Circumstellar Media in Late Stages of Stellar Evolution (Proceedings of the 34th Herstmonceux Conference, Held in Cambridge, July 12-16, 1993)

In the last throes of their lives, how do low- and high-mass stars interact with their immediate surroundings? How do a...

Author: R. E. S. Clegg | I. R. Stevens | W. P. S. Meikle

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In the last throes of their lives, how do low- and high-mass stars interact with their immediate surroundings? How do asymmetric stellar winds and the circumstellar medium affect the shape of a nebula? How are supernovae affected by a dense medium? And what do we understand of how stellar winds interact with their environments? These and many other exciting issues are addressed in these proceedings, from the 34th Herstmonceux conference held in Cambridge. Highlights of developments covered include the latest observations (including those with the Hubble Space Telescope) of stellar ejecta in planetary nebulae, novae, ring nebulae and supernovae, and a unified view of the physical processes involved; as well as the latest results on the media around supernovae 1987A and 1993J. This timely volume provides review articles that serve both as an excellent introduction for graduate students, and a handy reference for researchers; and up-to-date research papers for those who want to keep abreast of developments in the field.

Circumstellar Media in the Late Stages of Stellar Evolution

Circumstellar Media in the Late Stages of Stellar Evolution Proceedings of the 34th Herstmonceux conference, held in Cambridge, July 12-16, 1993

Edited by R. E. S. CLEGG Royal Greenwich Observatory I. R. STEVENS University of Birmingham W. P. S. MEIKLE Imperial College, London

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 for this book is available from the British Library Library of Congress cataloguing in publication data available

ISBN 0 521 46551 6 hardback

Contents

Preface

page ix

Conference Photograph

x

Conference Participants

xiii

Part one: Stellar Evolution and Wind Theory Evolution of massive stars N. Lunger Evolution of AGB stars P. R. Wood Hot star winds J. E. Drew Axisymmetric outflows from single and binary stars M. Livio Flows in clumpy CSM J. E. Dyson & T. W. Hartquist

1 15 27 35 52

Part two: Wolf-Rayet Ring Nebulae Ring Nebulae around LBVs and WR, stars L. J. Smith WR stars in the LMC T. A. Lozinskaya el al. WR Shell Nebulae R. Dufour Three-wind model for WR bubbles G. Garcia-Segura & M. Mac Low Si 19: a new Luminous Blue Variable? A. Notn et al HST images of Eta Carinae D. Ebbets et.nl

64 73 78 85 89 95

Part three: Supernovae Supernovae and their circumstellar environment B. Leibundgut... Radio supernovae and progenitor winds 5. Van Dyk et al Circumstellar interaction in supernovae C. Fmnsson SN progenitor winds J. Blondin Supernovae with dense circumstellar winds Ar. N. Chugai

100 112 120 139 148 vn

Compact supernova remnants R. J. Terlevich 153 The evolution of compact supernova remnants G. Tenorio-Tagle..l66 Massive supernovae in binary systems P. C. Joss et al 179 The progenitor of SN 1993J Ph. Podsiadloivski et al 187 Narrow lines from SN 1993J R. J. Cumming et al 192 UV spectroscopy of SN 1993J G. Sonneborn et al 198 Ryle Telescope observations of SN 1993J D. A. Green & G. G. Pooley 203 SN 1993J - early radio emission K. W. Weiler et al 207 The circumstellar gas around SN 1987A and SN 1993J P. Lundqvist 213 X-ray emission from SN 1987A and SN 1993.1 T. Suzuki et. al, .. .221 The interstellar medium towards SN 1993J in M81 D. L. King et al 227

Part four: Asymptotic Giant Branch stars Mass loss from late type stars /. Cherchneff & A. G. G. M. Helens 232 Kinematics and structure of circumstellar envelopes //. Olofs$on..246 Circumstellar shells of Long-Period Variables A. Ganger et al. ... 262 Observation of circumstellar shells with the IR AM telescopes M. Gue'lin et al 266

Part five: Planetary Nebulae Morphology and kinematics of PNe H. E. Schwarz FLIERs in elliptical Planetary Nebulae M. Perinotto et al Circumstellar dust in PN and PPN S. Kwok et al H-poor ejecta in A30 and A78 ./. P. Harrington et al The neutral envelopes of PNe P. J. Muggins et al Magnetic shaping of Planetary Nebulae /?. Chevalier Aspherical two-wind configurations V. lake

274 291 296 300 304 308 309

Part six: Novae and Symbiotic Stars Novae as tracers of mass loss M. F. Bode Light scattering in symbiotic stars //. M. Schmid & II. Schild

321 331

Poster Papers

335

Author Index

341

Object Index

343

vin

Preface

The historic series of Herstmonceux Conferences was started by the new Astronomer Royal, Sir Richard Woolley, in the late 1950's. The Royal Greenwich Observatory had by then recently finished moving its scientific operations from Greenwich to Herstmonceux Castle in East Sussex. Evidently the first few such conferences were relatively small private affairs, and there are few written records of them, but in later years they grew in size. After the moving of the RGO to Cambridge in 1989, they have been organised jointly with the University's Institute of Astronomy. Our idea for the 34th Conference was to bring together different astronomical communities who study stellar evolution, stellar winds and the physics of circumstellar media, and to bring out the common physics affecting matter around both high and low-mass stars. This volume presents the proceedings. We have included all the invited reviews and the contributed oral talks, and there is a summary listing the titles and authors of all the poster papers. Thanks are due to many people for helping to put together what was the largest-ever Herstmonceux Conference. The Organising Committee were Robin Catchpole, Robin Clegg, Peter Meikle, Jim Pringle, Anne Reynolds and Ian Stevens; Anne did a huge job as Conference Secretary and deserves special mention. The Advisory Committee were John Dyson, Alex Filippenko, Claes Fransson, Harm Habing, Alain Omont and Guillermo TenorioTagle. The RGO and the IoA gave financial support, and the International Science Foundation funded speakers from Russia. We are also grateful to Alec Boksenberg for opening the Conference and to Sir Martin Rees and Jasper Wall for speeches at the Conference Dinner, held in King's College. Finally, thanks to our editor at CUP, Adam Black, for his patient assistance. Robin Clegg, Ian Stevens and Peter Meikle

IX

J 1. O. Hashimoto 2. N. Patat 3. B. Cadwell 4. G. Garcia-Segura 5. R. Corradi 6. R. Oudmaijer 7. K. Borkowski 8. 3. Blondin 9. D. Green 10. J. Solf 11. T. Forveille 12. V. Icke 13. P. Tuthill 14. J. Blommaert 15. A. Gauger 16. P. Podsiadlowski 17. R. Wolstencroft

18. 19. 20. 21. 22. 23. 24. 25. 26. 27. 28. 29. 30. 31. 32. 33. 34.

K. Justtanont W. van der Veen P. Joss M. Kolman S. Zhekov F. Fagotto A. Wootten K. Exter C. Fransson R. Tuffs M. Perinotto J. Pringle G. Wynn-Williams Xiao-Wei Liu H. Habing E. Bakker R. Cid Fernandez

35. 36. 37. 38. 39. 40. 41. 42. 43. 44. 45. 46. 47. 48. 49. 50. 51.

N. Langer D. King P. Harrington J. Walsh R. Tweedy R. Chevalier R. Cumming S. Scuderi R. Schulte-Ladbeck H. Schwarz F. Pijpers T. Harries M.. Clampin C. Koempe D. Ebbets S. Lorenz-Martins F. de Aravjo

52. 53. 54. 55. 56. 57. 58. 59. 60. 61. 62. 63. 64. 65. 66. 67. 68.

P. de Laverney E. Huguet P. Wood L. Stanghellini P. Garcia-Lario W . Maciel P. Pavelin P. Huggins H. Olofsson J. Yates P. Vilchez G. Sonneborn D. Schaerer M . Groenewegen M . Barlow J. Drew P. Williams

69. 70. 71. 72. 73. 74. 75. 76. 77. 78. 79. 80. 81. 82. 83. 84. 85.

I. Yamamura D. Hutsemekers P. Lundqvist N. Chugai K Weiler L. Smith A. Skopal A. Harpaz M . Lewis C. Charbonnel S. Van Dyk E. Terlevich R. Terlevich L. Bianchi A. Not a A. Pasquali G. Mellema

86. M. Kato 87. H. Walker 88. J. Pacheco 89. C. Rossi 90. A. Heske 91. R. Dufour 92. A. Manchado 93. A. Michalitsianos 94. B. Pagel 95. M. Wrigge 96. M. Bryce 97. C. Guilain 98. A. Frank 99. N. Berruyer 100. J. Nichols 101. T . Lozinskaya 102. MI. Livio

103. 104. 105. 106. 107. 108. 109. 110. 111. 112. 113. 114. 115. 116. 117. 118. 119.

