factor of 4 within a few hours has been detected in a 20 ksec pointing on the AGN IRAS 13224-3809. The optical spectrum indicates that IRAS 13224-3809 is a narrow-line Seyfert 1 galaxy with strong permitted Fell emission, a member of the unusual I Zw 1 class objects. IRAS 13224-3809 appears to be the most rapidly varying AGN known so far. This is the first time that variability on a time scale smaller than 1000 s is reported at such high X-ray luminosity [1(0.1 - 2.4 keV ) = 3 • 1044 erg s"1] in Seyfert galaxies. It is also the first reported X-ray variability in I Zw 1 class objects. The At = 800 s variation indicates that the X-rays come from a compact region of about 17 light minutes in size. Our results from the X-ray spectral analysis favour a scenario in which a hard X-ray source irradiates the accretion disk which reemits at soft X-ray energies. The absence of broad HI wings can be explained only if a part of the BLR, far from the centre, is observed and the bulk of the region, which emits the wings, is hidden. We want to draw attention to the fact that rapid X-ray variability could also be connected with the absence of broad HI lines in IRAS 13224-3809.

1. Introduction Variability studies are unique tools for studying the physical conditions in astrophysical objects. The time scales on which variability occurs in extragalactic sources give constraints on the sizes of the emitting regions and on the transfer function of the surrounding material, and they can be used to estimate the mass of the central black hole. Usually relativistic beaming is assumed to explain the fast variability in BL Lac objects, first discussed by Blandford and Rees (1978) and Rees (1984). Instabilities in the accretion disk near the central source are supposed to cause the X-ray variability in Seyfert galaxies or QSOs. Most Seyfert galaxies, quasars and BL Lacs show X-ray variability on time scales between 104 s and 106 s. Only a few AGN have reported variations on time scales less than about 1000 s. The BL Lac object H0323+022 (Feigelson et al. 1986), perhaps the most rapidly variable extragalactic source known, varies by a factor of 3 in about 30 s. The shortest time scales found in Seyfert galaxies are about 1000 s in NGC 4051 and MCG-6-30-15 (Matsuoka et al. 1990). In this contribution we report on the X-ray timing and spectral properties of IRAS 13224-3809 and the results from the optical follow-up spectroscopy.

2. X-ray observation The X-ray observation was performed between August 10 and August 12, 1992, using the PSPC detector (Pfeffermann et al. 1987) on the ROSAT satellite (Trumper 1983) in a 20 000 s pointing on IRAS 13224-3809. The mean count rate in the ROSAT 0.1 - 2.4 keV energy band is 0.31 counts s" 1 . The resulting X-ray luminosity is 3 • 1044 erg s" 1 assuming a Galactic column density of 4.85 • 1020 cm"2 and a photon index of 4.4. Figure 1 shows the X-ray (0.1 - 2.4 keV) light curve of IRAS 13224-3809. Coherent variations on time scales of about 800 s by a factor of 2 between epochs at 18 000 and at 128 000 410

Boiler & Triimper: Most Rapidly Varying Seyfert Galaxy

411

0.7 0.6 O.S 0.4 2000.0

40OO.0

6000.0

8000.0

126000.0 128000.0 130000.0 132O00.0 134000.0

1O0OO. 0

O.S 0.3 0.4 0.2

CD

0.3

•(->

0.2

0.1 10000.0

12000.0

HOOD. 0

16000.0

18000.0

20000.0

170000.0 172O0O.O 17400O.O 176000.0 178000.0 1B00OO.O

3 0.3

0.3

0.2

0.2

0.1

0.1

36000. 0

38000.0

40000.0

42000.0

44000.0

46000.0

180000.0 182000.0 184000.0 186000.0 188000.0 190000.0

0.7 0.3

0.6

0.2

O.S

0.1

0.4 0.3

90000.0

92000.0

94000.0

96000.0

98000.0

100000.0

236000.0 238000.0 240000.0 242000.0 244000.0

time [ s ] FIGURE

1. X-ray light curve of IRAS 13224-3809 (count rate versus time in seconds). The doubling time scale is only 800 s.

s after the beginning of the observation have been detected. Amplitude variations of a factor of about 4 can be seen on time scales between 8 and 16 hours.

3. Optical observations A CCD image of IRAS 13224-3809 and two long-slit spectra were taken with the La Silla 3.6-m telescope using the EFOSC spectrograph on 1992 December 24 (see Figure 3 of Boiler et al. 1993). Numerous emission lines are seen in the galaxy spectrum. The Fell lines from the multiplets around 4570 and 5270 A, Hp and [OIII] A5007, relevant to the present study, were measured. IRAS 13224-3809 appears to be one of the most intense Fell emitters with FeII(4570)/HJg = 2.4. The intensity of the narrow H.p emission is only slightly broader than the forbidden line [OIII] A5007. No indications for broad wings of Hp are apparent.

4. Discussion The At ~ 800 s X-ray variations with AL = 1.5 • 1044 erg s" 1 indicate that the X-rays come from a compact region of about R ~ At c w 10~ 5 pc in size. The small size of the emitting region can be taken as the strongest proof for the presence of an active nucleus. Starburst models for AGN have difficulties in explaining the X-ray variability with a doubling time scale of about 800 sec at a (0.1 — 2.4 keV) X-ray luminosity of a few l O ^ e r g s " 1 . The observed steep X-ray spectrum was compared to power-law, blackbody and Bremsstrahlung models. Only the Bremsstrahlung model gives an acceptable fit, but it must be rejected on the basis of the variability argument. Emission from a standard accretion

412

Boiler Sz. Triimper: Most Rapidly Varying Seyfert Galaxy

disk model fits the data. Again, however, from the observed rapid variability we are forced to rule out this model. A scenario in which a hard X-ray source irradiates the accretion disk, which reemits at soft X-ray energies can explain both the steep X-ray spectrum and the variability. The optical spectrum of IRAS 13224-3809 indicates that the source is a strong optical Fell emitter with FeII(4570)/H;g = 2.4. The Fell problem is an unsolved problem in AGN studies. From observations it is known that the total strength of the Fell blends can equal the H a intensity (Collin-Souffrin et al. 1986) while the strength calculated from photoionization models is only about 0.5 (Collin-Souffrin et al. 1986 and references therein). Collin-Souffrin and collaborators (1988) have shown that the observed FeII/H a ratio may result from hard X-ray reprocessing in the outer region of the accretion disk, where physical conditions are different from those of normal broad-line region clouds, i.e. the gas is in hydrostatic equilibrium and the density is higher. Such conditions selectively increase Fell emission with respect to Balmer line emission. It is of some interest that our spectral analysis of the PSPC data leads to a similar scenario, namely that IRAS 13224-3809 may host a hard X-ray source which irradiates the disk. The assumptions of Collin-Souffrin et al. (1988) solve the Fell problem in the majority of AGN. The Fell problem is not, however, completely solved in the case of intense Fell emitters with Fell A4570/H> > 0.8 (Joly 1987, her Figure 4). The combination of unusual observational parameters, the rapid X-ray variability, the steep X-ray spectrum, and the extreme Fell /Up ratio has brought us to search for possible links between these parameters in IRAS 13224-3809. We would add another speculative explanation for the existence of narrow HI lines and pronounced optical Fell emission in IRAS 13224-3809. We have indications that in IRAS 13224-3809 superEddington conditions may be reached, which could lead to fast variability and may change the conditions in the broad-line region. We obtain 0.1 < L/Lsdd < 10. From the standard disk model fit we have L/Lsdd — 0.7, which is close to the Eddington limit, too. Under super-Eddington conditions the broad-line region may be affected by outflow of gas and dust from the central source depressing the broad Balmer line emission. As mentioned above Fell could arise mainly from the outer regions of the disk and would therefore not be influenced strongly by the outflow. The detailed discussion can be found at Boiler et al. (1993).

REFERENCES BLANDFORD, R. D., REES, M. 1978 In Pittsburgh Conference on BL Lac Objects (ed. A. M.

Wolfe), p. 328. Univ. Pittsburg Press, Pittsburgh. BOLLER, TH., TRUMPER, J., MOLENDI, S., FINK, H., SCHAEIDT, S., CAULET, A. & DEN-

NEFELD, M. 1993 Astron. & Astrophys. 279, 53. S., DUMONT, S., JOLY, M. &; PEQUIGNOT, D. 1986 Astron. & Astrophys. 166, 27. COLLIN-SOUFFRIN, S., HAMEURY, J.-M. & JOLY, M. 1988 Astron. Sc Astrophys. 205, 19. FEIGELSON, E. D. ET AL. 1986 Astrophys. J. 302, 337. JOLY, M. 1987 Astron.
PFEFFERMANN, E., BRIEL, U. G., HIPPMANN, H., KETTENRING, G., METZNER, G., PREDEHL, P., REGER, G. & STEPHAN, K.-H. 1987 MPE print 81.

REES, M. J. 1984 Ann. Rev. Astron. & Astrophys. 22, 471. J. 1983 Adv. Space Res. 4, 241.

TRUMPER,

QSO Evolution: a Link with Starbursts? By B. J. BOYLE Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge, CB3 OHA, UK I review recent results on the cosmological evolution of QSOs identified at optical, X-ray and radio frequencies. In all these regimes, it is now clear that the redshift range 2 £ z < 3 corresponds to the epoch of maximum QSO activity. I demonstrate that QSO models invoking supermassive black holes or the starburst cores of young elliptical galaxies are equally successful at reproducing the observed space densities of even the most luminous QSOs in the Universe at these redshifts. In addition, both models can also account for the strong decline in QSO luminosity, L(z) oc observed in all regimes at lower redshifts (z~2). In the infrared and X-ray (1 + z)30±0A, passbands, recent results suggest that starburst galaxies may also exhibit a remarkably similar rate of evolution to QSOs, L(z) oc (1 + z)3±l.

1. Introduction In the last few years there has been a dramatic increase in the numbers of QSOs identified in complete spectroscopic surveys. In particular, a significant number of high redshift QSOs have now been identified, including more than 40 QSOs at z > 4. The mere existence of these high redshift QSOs, coupled with their inferred high metal abundances (Hamann & Ferland 1993; Ferland, this conference), indicates the presence of a significant number of massive, evolved systems only ~ 1 Gyr after the Big Bang. Conventional wisdom has it that such QSOs pose a significant challenge for many theoretical models of QSO formation, including both the conventional supermassive black hole model (see Rees 1984 for a review) and the less conventional starburst model, as revived by Terlevich k Melnick (1985). For example, in a cold dark matter (CDM) dominated universe (Davis et al. 1985) galaxy formation occurs at relatively low redshifts and it has been stated that it is difficult to form super-massive (~ 109A/©) black holes sufficiently quickly at high redshifts to account for the observed numbers of z > 4 QSOs. Alternatively, it has been argued that the starburst model cannot account for the extremely high bolometric luminosities (1048ergs~1f) inferred for such QSOs. However many of these claims have been made without detailed analysis of these models and the numbers of luminous high redshift QSOs that they actually predict. In this review, I shall therefore discuss the latest results on the observed space density and cosmological evolution of QSOs and compare these observations to the space densities predicted by both the black hole and starburst models. 2. QSO evolution 2.1. Optical

With ~ 2000 optically-selected QSOs now identified in complete spectroscopic surveys (see Figure 1), a consistent picture of QSO evolution in the optical passband is beginning to emerge. At z < 2, the evolution of the QSO luminosity function (LF) is dominated by strong luminosity evolution (Marshall et al. 1984; Koo & Kron 1988; Boyle, Shanks & Peterson 1988; Boyle et al. 1991). This evolution can be represented by a shift of the LF towards higher luminosities at higher redshifts parameterized by a power-law evolution t Unless otherwise stated, cosmological parameters Ho = 50 kms~1Mpc~1, qo = 0.5 and A = 0 are assumed throughout this paper. 413

414

Boyle: QSO Evolution

Luminosity (erg s"1)

Luminosity (erg s~') I I ' " 1 " . 1985

I , , ,|

1

I I'

I

|

I

I

II

|

I

i;

i

|

I.

1993 -_ 12

:''

o o

4 -

N

I

1047 _J I , | i

o: o: ":'

o l"

/

3 _-

4-

XI

W T3

1 /#•

-

>

o "I

_ 10

a

• •

J

•

?/ % • /*.•*

(0

•

a

«?«.'

1 --

N 3

3

t ill -

A

, , , 1i

i

(k'

•4- 9 -

8 7 ;- _

111

-20 -22 -24 -26 -28 -30 Absolute Magnitude (MB)

5

_5 1

I I I I I I I I I I I I |

I I I

-20 -22 -24 -26 -28 -30 Absolute Magnitude (MB)

1. Absolute blue magnitude (MB) V. redshift diagram for optically-selected QSOs identified in complete spectroscopic surveys a) prior to 1986 and b) to date (Jan 1994).

FIGURE

of the form Lopt(z) oc (1 + z)*°p« (see Figure 2). For q0 = 0.5 and a mean optical spectral index a opt = 0.5, the rate of optical evolution is kopt = 3.45 ± 0.1 (Boyle et al. 1991), although the derived rate could be significantly less if a sizeable dispersion exists in the optical spectral index (Giallongo & Vagnetti 1992; Francis 1993). Evidence for some luminosity dependence in this evolution at the very highest QSO luminosities (i.e. the most luminous QSOs evolve less strongly with redshift) has also been found by Hewett, Foltz k Chaffee (1993) and Goldschmidt et al. (1992). At z 12 there is general agreement that the strong evolution witnessed at lower redshifts slows down, although the precise form and nature of this "slow-down" is still the matter of some debate. Boyle et al. (1991) find that the observed number of z>2 QSOs in a compilation of surveys is consistent with an unevolving QSO population in the interval 1.9 < z < 2.9 i.e the comoving space density of QSOs at any given luminosity remains constant over this redshift range. Based on the QSOs identified in the Large Bright Quasar Survey, Hewett et al. (1993) fit an evolution model which includes a break to weaker luminosity evolution, Lopt(z) a (1 + z)1 5 , at z > 1.6, but this evolution continues on until z ~ 3. Warren, Hewett & Osmer (1994) also favour a more gradual slow-down, fitting a weaker exponential luminosity evolution model Lopt(z) oc e*T at z > 2 which peaks at z ~ 3.3. At the very highest redshifts (z > 3), observations indicate that QSO evolution is strongly dependent on luminosity. The most luminous QSOs (MB ~ -28) exhibit little

Boyle: QSO Evolution 0.001 p-r-T—T-

T '^

r

415

I'''I'''I'' I

0.0001 lO" 6 10- 6 o a, 2 lO" 8 10-"

H, o 0.30
lo-io

• 1.25
lO"11

-20

, I , . , I,

-22

-24

I

-28

-26

-30

FIGURE 2. The QSO optical luminosity function $OP\(MB)

plotted in four discrete redshift intervals over the range 0.3 < z < 2.9 (Boyle et al. 1991).

1-5

O 2.0
M

1

-

2.2
o

a •e-

-23

-24 -25 -26 absolute magnitude Mc(1216)

-27

FIGURE 3. The QSO optical luminosity function $op«(A/c) plotted in four discrete redshift intervals over the range 2 < z < 4.5 (Warren et al. 1994). An approximate conversion between Me (1216) derived assuming Ho = 75kms" 1 Mpc~ 1 (plotted in this figure) and MB derived assuming #o = 50kms~1Mpc~1 (used throughout this paper) is given by: A/c(1216) ~ M B + 1.5. Reproduced with permission from Warren et al. (1994).

416

Boyle: QSO Evolution

- '

1 '

i

•

1

*

i

•

1 1994

o

•

o

'. oEMSS _.ROSAT

'

QP

-

^Oo

o -

j

1"

• <

sP -A M

•

42 44 46 1 log L,< (ergs s" )

-

• <

42 log

m44Ho

1

1

.