N. Netzer N. Soker R. Clegg Z. Han P. Murdin A. Reynolds I. Stevens P. Meikle U. Munari B. Leibundgut S. Etoka A. Tielens H. Schmid G. Tenorio-Tagle R. Bachiller T. Suzuki B. Hassall

Conference Participants

R. Bachiller Yebes, Spain. E. Bakker SRON, Netherlands. J. A. Baldwin MRAO, Cambridge, UK. M. J. Barlow University College London, UK. N. Berruyer Obs. de la Cote d'Azur, France. L. Bianchi STScI, USA. J. Blommaert Inst. d'Asirophysique, Paris, France. J. M. Blondin North Carolina State Univ., USA. M. F. Bode Liverpool John Moores Univ., UK. K. J. Borkowski Univ. of Maryland, USA. M. Bryce Univ. of Manchester, UK. B. C. R. N.

Cadwell Penn State, USA. Charbonnel O6s. de Geneva, Switzerland. A. Chevalier Univ. of Virginia, USA. N. Chugai Russian Academy of Sciences, Moscow, Russia. R. Cid Fernandez RGO, Cambridge, UK M. Clampin STScI, USA. R. E. S. Clegg RGO, Cambridge, UK. R. Corradi Univ. of Padova, Italy. P. Cox MPI f. Radioastronomie, Bonn, Germany. A. Crotts Columbia Univ., USA. R. J. Cumming RGO, Cambridge, UK. R. J. Davis NRAL Jodrell Bank, UK. F. de Aravjo Obs. de la Cote d'Azur, France. P. de Laverny Univ. Montpellier, France. J. S. B. Dick RGO, Cambridge, UK. J. E. Drew Oxford Univ., UK. R. E. Dufour Rice Univ., USA J. E. Dyson Univ. of Manchester, UK. D. Ebbets Ball Aerospace Systems Group, Boulder, USA. S. Etoka 06s. de Paris-Meudon, France. K. Exter Univ. St. Andrews, UK. F. T. A. C.

Fagotto Padova Observatorio, Italy. Forveille 06s. de Grenoble, France. Frank Univ. Minnesota, USA. Fransson Stockholm Observatory, Sweden.

G. Garcia-Segura Univ. of Illinois at Urbana-Champaign, USA. P. Garcia-Lario WE Observatory, Spain. A. Gauger Technische Universitat Berlin, Germany. A. Glassgold New York Univ., USA. D. A. Green MRAO, Cambridge, UK. M. Groenewgen Astronomical Institute, Amsterdam, Netherlands. M. Guelin IRAM, France. C. Guillain CNRS Toulouse, France. H. J. Habing Sterrewacht Leiden, Netherlands. Z. Han loA, Cambridge, UK. A. Harpaz Technicon, Israel. T. Harries University College London, UK. J. P. Harrington Univ. of Maryland, USA. 0 . Hashimoto Seikei University, Japan. B. Hassall RGO, Cambridge, UK. A. Heske ESTEC, Netherlands. R. Hills MRAO, Cambridge, UK. P. J. Huggins New York Univ., USA. E. Huguet 06s. de Paris-Meudon, France. D. Hutsemekers Astrophysical Inst of Liege, Belgium. V. Icke Sterrewacht Leiden,

Netherlands.

P. C. Joss MIT, USA. K. Justtanont Institut d'Astrophysique, Paris, France. M. Kato Keio Univ., Japan. D. L. King RGO, Cambridge, UK. H. Kley IoA, Cambridge, UK. C. Koempe Uni-Sternwarte, Germany. Sun Kwok Univ. of Calgary, Canada. N. Langer MPI f. Astrophysik, Germany. J. Lefevre 06s. de la Cote d'Azur, France. B. Leibundgut UC Berkeley, USA. B. M. Lewis Arecibo Observatory, USA. X-W. Liu University College London, UK. M. Livio STScI, USA, xm

xiv

Conference Participants

S. Lorenz-Martin 06s. de la Cote d'Azur, France. T. A. Lozinskaya Sternberg Inst., Russia. P. Lundqvist Stockholm Observatory, Sweden. W. Maciel IAG/USP, Brazil. A. Manchado Univ. of Illinois, USA. W. P. S. Meikle RGO, Cambridge, UK. G. Mellema Sterrewacht Leiden, Netherlands.

A. Michalitsianos NASA/GSFC, USA. U. Munari 06s. di Padova, Italy. P. G. Murdin Royal Observatory Edinburgh, UK. N. Netzer 0RT- Braude College, Israel. J. Nichols-Bohlin NASA/GSFC, USA. A. Nota STScI, USA. H. Olofsson Stockholm Observatory, Sweden. A. Omont Institut d'Astrophysique, Paris, France. R. Oudmaijer Kapteyn Lab. Groningen, Netherlands.

J. Solf MPI f. Radioastronomie, Bonn, Germany. G. Sonneborn NASA/GSFC, USA. L. Stanghellini Obs. di Padova, Italy. E. Stengler RGO, Cambridge, UK. I. R. Stevens IoA, Cambridge, UK. T. Suzuki Univ. of Tokyo, Japan. G. Tenorio-Tagle Univ. de La Laguna, Spain E. Terlevich RGO, Cambridge, UK. R. J. Terlevich RGO, Cambridge, UK. A. G. G. M. Tielens NASA Ames, USA. R. Tuffs MPI f. Radioastronomie, Bonn, Germany. P. Tuthill MRAO, Cambridge, UK. R. Tweedy Steward Observatory, USA. S. D. Van Dyk NRL, USA. W. van der Veen Columbia Univ., USA. S. Vessey MRAO, Cambridge, UK. J. M. Vilchez Inst. Astrofisica de Canarias, Spain.

J. Pacheco IAG/USP, Sao Paulo, Brazil. B. E. J. Pagel NORDITA, Denmark. A. Pasquali Univ. of Firenze, Italy. F. Patat Padova, Italy. P. Pavelin NRAL Jodrell Bank, UK. M. Perinotto Univ. of Firenze, Italy. F. Pijpers Uppsala Obs., Sweden. Ph. Podsiadlowski IoA, Cambridge, UK. J. E. Pringle IoA, Cambridge, UK. G. G. Pooley MRAO, Cambridge, UK.

H. J. Walker RAL, UK. J. R. Walsh ST-ECF Garching, Germany. N. Walton RGO, Spain. K. W. Weiler NRL, USA. P. M. Williams Royal Observatory Edinburgh, UK. R. D. Wolstencroft Royal Observatory Edinburgh, UK. P. Wood MSSSO, Australia. A. Wootten NRAO, Virginia, USA. M. Wrigge Hamburger Sternwarte, Germany. G. Wynn-Williams Univ. of Hawaii, USA.

A. Reynolds RGO, Cambridge, UK. C. Rossi Sapienza. Italy.

I. Yamamura Univ. of Tokyo, Japan. J. Yates NRAL Jodrell Bank, UK.

D. Schaerer Geneva Observatory, Switzerland. H. M. Schmid ETH Zentrum, Switzerland. R. Schulte-Ladbeck Univ. of Pittsburgh, USA. D. Schwarz IoA, Cambridge, UK. H. E. Schwarz ESO, Chile. S. Scuderi STScI, USA. A. Skopal Liverpool John Moores Univ., UK. L. J. Smith University College London, UK. N. Soker Oranim - University Division, Israel.

S. Zhekov Obs. di Arcetri, Italy.

Evolution of massive stars N. Langer MPIf. Astrophysik, Karl-Schwarzschild-Str. 1, D-8511,0 Garching, FRG

Abstract Recent results of the theory of massive star evolution are discussed. We divide the regime of massive stars in a "low" mass and a high mass part, and show that the evolution, the basic theoretical problems in their modeling, and the display of circuinstellar matter are quite different for stars from both parts. For stars in the lower considered mass regime, it is shown that our ignorance about, internal mixing processes is the main source of uncertainty in stellar model calculations. Mixing processes related to thermal convection are discussed, and their effect on the observable stellar parameters and presupernova structure are sketched. The role of mixing induced by differential rotation is also briefly described. We argue that the supernova stage is a good possibility to investigate the circumstellar material of these objects and describe its high diagnostic potential for the whole presupernova evolution. Our understanding of stars in the upper mass range, i.e. the most massive stars, also suffers from uncertainties related to internal mixing. However, we argue that it is the mass loss process which dominates their global evolution. The evidence that those objects do lose the major part of their initial mass before they collapse is discussed, together with the possibility of the display of circumstellar nebulae during their hydrostatic evolution. Finally, the fate of very massive stars is discussed.

1 Introduction As usual in the astrophysical literature, we mean stars sufficiently massive in order not to develop into white dwarfs here (i.e. Mz.4A/s ^ 8 A/©, cf. Iben and Renzini 1983; ZAMS stands for zero age main sequence) when we speak of massive stars. Thus, this review deals with stars in the mass range ~8M0— ~2OOM0, where the upper end of the regime is defined by the upper end of the stellar initial mass function. It is to be mentioned right at the beginning of this paper, that — even at a given constant metallicity Z — stars in different regimes of this mass range have very different evolutionary characteristics. This is true for their internal evolution as well as for the evolution of their surface properties (e.g. their tracks in the HR diagram): Massive stars with masses close to 8 MQ may explode due to carbon deflagration (Type 1^ SN; cf. Iben and Renzini 1983), or collapse due to electron capture (cf. e.g. Nomoto 1984). At somewhat higher mass, stars probably make "standard" Type II SNe due to iron core collapse (see 1

2

N. Langer: Evolution of massive stars

Woosley and Weaver 1986), while still more massive stars may form black holes (cf. Maeder 1992). Further up in the mass range stars may explode again since they end up as low mass objects, but they will appear as Type I SNe as they are devoid of hydrogen (cf. Woosley et al. 1993, Langer 1994; cf. also Sect. 3). And for the uppermost masses, there is (at low metallicity) still the possibility of exploding before the formation of an iron core, clue to the e ± -pair instability (cf. e.g. Woosley 1986, El Eid and Langer 1986). As diverse as the internal evolution of massive stars are their HRD tracks. The less massive of the massive stars evolve always to the right of the main sequence. They stay on that side for their whole life, most of them using up all the available space up to the Hayashi-line. The very massive stars, on the contrary, never get really cold, i.e. their surface temperature always remains hotter than ~ 10 000 K. Their tracks generally cross the main sequence once and lead into the hot part of the HR diagram. The described diversity makes the necessity evident to divide the considered mass range into peaces and to discuss stars in each piece separately. Since this paper is in a book about circumstellar matter, and since we can only perform a qualitative investigation of stellar evolution here, we just give a qualitative discrimination of stars of two different groups: the "just massive stars" (JMS), which are characterized by the fact that their total mass remains basically constant during their whole evolution form the main sequence to the supernova stage, and the "very massive stars" (VMS), which do lose the major part of their initial mass. The JMS, due to the fact that they keep most of their hydrogen-rich envelope, may become red supergiants during their post-main sequence evolution, and as such they may lose several solar masses of their material and develop a rather slow and dense circumstellar wind. This material is, however, unlikely to be directly observable, unless it is illuminated by the supernova outburst of the central star, as in the case of SN 1987A. Otherwise, the circumstellar matter may manifest itself only indirectly, e.g. due to an infrared excess in the spectrum of the central star (cf. Stencel et al. 1989). The VMS do not only accumulate much more circumstellar matter due to their large mass loss rates, but since they develop into very hot and compact stars they eject fast winds which can compress the previously ejected material, and radiate a hard photon field which can ionise it. Consequently, many of them develop nice ring nebulae or spherical circumstellar shells. There is of course no sharp borderline between JMS and VMS, which — according to recent stellar evolution calculations — is somewhere between 25 and 40 Me in our Galaxy (cf. Schaller et al. 1992, Woosley et al. 1993).

N. Langer: Evolution of massive stars

3

Note, however, that due to the metallicity dependence of stellar winds (Kudritzki et al. 1987, Leitherer and Langer 1991) this borderline depends strongly on the initial stellar metallicity. In the following Section we will try to illustrate the problems in modelling the JMS, which involves basically internal mixing processes. Sect. 3 deals with the VMS, where the main problem concerns the mass loss. Some final remarks are given in Sect. 4.

2 The "just massive stars" The JMS have been defined as massive stars which keep most of their mass until the end of their hydrostatic evolution. In fact, a. weaker condition may be chosen, though it is a bit harder to understand, namely that a substantial part of the hydrogen-rich envelope remains bound to the star: it ensures that the H-burning shell source is almost unaffected by the mass loss, and even more so is the helium core and thus all the burning stages following He-burning (Maeder 1981, Chiosi and Maeder 1986). With this definition, e.g. a 25 M@ star losing some 5-10 A/,v, during its evolution is still a JMS: it develops a He-core of ~ 8 M 0 (when adopting the Schwarzschild criterion, see below), i.e. it would retain a H-rich envelope of ~7-12 MQ in the above example. Since the topic of this conference concerns late stages of stellar evolution, we only want to mention some problems with JMS main sequence models (cf. Langer 1993, for more details): compared to observed main sequence JMS the models are — on average — too faint, too hot, and lack helium and nitrogen enrichment (cf. Herrero el al. 1992, Langer 1992). Increased mass loss in model calculations leads in principle to larger luminosity-tomass rations, cooler main sequence models, and eventually even to surface enrichment; however, the required mass loss rates are one order of magnitude larger than the observed rates, which are by far not so uncertain. I.e., mass loss is not the right solution. Internal mixing appears to be rather promising. Zahn (1983) and Chaboyer and Zalin (1992) have proposed that differential rotation may lead to substantial mixing in the stellar interior, and Maeder (1987) showed that — for M = 40 MQ — this may even lead to homogeneous evolution. Langer (1992) and Denissenkov (1993) have applied this mechanism to somewhat smaller initial masses and also found a substantial mixing to occur, while Fliegner and Langer (1994), with a selfconsistent treatment of rotating stars, showed that even at 5 MQ rotational mixing may substantially alter the stellar evolution on the main sequence. Certainly, also the post-MS evolution is affected by this kind of mixing (cf. Langer 1992);

4

N. Langer: Evolution of massive, stars

however, systematic studies are still lacking, and the parameters involved in the description of rotational mixing are still weakly restricted. It is to be mentioned here, that models of close binary systems also have the potential to cope with the above problems, but also here a thorough discussion is not yet available (cf. Langer 1992; Joss, this volume; and references therein). When a JMS has finished core hydrogen burning, it develops into a supergiant. Due to the fact that JMS models in the effective temperature regime 4.0 ^ log Tejj ^> 3.6 (the precise numbers depend on the internal composition profile and metallicity) are thermally unstable (cf. e.g. Tuchman and Wheeler 1989ab) — a fact which agrees remarkably well with the observed low population density in that regime of the IIR diagram; cf. Blaha and Humphreys 1989 —, only two long-lived states are possible for post-MS JMS: that of blue supergiants (BSGs) and that; of red supergiants (RSGs). However, despite of this simplicity (especially when compared with VMS; cf. Section 3) it is today by no means clear which of these two states are occupied by JMS of a given mass and metallicity, or in which order they occur during the life of a JMS. We want to emphasize that the situation is much different in the (often better known) case of intermediate and low mass stars (i.e. MZAMS <; 8 MQ): e.g. at Z — ZQ they always become red gia.nts/supergiants upon core hydrogen exhaustion, they generally develop blue loops above a certain mass, and turn into AGB stars after core helium depletion, independent of all the uncertainties which enter the respective stellar evolution calculations. In this case, the qualitative aspects are discussed as a function of the input physics, e.g. how far blue loops extend towards hot temperatures, or above which critical mass exactly the blue loops occur (e.g. Matraka et al. 1982, Huang and Weigert 1983, Bertelli et al. 1985, Stothers 1991). In the case of JMS, the discussion is still in the qualitative phase: we do not know where in the HR diagram He-burning is started, whether the evolution occurs as BSG—>RSG or vice versa (note that due to the dredge-up process in the R.SG stage one may hope to observation ally disentangle the post-RSG BSGs from the others on the basis of their surface abundances; cf. Lennon et al. 1992, 1993; Fitzpatrick and Bohannan 1992), and the case of SN 1987A reminded us that we are not even sure about which of the two possible states (ie. BSG or RSG) is the pre-SN stage. The basic problem which gives rise to the huge uncertainties quoted above exists as long as stellar models are evolved on the computer: which is the correct treatment of convection in the presence of a mean molecular weight (/t) gradient? The extreme answers to this quest ion are: 1) adopting the Schwarzschild criterion for convection, i.e. ignoring the stabilizing effect of

N. Longer: Evolution of massive stars

log -feff

Fig. 1. Tracks of three 20 A/0 sequences of Z = 0.005 in the HR.diagram, computed with different assumptions on the efficiency of convection in regions of varying mean molecular weight. The short-dashed line indicates a track computed with the Schwa rzschi Id criterion, in the dot-dashed track the Ledoux criterion was used, and for the solid track an intermediate case, namely the semiconvection model of Langer el al. (1983), has been applied. Big dots mark core helium ignition and exhaustion, and asterisks mark the pre-SN positions (cf. Langer el al. 1989, for details).