46 1 (ergs s" )

FIGURE 4. X-ray (0.5-2 keV) luminosity vs. redshift diagram for X-ray selected QSOs identified in complete spectroscopic surveys a) prior to 1990 (MSS: Maccarcaro et al. 1984) and b) to date (Jan 1994). Filled circles represent QSOs identified from a deep ROSAT survey (see Georgantopoulos et al. 1994) and open circles denote the EMSS QSOs (Maccacaro et al. 1991).

cosmological evolution over the redshift range 2 < z < 4, with the space density of luminous QSOs remaining unchanged over the redshift interval 2 < z < 4 (Mitchell et al. 1990; Irwin, McMahon k Hazard 1991). In contrast, less luminous QSOs (MB 26) undergo a significant decline in their numbers between z ~ 3 and z ~ 4 (Schmidt, Schneider & Gunn 1991; Warren et al. 1994). This evolution has been modelled by Warren et al. (1994) as an exponential decline in comoving space density, $Opti of the form $Opt(2) <x ea;p[(3.3 - z)ko], where kp = 3.3; corresponding to a drop in the space density of MB ~ -26 QSOs by a factor 6.5 between z = 3.3 and 2 = 4.0. The QSO LF for z > 2 QSOs derived by Warren et al. (1994) is plotted in Figure 3. 2.2. X-ray The last few years have also seen a significant increase in the number of X-ray selected QSOs identified in complete spectroscopic surveys (Figure 4). Based on the Einstein (0.5-2 keV) Extended Medium Sensitivity Survey (EMSS), Maccacaro et al. (1991) established that the X-ray evolution of QSOs at low redshifts z < 2 was well represented by luminosity evolution, but with a significantly slower evolution rate Lx(z) oc (1 + z^2.56±o.i tha.n derived in the optical. More recently, significantly faster rates of X-ray evolution have been derived from QSO samples identified with the ROSAT (0.5-2 keV) mission, Lx{z) oc (1 + zf ° ( Boyle et al. 1993, 1994a). The X-ray LFs derived from the EMSS and ROSAT QSO samples by Boyle et al. (1994a) are plotted in Figure 5. The reason for the difference between the evolution rates obtained from the EMSS and ROSAT samples is not yet understood, although it is possible that uncertainties in the X-ray spectral index for QSOs could play a role in the observed discrepancy between the

Boyle: QSO Evolution 0.001

fe

417

i i i .

1

I

0.0001 10"5 bO

o

10-6

a, 2

io- 7

g

io-8

o

10"9

qo=O.5 H0=50 o 0.0
1O -io

1.8
.

l

43

,

,

i

I

44 45 log I^ (ergs s"1)

46

FIGURE 5. The QSO X-ray luminosity function $(Lx) plotted in four discrete redshift intervals over the range 0.0 < 2 < 3.0 (Boyle et al. 1994a).

two surveys (Boyle et al. 1993). It has also been suggested by Francheshini et al. (1993) that the EMSS contains a two populations of QSOs; a population with hard X-ray spectra which exhibit weak cosmological evolution and a population with softer X-ray spectra which exhibit the strong LJC(Z) oc (1 + z) 3 evolution. This latter population dominates the fainter ROSAT samples. However, the fast rate of X-ray evolution also derived from a sample of QSOs detected at much harder energies (2-10 keV) by the HEAO Al mission (Marshall 1991) would appear to be inconsistent with this explanation. In common with the optical regime, the strong X-ray evolution also appears to "switch off' at z ~ 2. Tentative evidence for a slow-down in the evolution rate at these redshifts was first found by Delia Ceca et al. (1992), based on the QSOs identified in the EMSS. However, using additional high redshift QSOs detected by ROSAT, Boyle et al. (1994a) have been able to demonstrate that the number of z > 2 QSOs in the ROSAT and EMSS samples is consistent with an unevolving QSO population over the redshift range 1.7 < z<3. Although a few very high redshift X-ray selected QSOs have been found (see e.g. Henry et al. 1994), as yet insufficient QSOs at z > 3 have been identified in complete X-ray surveys to derive meaningful limits on the X-ray evolution of QSOs at these redshifts. 2.3. Radio Unlike the optical and X-ray regimes, complete flux-limited surveys of QSOs at radio frequencies are much more difficult to obtain. This is primarily due to the large number of faint (B > 21) optical counterparts which are found in any radio survey, even those conducted at relatively bright radio-flux limits. Nevertheless, using a sample of flat-spectrum (aR < 0.5) radio-selected (2.7 GHz) QSOs, Dunlop k Peacock (1990) were able to demonstrate that the evolution of radio-selected QSOs at z < 3 was well de-

Boyle: QSO Evolution

418

0.38£z<0.75 :

(a) 0.75§z<2.00 •

^

- <,

2.00Sz<3.00 . n nnsj<4. sn

IV

r o 2 a t 2 o

\\ \\ \\

-

\

1

10'

,,

i

- •

, . .t

10'

FIGURE 6. The QSO radio luminosity function P(LR) plotted in four discrete redshift intervals over the range 0.38 < z < 4.50 (Hook 1994). scribed by luminosity evolution of the form LR(Z) OC dex(1.18z — 0.2822). This evolution law corresponds closely to the familiar L(z) oc (1 + z)3 law at z < 1.8, but peaks at z ~ 2.5, after which the evolution goes into reverse, i.e. decreasing luminosity at higher redshift. Although hampered by the lack of high redshift QSOs with accurate spectroscopic redshifts in their sample, Dunlop & Peacock (1990) were able to demonstrate that some form of luminosity-dependent evolution was occurring at z > 2. Analogous to the optical regime, they found that high-luminosity QSOs (L/{ = 10 27 WHz -1 ) appeared to exhibit a less marked decline in their comoving space density between z = 2 and z = 3 than lower-luminosity QSOs (LR = 1026 WHz" 1 ). Indeed, to within the errors of the fit, the numbers of LR = 1027WHz~1 QSOs were consistent with an unevolving population over the redshift range 2 < z < 3. This general trend has been confirmed with a new sample of high-redshift (z > 3) radio-loud QSOs identified by Hook (1994). Hook finds no evidence for any significant decline between z = 2 and z = 4 in the comoving space density of the most luminous radio-selected QSOs (LR = 1027 2 5 WHz" 1 ). In contrast, the less luminous radio QSOs (LR = 1026-25 WHz" 1 ) appear to exhibit a decline in their space density by a factor 5-10 over the same redshift range (see Figure 6), although optical incompleteness in the sample could still influence this result (McMahon 1991). 2.4. Summary From the previous sections, it is clear that QSO evolution is remarkably similar throughout the optical, X-ray and radio regimes. At low redshifts z£ 1.8, the evolution in each regime is marked by a strong, approximately power-law increase in luminosity with redshift of the form L(z) oc (1 + z)k. In all cases, the derived range of evolution lies in the narrow range 2.6 < k < 3.4 and, given the systematic uncertainties present in the