/i-barriers altogether, or 2) adopting the Ledoux criterion, i.e. to prevent any mixing at all in the presence of (sufficiently strong) //-barriers. Following Kato's (1966) stability analysis, an intermediate treatment has been developed (Langer et al. 1983. 1985), and Fig. 1 illustrates the consequences for the evolution of the surface properties at the example of a 20 M.-, star at roughly LMC metallicity. The reason why the evolutionary tracks depend on the treatment of internal mixing are twofold: a) mixing at the H/He-interface during core Hburning (cf. e.g. Chiosi and Summa 1970; Stothers and Chin 1975, 1976) and — even more important — during the contraction phase towards helium ignition (Langer et al. 1985, Bressan et al. 1993) determine the conditions for the hydrogen burning shell during core helium burning, and b) mixing at the He/C+O-interface determines the size of the He-burning convective core and thereby the core He-burning life time, and the C+O-core mass, which in turn determines completely the internal evolution beyond core helium burning (cf. Fig. 2). Since core He-burning stars in the mass range 15 - 30A/(T, are in a. delicate balance between the BSG and RSG states in general, almost any change in the internal composition profile can have a large impact on the predicted BSG/RSG lifetime ratio (cf. Langer 1991). In addition to the problem of the feedback of /t-barriers with convection,

N. Langer: Evolution of massive stars 20M_sun — He — Ne — burning 9.0

•

'

^c^S1

'

log 18c

8.8 8.6 8.4 8.2

Jr

'-•

8.0 4 log

5 rho_c

Fig. 2. Evolution of two 20 MQ tracks in the logTc — \ogpc -plane. The upper track has been computed using the Schwarzschild criterion for convection, in the lower one semiconvection according to Langer et al. (1983) has been applied. The time distance between two neighboring crosses on each track is 5000 yr. further mixing processes are potentially important for JMS. One is the so called convective overshooting, which designates the possibility of convective motions to penetrate into radiatively stable stratified layers adjacent to convection zones. Only the effects of overshooting in stellar convective cores (e.g. Schaller et al. 1992, Bressan et al. 1993) and of convective envelopes (Stothers and Chin 1991; and references therein) have yet been systematically investigated, and at present the importance of this process for JMS is still a matter of debate (see e.g. Langer 1991, Schaller et al. 1992). A second family of mixing processes with almost unknown efficiency are the rotationally induced instabilities, which — at least in rapidly rotating stars — may significantly alter the whole internal chemical structure of the star (see above). Also on this subject, much work remains to be done. It has been mentioned at the beginning of this Section that JMS may well lose appreciable amounts of mass before they explode as supernova. It is important to note here that even for precisely known mass loss laws, i.e. mass loss rates as function of the stellar surface conditions, the amount of matter lost during the life of a star of given initial mass and chemical composition can not be accurately predicted due to the large uncertainty of the effective temperature evolution mentioned above. E.g., we know that the mass loss rate in the RSG stage exceeds by far the BSG mass loss rate. However, since we do not know which star spends which fraction of its lifetime as a RSG, the total amount of lost mass remains very uncertain. In this respect it is the circumstellar matter and the supernova event

N. Langer: Evolution of massive stars

7

which provide most useful constraints for the whole pre-SN evolution, and in particular for mass loss rates and total amounts of mass lost; many examples for this can be found in the present book: E.g., the early radio (van Dyk, this volume) and X-ray (Fransson, this volume) emission of supernovae provide stringent limits on the wind density M/v^ of the progenitor star. The same quantity may also affect the early optical light curve of Type II SNe, as shown by Hoflich et al. (1993) for the case of SN 1993J. SN 1987A provides a unique example which is widely discussed in the literature. Let us just mention that the progenitor mass loss history is traced by the famous ring nebula (cf. e.g. Lundqvist, this volume), and the light curve and spectral evolution, which provide envelope and He-core mass, allowed the best estimate of the total mass lost by a. massive star within its hydrostatic evolution (Arnett et al. 1989, Hillebrandt and Hoflich 1989). Also the SN emission at late times may yield information about the circumstellar matter density at moderate distance from the progenitor star and thus probe the stellar mass loss up to some 104 yr before the actual SN explosion (cf. Leibundgut, this volume; Clmgai, this volume). The quoted contributions contain evidence for unexpectedly large wind densities in quite a number of progenitor stars which probably fall into the .IMS mass range since they became Type II SNe. Many interesting results from this area of research are to be expected in the near future.

3 Very massive stars Since a very recent review on VMS evolution is published elsewhere (Langer 1994), we will focus here on topics related to circumstellar matter and late evolutionary stages of VMS. VMS have been defined as stars which do lose the major part of their initial mass before they collapse (cf. Sect. 1). Though such a scenario is largely favored by many observational facts (cf. Langer 1989, Meynet et al. 1994) as well as by theoretical considerations (cf. Langer 1989, 1989a), one would like to have a. direct "proof" for very massive stars evolving into rather low mass objects. If this scenario is indeed correct, then the uncertainties in internal mixing processes discussed for JMS in Sect. 2 play only a minor role for the understanding of VMS (though the VMS may be an ideal tool to investigate them, since the large mass loss allows us to quasi look deep inside the stars; cf. Langer 1991a). Instead, what almost completely determines the JMS evolution is the mass loss process: note that current models predict e.g. a 60 MQ star to evolve into a. star of roughly 5 M 0 (Langer 1989, Schaller et al. 1992, Woosley et al. 1993). Let us concentrate on two possible ways to "proof" this — at first glance

8

N. Langer: Evolution of massive stars

surprising — type of scenario, namely that VMS lose some 90% of their initial mass during long lasting evolutionary phases, i.e. core hydrogen and helium burning: one is to show that a correspondingly large amount of freshly ejected matter exists close to the stars which are potential descendants of once very massive stars, and the other is to look for supernovae from such objects.

3.1 VMS ejecta When a VMS — say a 60 M® star with Z = Z@ — ends up as a WC/WO star of ~ 5 M@, it loses some 55 MQ of material during its H- and He-burning phase of evolution. Since the resulting WR. star is very hot and luminous it may be expected that the ejected material appears as a visual nebula. In fact, we know that an appreciable number of WR stars are associated with ring nebulae (Chu et al. 1983, Chu 1991, Lozinskaya, this volume), many of which can be shown to have been ejected by their central WR object on the basis of their kinematical or chemical properties. What kind of nebula can we expect from current VMS models? Figs. 3 and 4 show the time dependence of the mass loss rate and the wind velocity for a 60 MQ sequence which was computed by using the latest available input physics, i.e. OPAL opacities (Iglesias et al. 1992), O star mass loss according to Lamers and Leitherer (1993) and due to pulsational instabilities (cf. Langer 1994, for details), and mass dependent WR mass loss (Langer 1989). The wind velocity has been computed as a function of the escape velocity, using data of Abbott (1982) and Kudritzki and Hummer (1990) for 0 stars and of Hamann et al. (1993) for WR. stars for calibration. We see from Figs. 3 and 4 that, during the first 2 million years, the star — corresponding to a normal 0 star in this period — has a very fast wind and a relatively small mass loss rate. This wind will certainly excavate a large bubble around the central star, which provides the action space for the later nebula formation. Using analytic relations for wind blown bubbles (cf. Shull 1993) and adopting typical parameters for our 0 star (M - 2 10~6 MQ yr~l, VQO = 3500km s'1) and an ISM density of n = lcm~3, the bubble size will be about 80 pc in diameter at an age of 2 106 yr. In the ensuing WNL phase the mass loss rate is by one order of magnitude enhanced, and the wind velocity is somewhat lower. Then, at core hydrogen exhaustion, a large amount of mass (5 —10 MQ) is lost with a rather small wind velocity (some 100 km s" 1 ) in a Luminous Blue Variable phase (cf. Langer 1994). This slow moving gas can be expected to be accumulated and compresses in a shell by the successive fast and energetic WR wind.

N. Langer: Evolution of massive stars

Fig. 3. Mass loss rate as a function of time for a. 60 M© sequence at Z=2% (cf. Langer 1994 for details). The first broad maximum corresponds to a first WNL stage. At t ~ 3.35 106 yr, core H-burning is terminated, and an LBV phase occurs. Finally, mass dependent VVR mass loss according to Langer (1989) has been used.

4OOO

3OO0

2OOO -

1ODO -

p

0

Fig. 4. Final wind velocity as a function of time corresponding to the mass loss rate shown in Fig. 3; see text for details. Garcia-Segura and Mac Low (this volume) performed 2D hydrodynamic computations of wind bubbles around massive stars, using a simplified schematic time behavior of the central star wind, which resembles that displayed in Fig. 3. They found a nice agreement of their results with observed WR nebulae like NGC 6888. Note that they assumed a RSG phase instead of the LBV phase in the above example. However, both, RSGs and LBVs, are characterized by large mass loss rates and by wind velocities which are slow compared to the very fast WR winds. I.e., also from post-LBV WRs one may expect the formation of a shell type nebula.

10

N. Langer: Evolution of massive stars

It remains to be investigated whether the expected life times of such nebulae agree with the observed number of them (cf. Lozinskaya, this volume), and whether they occur at the evolutionary stage/observed WR subtype where they are expected. However, the basic fact of the presence of many of them, and the success of the theoretical models so far (cf. Garcia-Segura and Mac Low, and references therein) provide a basic support of the idea that VMS do lose the major part of their mass before they explode. In fact, further studies of WR nebulae promise much insight into the details of the progenitor mass loss as a function of time. A second and independent possible evidence in favor of large amounts of mass being lost during the VMS evolution may be through the detection of 7-rays form the radioactive decay of 26Al: 26.4/ is synthesized during core hydrogen burning in massive stars (see e.g. Prantzos et al. 1986), and its long half life of roughly 106 yr provides the possibility that much of the 2eAl decays only after it has been ejected into space so that the produced l.SMeV photons can escape without interaction. Note that galactic 26Al emission has been discovered by HEAO-3 (cf. Mahoney 1982), and — among other sources — WR stars have been proposed to eject radioactive 26Al into the ISM. Fig. 5 shows the expected time dependence of the number of 7-photons from the 26Al decay for the same 60 MQ calculation for which the mass loss rate is shown in Fig. 3 (cf. Langer 1994, for details), assuming a distance to the source of 1 kpc. The flux is turned on in the WNL stage and reaches a maximum in the ensuing LBV phase. During the following WNE and WC stages, the flux declines exponentially with time, since basically all the produced 26Al has been ejected before. Note that at t ~ 4.1 106 yr the star explodes as a supernova, which adds an amount of 5 10~5 MQ of explosively synthesized 26Al according to Woosley et al (1994). Clearly, the detection of a compact source of l.SMeV photons related with a VMS type object would be a strong evidence for significant mass loss. The photon flux, together with the distance, would give hard lower limits to the amount of ejected matter. Note that 7-ray astronomy is in the shape to perform a detailed spatially resolved mapping of the l.SMeV flux at present and individual sources may be recognized in the near future (Diehl et al. 1993); a potential detection of the WC star 7 Velorum has been recently reported (Diehl, priv. communication). Our calculations point out that not only WR stars but also LBVs can be expected to be prominent 1.8 MeV photon sources. In this respect, P Cygni is the most outstanding candidate.

N. Langer: Evolution of massive stars

11

26AI — d e c a y — p h o t o n — f l u x

4x1 0 c

2x10

6*10c

time

Fig. 5. 1.8 MeV photon flux from the decay of '-''Al as function of time (in yr), computed for a 60 MQ sequence at Z=2% (c.f. Fig. 3, and Langer 1994, for details). The nuclear reaction rates of Caughlan and Fowler (1988) have been used in that model. Note that the smaller -6Al production rate proposed by Iliadis et al. (1990) would reduce the expected flux by about a factor of 4. For the absolute flux value, a distance of 1 kpc has been assumed.

3.2 Supernovae from VMS ? The question whether or not VMS may develop into supernovae has been investigated recently by Woosley et al. (1993, 1994). Here, we just want to briefly outline the main conclusions (cf. also Langer 1994). Whether or not the core collapse in a given massive star yields a SN explosion depends sensitively on its iron core mass: larger iron cores are less likely to explode (Woosley and Weaver 1986). Though the iron core mass of (non-mass losing) massive stars is not necessarily a monotonic function of the stellar mass throughout, more massive stars develop more massive iron cores in general (cf. Woosley 1986). I.e., from non-mass losing stars one would expect the existence of a critical stellar mass above which core collapse leads to black hole formation and only below a core collapse SN occurs. How does the VMS mass loss affect this picture? The mass dependent WR mass loss leads to a mass convergence, i.e. the final stellar masses of VMS are almost independent of the initial stellar mass for a wide range of ZAMS masses, with typical final masses of the order of 5M© (cf. Schaller et al. 1992, Meynet et al. 1994). Since the pre-SN models are devoid of hydrogen, their total mass is equal to their He-core mass (cf. Table 1). One is thus tempted to conclude that, since the Hecore mass greatly determines the evolution through the final burning stages (cf. Arnett 1978, Sugimoto and Nomoto 1980), a supernova explosion is the

12

N. Langer: Evolution of massive stars Table 1. Core masses as function of the initial mass (in MQ) MzAMS

Mfinal

1

35 40 1 60 1

15.2 11.1 4.3

502

50.0

MC/o

MFe

12.2 11.1 4.3

3.7 3.2 3.0

1.45 1.42 1.40

23

-15

2.45

MHe

Notes: 1) From Woosley et al. (1993). 2) From Woosley (1986). likely result of VMS evolution, since a 5 MQ He-core also corresponds to a 15 MQ ZAMS mass, and SN 1987A showed that 15 — 20 MQ stars do in fact explode as supernovae. However, things are a bit more complicated. Woosley et al. (1993) have shown that the late evolution of low mass VVR. stars descendant of VMS of a certain final mass (4.25 MQ in their example) is quite different from the late evolution of a star of the same final mass which was not very massive in the past. The reason is that the large mass loss during core helium burning of the once very massive star establishes a. very shallow He-gradient, which allows the formation of a much larger C/O-core in this case, compared to the star which was never very massive. Since mass loss beyond core helium exhaustion is negligible, the C/O-core mass is then a good parameter in the sense that its value describes the late evolutionary phases almost independently of the previous evolution (see Table 1). Nevertheless, VMS mass loss reduces the final C/O-core mass to values (~ 3 MQ) which correspond to ZAMS masses of the order of 20 MQ in constant mass calculations, which still makes a SN explosion very likely, at least compared to the case of the evolution of very massive stars at constant mass (cf. Table 1). However, note that not only the iron core mass determines whether or not a SN explosion may occur, but to some extent also the density profile above the iron core (cf. Woosley et al. 1993); i.e., in order to settle the question which stellar models exactly would explode one needs a better understanding of the core collapse supernova mechanism.

4 Final remarks In the previous Sections we have shown that the theory of massive star evolution is still far from being in a final shape. Today, very important

N. Langer: Evolution of massive stars

13

parts in such a theory are still lacking or are oversimplified. This concerns points which are generally designated as "input physics" like mixing or mass loss theories, but also genuine stellar structure problems as a theory for the stellar radius evolution or the inclusion of effects of rotation. We have concentrated on current problems of stellar evolution theory since this may be the most efficient way to show its state of the art; it does by no means imply that it has not been also very successful in the recent past. However, it implies that still much work is to be done before massive star evolution models can be used as input for other astrophysical applications without asking for the remaining uncertainties in such models. It has been shown above, and by many contributions to this conference, that the circumstellar media can be considered as one of the best diagnostic tools to test and further improve stellar evolution models.

Acknowledgements This work has been supported by the Deutsche Forschungsgemeinschaft through grant No. La 587/8.

References Abbott D.C. (1982). Astrophys. J., 259, 282. Arnett W.D. (1978). Astrophys. J., 219, 1008. Arnett W.D., Bahcall J.N., Kirshner R.P. & Woosley S.E. (1989). Ann. Rev. Astron. Astrophys., 27, 629. Bertelli G., Bressan A.G. & Chiosi C. (1985). Astron. Astrophys., 150, 33. Blaha C. & Humphreys R.M. (1989). Astron J., 89, 1598. Bressan A., Fagotto F., Bertelli G. k Chiosi C. (1993). Astron. Astrophys. Suppl., 100, 647. Caughlan G.A. & Fowler W.A. (1988). Atomic Data and Nucl. Data Tables, 40, 238. Chaboyer B. & Zahn J.P. (1992). Astron. Astrophys.. 253, 173. Chiosi C. fc Summa C. (1970). Astrophys. Space Sci., 8, 478. Chiosi C. & Maeder A. (1986). Ann. Rev. Astron. Astrophys.. 24. 229. Chu Y.-H. (1991). in: IAU-Symp. 143, 349. Chu Y.-H., Treffers R.R. & Kwitter K.B. (1983). Aslrophys. J. Suppl.. 53, 937. Denissenkov P.A. (1993), Astron. Astrophys., (submitted). Diehl R., et al. (1993). Astron. Astrophys. Suppl., 97, 181. El Eid M.F. & Langer N. (1986). Astron. Astrophys., 1C7, 274. Fitzpatrick E.L. & Bohannan B. (1992). Astrophys. J., 404, 734. Fliegner J. &i Langer N. (1994), in: IAU-Symp. 1C2, (in press). Hamann W.-R., Koesterke L. & Wessolowski U. (1993), Astron. Astrophys., 274, 397. Herrero A., Kudritzki R.P., Vilchez J.M., Kunze D., Butler K. fc Haser S. (1992). Astron. Astrophys., 261, 209. Hillebrandt W. & Hoflich P. (1989). Rep. Prog. Physics, 52, 1421. Hoflich P., Langer N. L Duschinger M. (1993). Astron. Astrophys. Letter, 275, L29. Huang R.Q. & Weigert A. (1983). Astron. Astrophys., 127, 309. Iben I. Jr. &; Renzini A. (1983). Ann. Rev. Astron. Astrophys., 21, 271. Iglesias C.A., Rogers F.J. fc Wilson B.G. (1992). Astrophys. J., 397, 717.