Boyle: QSO Evolution

419

various analyses (e.g. spectral index effects), a single value of Jb ~ 3 for all regimes is not inconsistent with the data. In each passband, the rate of evolution also appears to decline (or switch off) at z> 1.8, with the redshift interval 2 £ z 13 representing the period of peak of QSO activity in the optical, X-ray and radio. For the most luminous QSOs, this period could extend as far as z ~ 5, with no decline in their space density at high redshifts being observed in either the optical or radio regimes. In contrast, lower-luminosity optical and radio selected QSOs appear to undergo a signifcant drop in space density (by a factor ~ 5) between z = 3 and z = 4. The (1 + z)3 QSO evolution law also appears to apply to starburst galaxies. Saunders et al. (1990) derive an infrared evolution of the form IIR(Z) oc (1 + z) 3 ± 1 for luminous (L > 1011 LQ) starbursts selected at 60 fim from the IRAS survey. The large error on the evolution rate reflects the small redshift range (z < 0.1) over which the evolution was determined. More recently, Boyle et al. (1994b) have also obtained a similar rate of evolution Lx{z) oc (1 + z) 2 6 ± 1 for a sample of narrow emission line, X-ray luminous galaxies (1042 — lO^ergs" 1 ) with z < 0.5. These objects were selected from the EMSS and the deeper Cambridge-Cambridge ROSAT serendipity survey (McMahon et al., in preparation). Although a precise identification of these galaxies requires further observations, detailed spectroscopy of at least one member of the sample reveals it to be a "weak-[OI] LINER", possibly indicative of a young starburst (Filippenko & Terlevich 1992). 3. Interpretation Signifying the epoch of peak QSO activity, the ability to reproduce the space density of QSOs in the redshift range 2 < z < 3 therefore provides the most severe test of any model seeking to explain the physical processes responsible for the QSO phenomenon. Similarly, any model should also be able to reproduce the subsequent evolution of the QSO population. In the supermassive black hole model, Efstathiou & Rees (1988) have shown that a constant comoving space density of even the most luminous QSOs between z = 2 and z = 4 is not inconsistent with the formation of super-massive black holes (~ 109 M©) in a Universe dominated by cold dark matter. More detailed modelling by Haehnelt & Rees (1993) has demonstrated that the QSO luminosity function in the redshift range 2 < z < 3 can be accurately reproduced by a model in which the supermassive black holes responsible for the initial generation of QSOs form quickly, within ~ 108yr, at z*8. During this period, there is sufficient time for the first generation of stars to produce the high metallicities observed in high redshift QSOs (Hamann k. Ferland 1993). In the Haehnelt & Rees (1993) model, QSOs are short lived (~ 108yr) and they all radiate at approximately the same ratio to their Eddington luminosity (L/LEdd ~ 0.01). The subsequent evolution of the QSO LF at z < 2 is therefore driven by the statistical evolution in the mass of the black hole formed for each new generation of QSOs (~ 100 generations by the present epoch). For the observed luminosity evolution this requires that black holes with smaller masses form at later epochs (i.e. lower redshifts). However, in a hierarchical framework such as CDM, the collapsing halos from which these black holes form will have larger masses at later epochs. Haehnelt & Rees (1993) therefore argue that the efficiency of black hole formation must decrease dramatically at lower redshift and obtain a good fit to the observed QSO LF (see Figure 7) for an efficiency which drops off as (1 + z) 5 5 ; approximately equal to the square of the density in the collapsing halo.

420

Boyle: QSO Evolution

......

4

10" 10 -

:

'

'

10-

1

'

'

'

1

5

i

[

^ •

<*\v.

7

r

•

•

•

i

*

0.30
0

10" 6 o

'

•

-

0.70<2
:

1.25
^

2.0<2<2.9

:

r

1

Y\\\

0" 8 10-

\

!

7 - 0?

10- 1 0 10'

\ \

i - 1.025 ;

z - 1.625

r

z - 2.45

>

*

\

•

1

•

-20

-22

-24

.

i

.

i

A

it.

i

•.

:

'.

•

i

. •• i

-28

-26

-

':

'. >

-12

:

\

*

10

•

'.

\

.

-30

FIGURE 7. The predicted optical QSO luminosity function and evolution (lines) based on the supermassive black hole model of Haehnelt & Rees (1993). The observed LF derived by Boyle et al. (1991) is also shown (symbols). Reproduced with permission from Haehnelt & Rees (1993).

10 '

1

10"' !T

Elliptical Galaxy LF (+ cD) with t"* Evolution

' I ' ' ' I ' ' ' I^^ T ' I n ' ' I ' -=

O 0.30
• 0.70<2<1.25 o 1.25
10

-20

-22

-24

-26

-28

-30

FIGURE 8. The predicted optical QSO luminosity function and evolution (dotted lines) based on the starburst model of Terlevich & Boyle (1993). The observed LF derived by Boyle et al. (1991) is also shown (symbols).

Boyle: QSO Evolution

421

Terlevich k Boyle (1993) have also attempted to reproduce the QSO LF at 2 < z < 3 based on the assumption that high redshift QSOs are not fuelled by accretion onto a supermassive black hole, but rather by a violent burst of star formation occuring in the cores of elliptical galaxies following their initial collapse. Reviews of the starburst model are given elsewhere (e.g. Filippenko 1992; see also the contributions by Terlevich and his collaborators in these proceedings) and so only those details pertinent to the fitting of the QSO LF will be given here. In order to predict the cores luminosities of the elliptical galaxies at 2 < z < 3, Terlevich & Boyle (1993) scaled the present-day elliptical galaxy LF in luminosity using observed mass-to-light ratios for present day ellipticals (Terlevich 1992) and the predicted mass-to-light ratios for young stellar clusters (Larson & Tinsley 1978). The LF was also scaled in space density using the predicted QSO duty-cycle at high redshifts. This was obtained by assuming that every elliptical formed by 2 = 2, and that each formation was accompanied by a single violent core starburst lasting 5 x 107yr in the QSO phase (set by the ZAMS lifetime for the minimum mass projenitor of a Type II supernova). For a typical core containing 5% of the galaxy mass, Terlevich k. Boyle (1993) found that the scaled core LF at 2 < z < 3 was in good agreement with the observed QSO LF (see Figure 8). Furthermore, the evolution of the QSO LF in a go = 0.5 universe could be accurately reproduced by the predicted decline in star formation rates (SFR a t~2) in dissipative models of galaxy formation (Larson 1974). However, Larson's models are for a continuous burst of star formation, whereas this model implies episodic bursts of star formation (~ 3 — 5 bursts/galaxy by the present epoch) with each burst occurring in a core in which the SFR has declined by t~2 since the previous burst. Such a model predicts that that QSOs at the present epoch should also be the cores of elliptical galaxies, contrary to the observation that most Seyferts are found in spiral galaxies. However, there is still sufficient uncertainty in the modelling (particularly at low luminosities) to accommodate a significant number of Seyfert nuclei in this model. Indeed, it is possible that starbursts in elliptical cores dominate the QSO LF at high luminosities (or high redshifts) with starbursts in spiral galaxies playing an increasingly dominant role at low luminosities.

4. C o n c l u s i o n s Both the super-massive black hole and starburst models are capable of reproducing the observed space densities of even the most luminous QSOs at the epoch when QSO activity was at its peak. Moreover, both models can (at least qualitatively) also account for the subsequent strong evolution of the QSO population at z < 2. However, the evolution at low redshifts is not a "natural" prediction of the supermassive black hole model (particularly in hierarchical models of galaxy formation such as CDM) . Similarly, the identification of spiral galaxies as the hosts of present day AGN is not consistent with the simple evolution of elliptical cores as QSOs in the starburst model. Clearly, further work needs to be done. Particular questions which should be addressed are: • Can either model explain the apparent differences between the evolution of low- and high-luminosity QSOs at high redshift? • Can the departures from pure luminosity evolution at low redshift be explained by either model? • How does star formation in the cores of spiral galaxies fit into the starburst model? • Why is the rate of QSO evolution at z < 2 so similar to that observed for starburst galaxies? Indeed, the agreement between QSO and starburst evolution could well be taken as an indication that the same physical processes are responsible for both phenomena. However,

422

Boyle: QSO Evolution

it should also be noted that the radio evolution of radio-loud QSOs (which currently can not be explained by the starburst model) is also identical to the evolution of optical/X-ray QSOs and starburst galaxies. In the modelling of QSO (and starburst galaxy) evolution perhaps the next step should therefore be to explain the remarkable ubiquity of the L(z) oc (1 + z)3 evolution law. I am indebted to Roberto Terlevich for providing much of the inspiration behind this work. I thank Paul Hewett, Richard McMahon and Martin Rees for allowing me to reproduce some of their diagrams in this paper. I would also like to thank the conference organisers for providing an extremely stimulating and enjoyable meeting. I was supported by a Royal Society Fellowship during the course of this work. REFERENCES BOYLE,

B.J.,

SHANKS,

T. &

PETERSON,

B.A. 1988 Mon. Not. R. Astron. Soc. 235, 935.