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Iliadis Ch., et. al. (1990), Nucl. Phys. A, 512, 509. Kato S. (1966). P.A.S.J., 18, 374. Kudritzki R.P., Pauldrach A. k Puls J. (1987). Astron. Astrophys., 173, 293. Kudritzki R.P. k Hummer D.G. (1990). Ann. Rev. Astron. Astrophys., 28, 303. Lamers H.J.G.L.M. &; Leitherer C. (1993). Astrophys. J., 412, 771. Langer N. (1989), Astron. Astrophys. 220, 135. Langer N. (1989a), Astron. Astrophys. 210, 93. Langer N. (1991), Astron. Astrophys. 252, 669. Langer N. (1991a), Astron. Astrophys. 248, 531. Langer N. (1992), Astron. Astrophys. Letter 265, L17. Langer N. (1993), in: lAU-Colloq. 137, 426. Langer N. (1994). In: Evolution of Massive Stars: A Confrontation between Theory and Observations, eds. D. Vanbeveren et al., Space Set. Rev., (in press). Langer N., Sugimoto D. k Fricke K.J. (1983). Astron. Astrophys., 126, 207. Langer N., El Eid M.F. k Fricke K.J. (1985). Astron. Astrophys., 145, 179. Langer N., El Eid M.F. k Baraffe I. (1989). Astron. Astrophys. Letter, 224, L17. Leitherer C. k Langer N. (1990), in: lAU-Symp. 148, 480. Lennon D.J., Dufton P.L. k Fitzsimmons A. (1992). Astron. Astrophys. Suppl., 94, 569. Lennon D.J., Dufton P.L. k Fitzsimmons A. (1993). Astron. Astrophys. Suppl., 97, 559. Maeder A. (1981). Astron. Astrophys., 99, 97. Maeder A. (1987). Astron. Astrophys., 178, 159. Maeder A. (1992). Astron. Astrophys., 264, 105. Mahoney W.A., Ling J.C., Jacobson A.S. k Lingenfeltcr R.E. (1982). Astrophys. J., 262, 742. Matraka B., Wassermann C. k Weigert A. (1982). Astron. Astrophys., 107, 283. Meynet G., Maeder A., Schaller G., Schaerer D. fc Charboimel C. (1994). Astron. Astrophys. Suppl, 103, 97. Nomoto K. (1984), Astrophys. J., 277, 791. Prantzos N., Doom C , Arnould M. & de Loore C. (198(5), Astrophys. J., 304, 695. Schaller G., Schaerer D., Meynet G. k Maeder A. (1992). Astron. Astrophys. Suppl., 96, 269. Shull J.M. (1993). ASP Conf. Set:, 35, Cassinelli et al., eds., p. 327. Stencel R.E., Pesce J.E. k Bauer W.H. (1989). Astron. J., 97, 1120. Stothers R.B. (1991). Astrophijs. J., 383, 820. Stothers R.B. k Chin C.-w. (1975). Astrophys. J., 198, 407. Stothers R.B. k Chin C.-w. (1976). Astrophys. J., 204, 472. Stothers R.B. k Chin C.-w. (1991a). Astrophys. J., 374, 288. Sugimoto D. k Nomoto K. (1980). .Space Set. Rew., 25, 155. Tuchman Y. k Wheeler J.C. (1989a), Astrophys. J., 344, 835. Tuchman Y. k Wheeler J.C. (1989b), Astrophys. J., 346, 417. Woosley S.E. (1986). In: 16th Advanced Course of the Swiss Soc. of Astron. and Astrophys., Saas-Fee Lecture Notes, A. Maeder et al., eds., Geneva Observatory. Woosley S.E. k Weaver T.A. (1986). Ann. Rev. Astron. Astrophys., 24, 205. Woosley S.E., Langer N., k Weaver T.A. (1993). Astrophys. J., 411, 823. Woosley S.E., Langer N., k Weaver T.A. (1994). Astrophys. J., (in preparation). Zahn J.P. (1983). In: 13th Advanced Course of the Swiss Soc. of Astron. and Astrophys., Saas-Fee Lecture Notes, B. Hauck et al., eds., Geneva Observatory.

Evolution of AGB Stars with Mass Loss P. R. Wood Mount Stromlo and Siding Spring Observatories, Private Bag, Weston Creek PO, ACT 2611, Australia

Abstract Observational and theoretical estimates for mass loss rates from AGB stars are discussed. Then models for the evolution of AGB stars including mass loss and the effects of helium shell flashes are presented. Finally, the possibility of mass loss by binary mass transfer is discussed.

1 Introduction It is well established that the bulk of mass loss from low and intermediate mass stars occurs during the asymptotic giant, branch (AGB) stage of evolution, leading to the well-defined sequence of mass-losing stars in the IRAS two-colour diagram (van der Veen and Habing 1988) and the formation of planetary nebulae (Abell and Goldreich 1966; Renzini 1981). However, a reliable theoretical understanding of the causes of mass loss is still not available, although progress is being made. An additional complication is that the time history of mass loss during AGB evolution is quite complex since AGB evolution is modulated by helium shell flashes which control the surface luminosity and thereby the mass loss rate. In this paper, mass loss rates from AGB stars are discussed and the effects of helium shell flashes on the mass loss are described and compared with observations. Time dependent winds produced by AGB stars are reviewed. Finally, the evolution of AGB stars that lose mass in binary mass transfer events is briefly described. 2 A G B mass loss rates IRAS observations of stars in the solar neighbourhood indicate that those stars with substantial circumstellar shells - those with high mass loss rates - are nearly all AGB stars undergoing large-amplitude pulsation (Habing 1990). These stars are the long-period variables (LPVs) with pulsation periods > 1 year. This observational result is in agreement with theoretical calculations which indicate that high mass loss rates from AGB stars can only be produced by the combined effects of pulsation and radiation pressure on grains (Wood 1979; Bowen 1988; see Holzer and MacGregor 1985 for a review of possible mass loss mechanisms). However, the quantitative mass 15

16

P. R. Wood: Evolution of AGB stars

loss rates resulting from theoretical calculations are very uncertain as they depend sensitively on the details of grain formation and the rate of cooling of atmospheric material heated by the passage of the shock front that propagates through the stellar atmosphere once per pulsation cycle. Wood (1979) found that in the limit of no cooling, a. typical Mira would drive a mass loss rate of ~0.01 MQ yr" 1 ! In the opposite limit of instantaneous cooling, but with no dust present, M ~ 10~12 MQ yr" 1 . Allowing dust to form boosted the mass loss rate to M ~ 3 x 10~7 MQ yr" 1 which is perhaps 10 times smaller than the observed mass loss rates typical for the type of Mira star considered. Bowen (1988) introduced a cooling function whose dependence on T and p was adopted as physically plausible, but whose normalizing factor is very uncertain. The results were extended to a range of models of various masses and luminosities to see how mass loss varied with stellar parameters (see also Bowen and Willson 1991). The Bowen models produce quite extended hot post-shock regions whose pressure obviously plays a considerable role in driving the mass flow since Bowen (1988) and Bowen and Willson (1991) find that dust is not necessary to drive significant mass loss (although the dust enhances the mass loss rate). However, observations suggest that there is no hot region above the photosphere in Mira atmospheres. The inverted temperature profile of this region would produce molecular bands in emission (Bessell et al 1989), a. spectral feature that is never observed. Theoretical work on the production of large AGB mass loss rates is continuing. Feuchtinger et al (1993) have developed a code for studying dust formation and pulsation in Mira atmospheres; this code contains a more elaborate treatment of radiative transport than the codes of Wood or Bowen. Other current developments related to the theory of mass loss from AGB stars are further reviewed in these proceedings by Gauger et al and Cherchneff and Tielens. Another approach to estimating mass loss rates in AGB stars is to use observations. Over the last ~ 10 years, a. large number of mass loss rate estimates have been derived for AGB stars from CO microwave emission (e.g. Knapp and Morris 1985), and from infrared emission from dust in the cases of very high mass loss rates where the CO mass loss rate estimates become unreliable (Heske et al 1990). These mass loss rates are shown plotted against pulsation period (P) in Figure 1 (see Vassiliadis and Wood 1993 for the list of sources). Given that theory indicates that pulsation plays a prominent role in the production of the large mass loss rates observed in LPVs, it has been assumed that the mass loss rate is primarily a. function of pulsation period. The mass loss rate shown in Figure 1 shows a remarkably rapid increase