BOYLE, B. J., JONES, L. R., SHANKS, T., MARANO, B., ZITELLI, V. & ZAMORANI, G.

1991

In The Space Distribution of Quasars, ASP Conference Series No. 21 (ed. D. Crampton), p. 191. ASP, San Francisco. BOYLE, B. J., GRIFFITHS, R. E., SHANKS, T.,

STEWART, G. C. k. GEORGANTOPOULOS, I.

1993 Mon. Not. R. Astron. Soc. 260, 49. BOYLE, B. J., SHANKS, T.,

STEWART, G. C ,

GEORGANTOPOULOS, I. & GRIFFITHS, R.

E.

1994a, Mon. Not. R. Astron. Soc. Ssubmitted. BOYLE, B. J., MCMAHON, R. G., WILKES, B. J. & ELVIS M. 1994b Mon. Not. R. Astron. Soc. Submitted. DAVIS, M., EFSTATHIOU, G., FRENCK, C. S. & WHITE, S. D. M. 1985 Astrophys. J. 292, 371. DELLA CECA, R., MACCACARO, T., GIOIA, I. M., WOLTER, A. & STOCKE, J. T. 1992 Astro-

phys. J. 389, 491. J. S. & PEACOCK, J. A. 1990 Mon. Not. R. Astron. Soc. 247, 19. EFSTATHIOU, G. fc REES, M. J. 1988 Mon. Not. R. Astron. Soc, 230, 5P. FILIPPENKO, A. V. 1992 In Relationship Between Active Galactic Nuclei and Starburst Galaxies, ASP Conference Series No 31 (ed. A. V. Filippenko), p. 253. ASP, San Francisco. FILIPPENKO, A. V. & TERLEVICH, R. 1992 Astrophys. J. Lett. 397, L79. FRANCESCHINI, A., MARTIN-MIRONES, J. M., DANESE, L. L DE ZOTTI, G. 1993 Mon. JVot. R. Astron. Soc. 264, 35. FRANCIS, P. J. 1993 Astrophys. J. 407, 519. DUNLOP,

GEORGANTOPOULOS, I., STWEART, G. C ,

SHANKS, T.,

BOYLE, B. J. & GRIFFITHS, R.

E.

1994 Mon. Not R. Astron. Soc. Submitted. GIALLONGO, E. & VAGNETTI, F. 1992 Astrophys. J. 396, 411. GOLDSCHMIDT, P., MILLER, L., LA FRANCA, F. & CRISTIANI, S. 1992 Mon. Not. R. Astron. Soc. 256, 65P. HAENELT, M. G. & REES, M. J. 1993 Mon. Not. R. Astron. Soc. 263, 168. HAMMAN, F. & FERLAND, G. 1993 Astrophys. J. 418, 11. HENRY, P. ET AL. 1994 Astrophys. J. In press. HEWETT, P. C , FOLTZ, C. B. & CHAFFEE, F. H. 1993 Astrophys. J. Lett. 406, L43. HOOK, I. M. 1994 PhD Thesis, University of Cambridge. IRWIN, M. J., MCMAHON, R. G. & HAZARD, C. 1991 In The Space Distribution of Quasars, ASP Conference Series No. 21 (ed. D. Crampton), p. 117. ASP, San Francisco). Koo, D. C. & KRON, R. G. 1988 Astrophys. J. 325, 92. LARSON, LARSON,

R. B. 1974 Mon. Not. R. Astron. Soc. 166, 585. R. B. &; TINSLEY, B. M. 1978 Astrophys. J. 219, 46.

Boyle: QSO Evolution MACCACARO,

423

T., GIOIA, I. M. & STOCKE, J. T. 1984 Astrophys. J. 283, 486.

MACCACARO, T., WOLTER, A.

DELLA CECA, R.,

GIOIA, I. M.,

MORRIS, S. L.,

STOCKE, J.

T.

&s

1991 Astwphys. J. 374, 117.

MARSHALL, H. L., AVNI, Y., BRACCESI, A., HUCHRA, J. P., TANANBAUM, H., ZAMORANI, G.

& ZlTELLI, V. 1984 Astrophys. J. 283, 50. MARSHALL, H. L. 1991 In The Space Distribution of Quasars, ASP Conference Series No. 21 (ed. D. Crampton), p. 184. ASP, San Francisco). MCMAHON, R. G. 1991 In The Space Distribution of Quasars, ASP Conference Series No. 21 (ed. D. Crampton), p. 129. ASP, San Francisco). MITCHELL,P.S., MILLER.L. AND BOYLE,B.J. 1990 Mon. Not. R. Astron. Soc. 244, 1. REES, M. J. 1984 Ann. Rev. Astron. Astrophys. 22, 471. SAUNDERS, W., ROWAN-ROBINSON, M., LAWRENCE, A., EFSTATHIOU, G., LIS, R. S. & FRENK, C. S. 1990 Mon. Not. R. Astron. Soc. 242, 318.

KAISER, N., EL-

M., SCHNEIDER, D. &; GUNN, J. 1991 In The Space Distribution of Quasars, ASP Conference Series No. 21 (ed. D. Crampton), p. 109. ASP, San Francisco. TERLEVICH, R. J. 1992 In Relationship Between Active Galactic Nuclei and Starburst Galaxies, ASP Conference Series No 31 (ed. A. V. Filippenko), p. 133. ASP, San Francisco. TERLEVICH, R. J. & BOYLE, B. J. 1993 Mon. Not. R. Astron. Soc. 262, 491. TERLEVICH, R. & MELNICK, J. 1985 Mon. Not. R. Astron. Soc. 213, 841. WARREN,S. J., HEWETT.P. C , & OSMER, P. S. 1994 Astrophys. J. 421, 412. SCHMIDT,

Evolution of Elliptical Galaxies — a Chemo-Dynamical Model By A. C. S. FRIAgAf AND R. J. TERLEVICH Royal Greenwich Observatory, Madingley Road, Cambridge, CB3 OEZ, UK A chemical evolution model is combined with a fully hydrodynamical code to follow the evolution of elliptical galaxies from the protogalaxy stage. In this way, the single-zone assumption, usual in chemical evolution models, is dropped. This allows the investigation of radial metallicity gradients, and, in particular the formation of the high-metallicity core in ellipticals. The star formation rate and the subsequent supernova heating regulate the episodes of wind, outflow, and cooling flow, thus affecting the recycling of the gas and the chemical enrichment of the intergalactic medium.

1. Introduction In this work, a chemical evolution model is combined with a hydrodynamical model to follow the evolution of elliptical galaxies from the protogalactic stage. One motivation for the present study comes from the starburst model for QSOs (Terlevich & Boyle 1993 and references therein). In this model, the QSOs are the young cores of massive ellipticals forming most of the dominant metal-rich population in a short starburst. Since QSOs are seen up to redshifts of z ~ 5, the suprasolar metallicities required by this models should be reached by ~ 1 Gyr since the epoch of galaxy formation. This evolutionary time scale is an important constraint in chemical enrichment models. Hamann & Ferland (1992) have used one-zone chemical evolution models to investigate the chemical history of QSOs with results consistent with the starburst model of QSOs. However, during the early evolution of an elliptical galaxy, several episodes of outflow and inflow are expected and these render the one-zone model unrealistic. Therefore, in this work, we develop a multi-zone chemo-dynamical model, in which a chemical evolution model is joined to a hydrodynamical code. This model constitutes a useful tool for the study of young elliptical galaxies and their relation to QSOs. In this way we have investigated: 1) the formation of the high-metallicity core in ellipticals; 2) the radial metallicity gradients in ellipticals; 3) the enrichment of the intergalactic medium by ellipticals. The chemo-dynamical evolution model is discussed in §2. The results of some computed models are presented in § 3, and in § 4 some consequences of the present study are discussed. 2. T h e model This work is aimed at modelling the chemical evolution of elliptical galaxies with a multi-zone model. The elliptical galaxy, assumed to be spherical, is subdivided into several spherical shells and the hydrodynamical evolution of its interstellar medium (ISM) is calculated (see Friac,a 1986, 1990, 1993). In this way, hydrodynamical consistency is achieved for gas flow from and into the spherical zones. Then, the chemical evolution equations are solved for each zone, taking into account the gas flow. The number of zones chosen here was 100, with the inner boundary at 100 pc and the outer boundary at 300 kpc. t On leave from Instituto Astronomico e Geofisico, Caixa Postal 9638, 01065-970 Sao Paulo SP, Brazil 424