P. R. Wood: Evolution of AGB .stars

17

-3.0

1200 P(days)

1600

2000

2400

Fig. 1. Observationally determined mass loss rates in AGB stars plotted against pulsation period P (dots). The dashed line is the mass loss rate according to Reimers' law (with rj — 1/3), and the solid line is the radiation pressure driven mass loss rate M = L/cvexp, where an expansion velocity vexp = 10 km s" 1 has been assumed.

with P up to ~ 500 days (3 orders of magnitude increase in M for an increase in P of ~ 250 days) but for longer periods, the mass loss rate seems to be relatively constant at large values which have become synonomous with the term 'superwind' (Renzini 1981). Presumably, the transition between the two mass loss regimes corresponds to a change in the physical processes producing the wind, perhaps because the dust formation region moves in to the stellar photosphere rather than lying above it in the optically thin region which is supported and extended by pulsation (Bedijn 1988). A widely used formula, for computing mass loss rates in AGB (and first giant branch) stars is that of Reimers (1975). This formula (M = 4 x 1Q~13T]LR/M, with M in M 0 yr" 1 and L. R. and M in solar units) is also shown in Figure 1, where (Wood 1990) L is obtained from the (M(,o/, P) relation for Mira variables. R. is obtained from r^ejj on the giant branch and the definition I, = 47raR 2 T^ / , and M is obtained from the fit M = 0.013 p0.75 normalizing factor 7/ = 1/3 has been used in Figure 1 as this is the value required to give the amount of total mass loss of ~ 0.2 M,^ observed on the first giant branch of globular clusters (Wood and Calm 1976; Sweigert, Greggio and Renzini 1990).

18

P. R. Wood: Evolution of AGB stars

The most significant feature of Figure 1 is that the Reimers' mass loss formula produces a mass loss rate that increases much more slowly with P (or L) on the AGB than is indicated by modern observations. This means that if the AGB is to be terminated at luminosities that agree with observed maximum luminosities in Magellanic Cloud clusters, then the mass loss will occur over a longer time interval, and at lower mean luminosities, than occurs with the formula shown in Figure 1. This may have important implications for the formation of carbon stars, since the ease of dredge-up of 12C is enhanced by a large envelope mass (Wood 1981). Attempts to use Reimers' mass loss law with studies of dredge-up (Groenewegen and de Jong 1993) may result in the requirement for enhanced dredge-up efficiencies to compensate for the excess mass loss at low luminosities. Finally, it should be noted that it has long been realized that a Reimers' law is incapable of producing 'superwind' mass loss rates without an enormous enhancement factor t] (Renzini 1981). Although there is clearly a strong dependence of mass loss rate on P as shown in Figure 1, there are also other factors that influence the mass loss rate. A parameter that has been repeatedly shown to influence the mass loss rate of LPVs is the light curve shape, in particular, the fraction / of the period during which the light is rising. Those LPVs that have rapidly rising light curves always show a higher mass loss rate at a given P than the rate exhibited by LPVs with less rapidly rising light curves (Bowers and Kerr 1977; Onaka, de Jong and Willems 1989; Wood 1990; Jura and Kleinmann 1992). Just which intrinsic property causes the light curve to rise more rapidly in some LPVs than others is at present unknown, but a likely candidate is stellar (envelope) mass. Another factor that clearly influences the mass loss rate is the pulsation amplitude (in K, which is close to the bolometric light amplitude) (Whitelock 1990). Once again, at a given period, mass is likely to be the stellar property that influences the pulsation amplitude. The flat 'superwind' part of Figure 1 seems to correspond to stellar winds in which the momentum has been imparted by radiation pressure on dust grains. The ratio of wind momentum to photon momentum (3 = Mvexpc/L was shown by Knapp (1986) to exhibit a strong preference for a. value around 1.0 for Galactic AGB stars, although a few objects had ratios > 10. In the Galactic Bulge, presumably more metal-rich than the solar vicinity, Whitelock et al (1991) found 1/2 < (3 < 2 while Wood et at. (1992) found a similar result in the LMC where the metal abundance is ~ 1/2 solar. To a first approximation then, the mass loss rate in the superwind phase appears to

P. R. Wood: Evolution of AGB stars

19

be given by M = /3L/cvexp, where f3 = 1, although it should be remembered that values of /? > 10 may be possible (Netzer and Elitzur 1993). A final property of AGB winds that needs to be examined is the expansion velocity vexp. The expansion velocity of AGB winds can be readily obtained from wind spectral features such as the width of the microwave CO lines (which originate at radii of a few times 10 17 cm) or the separation of the twin 1612 MHz OH maser lines (which originate at radii of a few times 10 16 cm). When plotted as a function of pulsation period (Sivagnanam et al 1989), the expansion velocities show quite different behaviours for P < 500 days and P > 500 days. For P < 500 d, the expansion velocity increases rapidly with P from a few km s" 1 at 300 days to ~ 15 km s" 1 at 500 days. For P > 500 d, the expansion velocity remains approximately constant at vexp ~ 15 km s" 1 . This behaviour mirrors that of the mass loss rate itself (which increases exponentially to P ~ 500 d and remains roughly constant thereafter), adding further weight to the suggestion that the mass loss mechanism changes as the pulsation period increases through 500 days. Finally, it is worth noting that abundance as well as pulsation period appears to affect the wind expansion velocity. This is shown clearly by the OH/IR stars in the LMC (Wood et al 1992) which have mean expansion velocities of ~ 9 km s" 1 compared to ~ 15 km s" 1 for comparable Galactic OH/IR stars. The most likely source of this effect is the lower abundance in the LMC (~ 1/2 solar). It is interesting to note that if the LMC OH/IR stars have the same ratio of (3 = Mvexpc/L as the Galactic stars, then their mass loss rates will be higher for similar luminosities during the superwind phase. On the other hand, higher luminosities might be required in order to produce a superwind in lower metallicity stars.

3 AGB evolution with empirical mass loss rates Using the mass loss rates shown in Figure 1, Vassiliadis (1992) and Vassiliadis and Wood (1993) have computed the evolution of AGB stars including in full the effects of the luminosity variations associated with helium shell flashes. A typical example of the results obtained is shown in Figure 2. The evolutionary sequence shown begins at the first shell flash; it is found that mass loss up to this time is always insignificant. Indeed, the main feature of the mass loss rates resulting from the empirical formula shown in Figure 1 is that mass loss is insignificant except for the last few shell shell flash cycles, and within those cycles the mass loss is concentrated to the high luminosity quiescent phases which precede each flash. The high mass loss (superwind) phases are clearly discrete, and the potential for multiple shell ejections is

20

P. R. Wood: Evolution of AGB stars

"3 QO O

3.6 3.5

3.4 4 >

bo

o

I

I I

I

I I

H—I—h

H—I

-HH

3.5 3 H

h

I—I—I

1000 500

0 15

I—I—I—|—I—I—I—|—I—I—I—|—I—I

10 5 I

1

i

i

=+=+

0.8 0.6

H—i—i—|—i—i—i—|—^

8 6 4 _!__!

1

I

2

I i I

'

>

'

1J L

4 6 Time (10 5 years)

8

Fig. 2. The effective temperature T e /y, luminosity L. pulsation period P (days), expansion velocity v,,rp (km s" 1 ), mass M and mass loss rate M (lO-"'3/V/0 yr" 1 ) plotted against, time on the AGB for a star of mass 1 M 0 , helium abundance Y = 0.25 and metal abundance Z=0.00'4.

clear, although the time interval between superwind phases is large, ~ 105 years'for stars of initial mass ~ 1M 0 and decreasing only slowly with initial mass to ~ 7 x l O 4 yr at 2.5 MQ. The reason for the very rapid onset of the superwind phase is largely the exponential increase in A/ with period shown in Figure 1. However, as can be seen in Figure 2, the pulsation period (taken to be the fundamental mode in these calculations) varies enormously during one shell flash cycle