Friac.a ic Terlevich: Evolution of Elliptical Galaxies

425

2.1. The hydrodynamical evolution The hydrodynamical evolution of the spherically symmetric ISM is given by solving the fluid equations of mass, momentum and energy conservation:

t

CVO"*.—*

(21)

where p, u, p and U are the gas density, velocity, pressure and specific internal energy, respectively. The equation of state U = (3/2)p/p completes the system. The total binding mass M is the sum of three components: gas, stars and a dark halo. Whereas there is exchange of mass between the gas and star components through star formation and stellar mass losses (supernovae, planetary nebulae, and stellar winds), the dark halo does not take part in the processes of star formation and death and is given by a static mass distribution whose density varies as PH(r)=ph0[l + (r/rh)2]-\

(2.4)

where pho is the central halo density and r*, is the halo core radius. The star formation process and the restoring of mass to the ISM by dying stars are represented by the specific gas removal rate u and the specific gas injection rate a. The terms ap, and vp in the continuity equation couple the gas density to the star mass density p». The stars formed are further assumed to follow purely radial orbits. The stars are assumed to die either as supernovae (SNe) or as planetary nebulae, when instantaneous ejection of mass and energy occurs. The energy per unit mass injected into the gas by the dying stars can be divided in three contributions due to Type I SNe, Type II SNe and quiescent stellar mass loss (planetary nebulae and stellar winds): Uinj = (OISNIESNI/MSNI

+ asNiiEsNii/MsNii + a»Uinj,*)/a,

(2.5)

where OISNI, <*SNII and a . are the specific gas injection rate by Type I SNe, Type II SNe and quiescent stellar mass loss, respectively (a = asm + <*SNII + <**)• MSNI, MSNII and ESNI, ESNII are the mass and kinetic energy of the supernova ejecta, respectively. Here, Type I SN stands for Type la only, whereas Type Ib is included in Type II SN. The gas which is lost from stars as wind or planetary nebulae is assumed to be thermalized to the temperature given by the velocity dispersion of the stars; i.e. Uinj,* = (3/2)2 is the cooling rate per unit volume, takes into account the variations of the gas abundances predicted by the chemical evolution calculation. For the sake of simplicity, instead of considering the abundances of all elements included in the chemical evolution calculations our cooling flunction depends only on the abundances of O and Fe, which are the main coolants for T > 105 K. In the evaluation of the cooling function, the abundances of elements other than Fe and O have been scaled to the 0 abundance as j/.- = yt.p + (yt,Q - yi,p)yo/yo,Q,

(2.6)

where y, is the abundance by number of the element t (t = H, He, C, N, O, Mg, Fe), and J/,P and y,0 are the primordial (i.e. Y=0.24 and Z=0.0) and solar abundances of the element i, respectively. The atomic database used in the determination of the cooling function comes from the photoionization code AANGABA (Gruenwald & Viegas 1992).

426

Friac.a k Terlevich: Evolution of Elliptical Galaxies

As for initial conditions, the galaxy is initially purely gaseous (M. = 0), the gas has temperature To = 104 K, and its distribution is the same of the dark halo. The models are evolved until the present epoch (t = 13 Gyr) or until t — 1 Gyr (z ~ 5). Ho = 50 km s" 1 Mpc" 1 was adopted throughout this paper. 2.2. The chemical evolution The evolution of the abundances of six chemical species (He, C, N, O, Mg, Fe) has been calculated in each spherical zone. The basic equations of chemical evolution which have been solved are given in Matteucci k Tornambe (1987). A star formation rate (SFR) following a Schmidt law t/>(<) oc p" has been adopted. In particular, the SFR has the form rl>(t) = voip/po)"-1.

(2.7)

Chemical evolution models for the Galaxy require a weak dependence of the SFR on the density. A range n = 1 — 1.5 is found by Rana k Wilkinson (1986), whereas in the models of Matteucci k Francois (1989), n = 1.1 is favored over n = 2. In this work, n has been fixed as 1.2. A fiducial value for the specific SFR VQ = 10 Gyr" 1 was chosen following the results of models of chemical evolution of ellipticals, which require high stellar formation efficiency (Matteucci k Tornambe 1987; Matteucci 1992). The fiducial value for the gas density po is taken as the average gas density inside the halo core radius at the beginning of the calculations. We adopted a Salpeter IMF (x = 1.35) for stellar masses between 0.1 and 100 MQ. The main-sequence permanence time is from Burkert k Hensler 1987. Chemical enrichment occurs as stars evolve and eject gas back into the ISM via stellar winds, planetary nebulae and SNe. The restitution of mass into the ISM depends on the initial stellar mass. Single stars in the mass range 0.1 < M/MQ < 8 end their life as helium or C-0 white dwarf with mass smaller than 1.4 MQ . Single stars with masses above 8 MQ end their life as Type II SNe. No distinction is made between Type II and Type Ib, a Type Ib being considered a Type II which has lost its envelope prior the explosion. Type la supernovae are assumed to originate from binary systems of total mass in the range 3 < M/MQ < 8 in which the primary evolves until it becomes a C-0 white dwarf. Mass transfer from the slower-evolving secondary triggers C-deflagration onto the primary when the latter reaches the Chandrasekhar mass (Whelan k Iben 1973). The detailed calculation of the Type la SN rate can be found in Greggio k Renzini (1983). An important parameter in this scenario is A, the mass fraction of the IMF between 3 < M/MQ < 8 that goes to binary systems giving rise to Type la SNe. Following Matteucci k Tornambe (1987) and Matteucci (1992), we have chosen A = 0.1. This value should verified a posteriori by checking the predicted against the observed SN la rate in ellipticals. The nucleosynthesis prescriptions for single intermediate mass stars (0.8 < M/MQ < 8) are taken from Renzini k Voli (1981) (their ac = 1.5, t) = 0.33 model), those for Type II SNe are from Arnett (1991), and those for Type la SNe are from Nomoto, Thielemann k Yokoi (1984, their model W7).

3. Model results We have constructed a model grid with three free parameters: 1) M/ um = Mga, + M», the total luminous mass (here luminous stands for the X-ray or optical bands, in which the gas and the stars, respectively, are mainly observed); 2) Mh/Mium, the ratio of the halo to the luminous mass; 3) r/,. Since our models are intended to represent large ellipticals, we have chosen Mium = 2 x 1011 or 5 x 1011 MQ. Mh/Mium was set as 1 or

Friac.a & Terlevich: Evolution of Elliptical Galaxies

Mtum (lOnA/0) 2 2 5 5

1 1 3 1 3

(kpc) 5 5 7.5 7.5

Mc,

Und,cf

Mw

(10*A/ e 0 0.76 1.20 3.60 6.09

(Gyr)

(10"A/ Q ) 0.73 0.72 2.24 2.10

TABLE 1.

0.70 0.83 0.79 0.82

M

427

w

o

(lO*Af 0 ) 9.32 8.32 26.0 25.2

AW e (lO*A/ 0 ) 1.25 1.05 3.35 3.20

Model resultsi

3, representing a not too massive dark halo. From X-ray observations suggesting a dark halo more extended than the star distribution and the correlations between the galaxy luminosity and its linear size (as measured, for instance, by the de Vaucouleurs radius), we have chosen r/, = 5 kpc for M j u m = 2 x 10 11 MQ and n, = 7.5 kpc for M; u m = 5 x 10 11 MQ (see Sarazin & White 1987). Therefore, we have investigated a total of 4 models (see Table 1). We have taken as fiducial model that with M/ u m = 2 x 1011 MQ and Mh/Mium = 1.

Since the initial conditions are far from equilibrium, at the start of the calculations the gas falls towards the centre and then bounces giving rise to shocks, the gas being rapidly heated to approximately the virial temperature of the system. At the centre, the initial strong central compression gives rise to a sharp density peak. The highly efficient star formation which is occurring through the whole galaxy is even more vigorous at the centre. As a result, most of the gas in the inner few hundred pc is consumed in about 3 x 108 yr. Following the initial violent star formation burst, the first Type II SNe arise and heat the gas. The SN II rate reaches a maximum of 600 SNe (100 yrjn1 (lO 11 ^/©)- 1 at 3 x 107 yr and decays rapidly afterwards. The Type II SNe are followed by the Type I SNe, whose maximum rate of 7 SNe (100 yr)" 1 (lO11^/©)"1 is reached at 3x 108 yr. The SN I rate decreases much slower than the SN II rate. At 1 Gyr, it has decreased only 10% with respect to the peak value and at the present time (t = 13 Gyr) it is 0.22 SNe (100 yr)- 1 (lO 11 ^©)- 1 . This latter value corresponds to 0.22 SNU (1 SNU = lO" 1 0 ^© (100 yr) - 1 ) for M/LB = 8 and agrees well with the recent estimate of the SN I rate of 0.25 SNU by van den Bergh & Tammann (1991). Only the heating by Type I SNe is important in the later evolutionary phases (t > 1 Gyr) of the elliptical, when the star formation rate has become very low and, as a consequence, the SN II rate drastically decreases. In all our models, after the initial settling phase a partial wind appears; that is, a coolingflowin the inner regions of the galaxy and a wind in the outer regions. As more gas is consumed by star formation, the stagnation point separating the cooling flow from the wind moves inwards and by 1 Gyr the cooling flow region has shrunk so much that a total wind is established from 100 pc to 300 kpc. The model results are shown in Table 1. The first three columns give the model parameters: Mium, Af/,/M/um, and 77,. The following two columns indicate Mcj, the total mass collected in the inner 100 pc by the cooling flow, and tend,ef, the time for the onset of the total wind. The last three columns report mass loss quantities calculated at 1 Gyr: Mw, the total mass lost in the wind, MWto and MWiFe, the O and Fe masses ejected with the wind. From the mass loss results of Table 1, it is clear that the gas plays a very important role in the first Gyr of the galaxy's life. A significant fraction of the initial galaxy mass (approaching in some cases half the galaxy mass) is lost with the wind during this early phase. Moreover, the gas is metal rich and is therefore of great importance for the

428

Friac.a k. Terlevich: Evolution of Elliptical Galaxies

understanding of the chemical enrichment of the intergalactic medium. In view of this, if one wishes to model the evolution of galaxies at high redshifts, the present model should be preferred over models in which the galaxy is initially gasless, such as those developed by David, Forman k Jones (1990, 1991) and Ciotti et al. (1991). The early wind soon becomes very metal rich. At 1 Gyr, the fiducial model shows wind abundances of [Fe/H]=2 and [O/H]=l (metallicities are referred to the solar values). The metallicity of the gas is still higher at the galaxy centre; at r = 100 pc, [O/H]=4.5 whereas the iron shows a shallower gradient, being only 30% more abundant at 100 pc than at 100 kpc. The underlying stellar population also exhibits dramatic metallicity gradients. Again, for the fiducial model at 1 Gyr, [Fe/H] varies from 1.78 at 100 pc to 0.19 at 100 kpc, and [O/H] from 4.19 to 0.55 over the same radius range. A high-metallicity core is therefore rapidly built. The central oxygen abundance becomes suprasolar at 7 x 107 yr and the iron abundance at 5 x 108 yr, respectively.

4. Discussion The results of our models confirm the suggestion that elliptical galaxies could explain the iron abundances in the intracluster medium (ICM) of X-ray clusters of galaxies (David et al. 1990, 1991). It is interesting to note that the wind phase extending from t = 1 Gyr until now, makes a significant contribution to the iron enrichment. During this phase, a small amount of mass is ejected from the galaxy (1.45 x 1O1OM0 as compared to 7.3 x 1O1OM0 during the early phase). However, the amount of iron expelled with the wind, 8.5 x 108 MQ, is comparable to that dispersed in the early phase, 1.25 x 109MQ. This result is not unexpected, since iron arises mostly from Type la SNe and the longer lifetimes of the Type la SN progenitors result in a delayed Fe enrichment. The iron contribution to the ICM from ellipticals could be checked through the quantity (MFe/LB)ci] here, Mfe is the total iron mass in the ICM. Elliptical galaxies contribute most of the luminous mass in most X-ray luminous clusters. If we consider all the iron mass expelled by ourfiducialmodel until the present, we obtain Mf e /Mj u m = 1.04x 10~3. Considering this model as representative of cluster ellipticals, and assuming that essentially all the luminous mass in the cluster has [M/LB] = 8, typical of ellipticals, we get {MFel^B)c\ = 0.83 x 10~ 2 MQ/Z,© in very good agreement with the value inferred from observations, (MFe/Ls)ei « W~2MQ/LQ (David et al. 1990), thereby giving support to the thesis that enrichment by ellipticals accounts for the ICM iron abundances. The metallicity gradients predicted by our calculations are very plausible in view of those inferred from recent observational material (Davies, Sadler k Peletier 1993). From the reported Mg2 indexes, we have derived for NGC 4472 a metallicity of 3 (with respect to solar values, linear scale) in the nucleus and of 2 at the effective radius; for NGC 7626, the respective values are 3.1 and 1.3. The inferred galaxy metallicities are intermediate between the predictions of our fiducial model for the O and Fe: [Fe/H] varies from 1.78 at r = 100 pc to 0.51 at r = 10 kpc (the effective radius of both NGC 4472 and NGC 7626 is about 10 kpc), and [O/H] from 4.19 to 2.86 over the same radius range. The average abundance gradient (logarithmic scale) of—0.2 quoted by Davies et al. (1993) seems also to be intermediate between those predicted for the O and Fe by our model. For iron, the abundance gradient is —0.21 for 100 pc < r < 1 kpc and —0.32 for 1 kpc < r < 10 kpc; and for oxygen, the respective values are —0.17 and —0.25. The results of our models should be considered in the light of the finding of Davies et al. (1993), where the slope for the versus Mg2 within ellipticals is steeper than the equivalent relation for the nuclei of ellipticals. This has been interperted as evidence of an enrichment of Mg over Fe with respect to the solar value. Mg and O are created in

Friaca & Terlevich: Evolution of Elliptical Galaxies

429

the same Type II SNe and therefore their history is similar. In fact, in our models, the O/Fe ratio is suprasolar from r = 100 pc to r = 10 kpc. One additional important result is that a high-metallicity core is formed in a short time (t < 1 Gyr), giving support to the starburst model of QSOs. Stronger tests for the starburst model will come from the comparison of abundance ratios inferred from observations with those predicted by the present model. On the other hand, the chemical enrichment timescales predicted by our chemo-dynamical model will provide a chemical clock which could constraint cosmological scenarios. ACSF is supported by Conselho Nacional de Desenvolvimento Cientifico e Tecnologico (Brazil) under grant 20.1687/92.5. ACSF would like to thank the RGO for their hospitality.

REFERENCES ARNETT, D. W. 1991 In Frontiers of Stellar Evolution (ed. D. L. Lambert), p. 389. ASP Conf. Ser. 20. BUHKERT, A. & HENSLER, G. 1987 In Nuclear Astrophysics (ed. W. Hillebrandt, R. Kuhuss, E. Mfiller, k. J. W. Truran), p. 159. Springer-Verlag. CIOTTI, L., D'ERCOLE, A., PELLEGRINI, S. & RENZINI, A. 1991 Astrophys. J. 376, 380. DAVID, L. P., FORMAN, W. fc JONES, C. 1990 Astrophys. J. 359, 29. DAVID, L. P., FORMAN, W. & JONES, C. 1991 Astrophys. J. 369, 121. DAVIES, R. L., SADLER, E. M. & PELETIER, R. F. 1993 M. N. R. A. S. 262, 650.

A. C. S. 1986 Astron. Astrophys. 164, 6. A. C. S. 1990, In Chemical and Dynamical Evolution of Galaxies (ed. F. Ferrini, J. Franco & F. Matteucci), p. 561. ETS. FRIAQA, A. C. S. 1993 Astron. Astrophys. 269, 145. GREGGIO, L. & RENZINI, A. 1983 Astron. Astrophys. 118, 217. GRUENWALD, R. B. fc VIEGAS, S. M. 1992 Astrophys. J. Suppl. 78, 153. HAMANN. F. & FERLAND, G. 1992 Astrophys. J. Lett. 391, L53. FRIASA, FRIAQA,

MATTEUCCI, F. 1992 Astrophys. J. 397, 32. MATTEUCCI, F. & FRANgois, P. 1989 M. N. R. A. S. 239, 885.

F. & TORNAMBE, A. 1987 Astron. Astrophys. 185, 51. K., THIELEMANN, F.-K. & YOKOI, K. 1984 Astrophys. J. 286, 644.

MATTEUCCI, NOMOTO,

RANA, N. C. & WILKINSON, D. A. 1986 M. N. R. A. S. 218, 721. RENZINI, SARAZIN,

A. & VOLI, M. 1981 Astron. Astrophys. 94, 175. C. L., & WHITE, R. E., Ill 1987 Astrophys. J. 320, 32.

TERLEVICH, R. J. & BOYLE, B. J. 1993 M. N. R. A. S. 262, 491. VAN

DEN BERGH, S. & TAMMANN, G. 1991 Ann. Rev. Astron. Astrophys. 29, 363. J. & IBEN, I., JR. 1973 Astrophys. J. 186, 1007.

WHELAN,

Birth of Galaxies at z — 2 or Violent Star Formation at z = 0.4? By ARTHUR D. CHERNIN Sternberg Astronomical Institute, Moscow University, Moscow, 119899, Russia Observations of faint blue extended objects with the Hubble Space Telescope by Dressier et al. (1993) are discussed. It is argued that the objects may be regions of violent star formation in unseen galaxies at z = 0.4.

1. Nascent Galaxies? On a 6-hr Wide Field Camera exposure with the Hubble Space Telescope of the rich cluster CL 0939+4713, Dressier et al. (1993) have found an apparent group of 15-30 faint extended objects with magnitudes 22 < r < 25. The objects are typically 1" in angular size with bright central regions only a few tenths of an arcsecond in size, and are distributed over a region 40" x 20". The size and appearance of individual objects, their blue colors (known from broad-band ground-based imaging), and their clustering suggest, as Dressier et al. (1993) speculate, that they could be associated with each other, and - most importantly - that they are considerably more distant than the cluster CL 0939+4713 at z = 0.40. Furthermore, they assume that these objects could be physically associated with an unresolved, extremely blue object with the spectrum of a QSO at z — 2.055. If so, it means that at least a part of the objects may represent the early stage of formation of galaxies. Should this association of the objects with QSO be confirmed, it would have lead to the conclusion that the luminous parts of these objects have physical diameters of about 1-3 kpc. There are several cases where two or more objects are confined within a circle with a diameter less than 10 kpc. Dressier et al. (1993) hypothesize that many of the objects are "regions of intense star formation in young, coalescing systems", and that these systems "may be representative of a ubiquitous population of young or primeval galaxies" at z — 2.

2. Superassociations? Efremov & Chernin (1994) suggest an identification of these objects with, not a hypothetical, but well studied kind of region of intense star formation, i.e. superassociations that were discussed first by Baade (1963) and Ambartsumian (1964). Petrosian et al. (1984) studied 150 such objects in 57 spiral galaxies and found a mean diameter of 1 kpc, while the largest reach 4 kpc, and their typical absolute magnitudes are about —15. Perhaps the most intriguing examples of galaxies with the brightest superassociations are the systems known as clumpy irregulars. They contain several very bright supergiant HII regions. Some of them look like dense groups of galaxies; but a common rotation curve and (in some cases) a common HI envelope indicate that each such "groups" is rather a united giant irregular galaxy. A typical clump within a galaxy of this type (for example, Mrk 325 or Haro 15 - see Egiazarian et al. 1978) has a luminosity a hundred times that of 30 Dor and contains a hundred thousand OB stars (Heidmann 430

Chernin: Birth of Galaxies or Violent Star Formation?

431

1987; Efremov 1989). The size of a typical clumpy irregular is about 10 kpc. (Boesgardt and Edvards 1982). Superassociations with absolute magnitudes of around M — —15 can have visible magnitudes within the observed limits of 22 < r < 25 provided they are at the redshifts, which are much less than z = 2 of the QSO observed in the same field. Rather, it is closer to the redshift z — 0.4 of the cluster, mentioned as being in this field. If it is assumed that the objects have exactly the same redshift as the cluster, the absolute magnitudes of the objects prove to be in the interval —18 < M < —14 around the characteristic figure M = —15. (The standard Einstein-de Sitter cosmology with Hubble constant 50-100 km s" 1 Mpc" 1 and K-corrections are used here.) One can find also (in the same cosmology) the absolute linear dimensions for the objects at z = 0.4. The angular sizes of about one arcsecond or a few tenths of arcsecond correspond to 3-1 kpc, which is just characteristic for the bright superassociations. (Incidentally, the linear sizes are close to those for z = 2 because of their non-monotonic dependence on z.) The fact that the observations discovered two or more objects within an area which is about few times larger in size seems to be in agreement with what is known about superassociations in galaxies similar to, say, Mrk 325.

3. Conclusion Turning to the identification of the objects assumed by Dressier et al. (1993), it is found that they would have absolute magnitudes from M = — 2 3 t o M = —19 at 2 = 2 (with the same standard cosmology and K-corrections as above). How might several of such very luminous objects gather within a space volume of 10 kpc across? The observations by Dressier et al. (1993) treated in another, "modest", way (Efremov & Chernin 1994) reveal normal superassociations, well-known regions of violent star formation. It this case, superassociations are observed at z = 0.4, i.e. further than mentioned above. Note, that their parent galaxies present surrounding diffuse regions of lower surface brightness that are lost in the background noise of the WFC. So the galaxies themselves remain unseen; they are most probably of clumpy irregular type. The large regular galaxies of the cluster are well recognized in the same observations. To conclude, one should say that the definite test for the two versions of the identification discussed here can be provided by direct measurement of the redshifts of the objects, which seems to be a feasible task. Until then it is better to use the question marks in this note, as Dressier et al. (1993) do in the title of their paper.

REFERENCES V. A. 1964 In IAU Symp. No. 20, p. 122. Canberra, Austral. Acad. Sci. BAADE, W. 1963 Evolution of Stars and Galaxies Harvard Univ. Press, Cambridge, Mass. BOESGARDT, A. M. & EDVARDS, S. & HEIDMANN, J. 1982 Astrophys. J. 252, 487. DRESSLER, A., OEMLER, A., GUNN, J. E. & BUTCHER, H. 1993 Astrophys. J. Lett. 404, L45. EFREMOV, YU.N. 1989a Origins of Star Formation in Galaxies. Nauka, Moscow. EFREMOV, YU. N. & CHERNIN, A. D. 1994 In press. EGIAZARIAN, A. A., KAZARIAN, M. A. & KHACHKIAN, E. E. 1978 Astrofis. 14, 263. HEIDEMANN, J. 1987 In IAU Symp. 115: Star Forming Regions (ed. M. Peimbert & J. Jugaku), p. 599. Reidel, Dordrecht. PETROSIAN, A. R., SAAKIAN, K. A. & KHACHIKIAN, E. E. 1984 Astrofis. 21, 57. AMBARTSUMIAN,