P. R. Wood: Evolution of AGB stars

21

from ~ 150 to ~650 days (taking the last flash cycle as an example). This is because P a R2 <x L/T^JJ and Teff decreases while L increases towards quiescent maximum. Note that optically-visible Mira variables in the solar vicinity mostly have periods in the range ~ 250 to ~ 475 days (Wood and Cahn 1977) so that a given AGB star transits the complete range of Mira periods in one shell flash cycle . It should be noted that the periods shown in Figure 2 are those that the star would have if it were pulsating, but extant pulsation calculations are not capable of determining the stability of AGB stars with enough reliability to say whether the stars will actually be pulsating or non-variable. The superwind mass loss rate used by Vassiliadis and Wood (1993) assumed /? = 1. However, as noted above, (3 is observationally uncertain by a factor of ~ 2 and values of up to 10 are plausible theoretically. If (3 = 2 had been chosen for the AGB calculations, the pre-superwind mass loss rate would be unaffected but the mass loss rate would rise to higher values so that fewer flash cycles would be needed to remove the hydrogen-rich envelope and turn the AGB star into a planetary nebula nucleus. For stars of initial mass ~ 1 M Q , one superwind phase could potentially completely remove the envelope, the actual behaviour depending on the phase of the flash cycle at which the star first reached the superwind limit. In general, however, it would be difficult to avoid at least two superwind phases in most stars, meaning that towards the end of their lives, most ACiB stars should have at least one episode where they regress from being dusty objects in a superwind phase to ordinary optically-visible AGB stars. Many such AGB stars have now been found using the IRAS data base (Willems and de Jong 1986; Zijlstra et al 1992), the signature of these objects being a high 60//, to 25/i flux ratio resulting from the cool dust in the hollow shell formed by the cessation of mass loss. The results of Zijlstra et al (1992) show that the hollow shells are much more common among carbon stars than M stars. The calculations of Chan and Kwok (1988) show that the hollow shells should be observable in the IRAS 2-colour diagram for ~ 104 years. Given an interflash lifetime of ~ 105 years, with ~ 1/2 of that time being a stage when the superwind is not blowing, it is therefore reasonable to expect ~ 20% of optically-visible C stars to show evidence for a hollow shell. Zijlstra. et al (1992) give an estimate of ~ 40% of optically-visible C stars showing evidence for hollow shells, in reasonable agreement with theoretical expectation. It is well known that at helium shell flashes, there is a brief (~600 yr) spike in luminosity which has a peak a factor of 1.5 or more brighter than the quiescent luminosity maximum (see second panel of Fig. 2). This luminosity peak also gives a spike in the mass loss rate (bottom panel of Fig.2 ) but,

22

P. R. Wood: Evolution of AGB stars

due to its short duration, the amount of mass lost in the spike is small, being < O.OlM© in the calculations of Vassiliadis and Wood (1993). Olofsson et al (1992) (see also Olofsson, these proceedings) have found a small number of C stars in which a hollow mass loss shell has been spatially resolved with the SEST. Furthermore, the mass loss shells in these stars appear to be thin with a mass of ~ 0.02-0.04 M® and Olofsson et al suggest that they are the result of mass loss occurring during the luminosity spikes associated with helium shell flashes. If this is the case, then the total mass lost is somewhat greater than predicted by the Vassiliadis and Wood calculations. Another place that the mass loss rate enhancement produced by a luminosity spike might be expected is in the Mira variables It Aql and R Hya. These two Miras are currently rapidly decreasing in pulsation period and are in the phase of luminosity decline from a helium shell flash spike (Wood and Zarro 1981). In the case of R Aql, the observed mass loss rate (~ 10~6 M 0 yr" 1 ) is not at superwind levels but it is still some two orders of magnitude higher than expected for a Mira its period of ~ 280 d (Wood 1990). On the other hand, R Hya has a smaller mass loss rate than expected (~ 10~7 M s yr" 1 at ~ 380 days). In neither of these two cases will a significant amount of mass be lost in the luminosity spike interval of < 1000 yr. Figure 2 indicates that this behaviour would be expected if these stars are not undergoing one of their last few shell flashes. The significance of the luminosity spikes as producers of mass loss may well be affected by the abundance of the star. This is shown in Figure 3, where the quiescent luminosity maximum LQ for each shell flash cycle is plotted against core mass Mc. Also plotted in Figure 3 is the peak luminosity Lp in each luminosity spike. Both luminosities increase linearly with Mc but the height of the spike above quiescent maximum is much greater in the low metallicity stars. This may mean that low metallicity stars (for example, population II stars, globular cluster stars, or more massive stars in low metallicity systems) will throw off more mass at shell flashes than more metal rich stars. The larger the mass of an AGB star through the low mass range up to ~ 2.5 M 0 , the greater is the number of shell flash cycles required to remove the envelope in superwind phases (being ~ 3 flashes at 2.5 MQ in the Vassiliadis and Wood 1993 calculations). For even more massive AGB stars (>3.5 M 0 ), mass loss appears to behave quite differently because in these stars the external luminosity is almost unaffected by the shell flashes (except for the first few in each star). While mass loss remains insignificant in these massive AGB stars, the star slowly increases in luminosity. When the star reaches luminosities high enough to produce a superwind, the stellar mass

23

P. R. Wood: Evolution of AGB stars 1

1

0

/

3x10*

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y o x

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in

7~ L P

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/

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Of <:

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_ (M : Y; Z) <• ( 20.4-0; 0.3,0.02) * (1 75,20 iO-3iO.OO1) ' (1- 25,0.2 ; 0-001) » (1 0 ; 0 . 2 ; 0-0001)

0

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Fig. 3. Quiescent luminosity maximum LQ and peak surface luminosity at each shell flash LP plotted against core mass M c . The symbols are coded according to stellar mass M, helium abundance Y and metal abundance Z. Solid symbols indicate shell flashes where dredge-up of 1 2 C occurred.

24

P. R. Wood: Evolution of AGB stars

starts declining quickly and the stellar luminosity declines also (Vassiliadis and Wood 1993) i.e. the star evolves down the AGB. When the envelope mass is reduced to ~ 1 . 5 M©, the AGB behaviour changes to mimic that of lower mass stars, with large luminosity spikes appearing at the stellar surface at each shell flash. All throughout these final shell flashes, however, the star remains in the superwind phase and is presumably seen only as a dust-enshrouded object.

4 Time-dependent AGB winds Many authors have made steady-state models for AGB winds (e.g. see articles by Gauger et a/and by Cherchneff and Tielens in this volume). But given that AGB stars undergo large excursions in luminosity associated with helium shell flashes, one might ask how the winds will be affected by this behaviour. Vassiliadis and Wood (1992) computed simple hydrodynamic models of AGB winds in an attempt to assess the effect of these helium shell flash luminosity variations, taking the mass loss rate given by the relation in Figure 1 as the input at the inner boundary of the flow. As expected, the shell flash cycle time of ~ 105 years leads to the formation of hollow shells due to the cessation of the superwind following each shell flash. The calculations also reveal the possibility that during the extended luminosity dip following each shell flash, matter can actually fall back onto the central star. Whether this occurs in practice will depend on the opacity of the dust/gas mix in the wind. The outward force exerted on the circumstellar material is (o l)GM/R2, where a — nL/(4ncGM). Clearly, a must be greater than unity for outward mass flow, while the observational requirement that vexp < 15 km s" 1 demands a-1 be of order unity or less when the star is at quiescent luminosity maximum. Since L declines by a factor of 2-2.5 in the post-flash dip, a could easily decline to a value < 1 at this point, leading to material falling back onto the star. A final feature revealed by the time-dependent flow calculations is the occurrence of a faster wind generated by the luminosity spike of helium shell flashes (vexp ~ 20 km s" 1 , compared to ~ 1 5 km s" 1 at quiescent luminosity maximum). Because of the short duration of this pulse, there is no long term affect on shell structure. However, there are two OH/IR stars (IRAS 18520+0533 in Eder, Lewis and Terzian 1988, and IRAS 172532824 in te Lintel Hekkert 1990) which show twin pairs of OH maser peaks that would be expected for a wind flow containing an inner faster wind and an outer slower wind. The fraction of OH/IR stars showing twin flow

P. R. Wood: Evolution of AGB stars

25

velocities (2 in several hundred) is roughly the number expected given the short duration of the luminosity spike.

5 AGB evolution and binary mass transfer As noted in the last section, superwinds operating in single AGB stars remove the stellar envelope when the star is well into the helium shell flashing phase. The remnant star then has little nuclear fuel left to burn and it evolves through the region of the planetary nebula nuclei (PNNi) on timescales of order 103 to a few times 104 years, exciting the AGB wind material with UV photons as it evolves at high effective temperatures and producing a planetary nebula. If a star could be made to lose its hydrogenrich envelope during the early-AGB phase when the stellar luminosity is provided by helium shell burning following core helium exhaustion, the star would still move to high effective temperatures but it would still have plenty of nuclear fuel left to burn (helium in this case), leading to a long-lived object. Binary mass transfer is a mechanism for removing the H-rich envelope of early-AGB stars, although this mechanism can only apply to intermediatemass stars of initial mass >2.3 MQ which do not develop degenerate helium cores before helium core ignition. The low mass (M < 2.3 MQ) stars evolve to luminosities and radii on the first giant branch (FGB) which are larger than those during the early-AGB phase, so that for these stars the mass transfer would occur during the FGB rather than the early-AGB. Iben and Tutukov (1993) have examined in detail the evolution of intermediate mass stars undergoing early-AGB mass transfer (see also the contribution by Han et al in these proceedings), and find that this type of mass transfer will occur for semi-major axes A of 20 < A / R 0 < 1000. For larger orbital separations, mass transfer will occur during the thermally pulsing AGB, leading to evolution of the remnant star which is similar to that which would have occurred in a single star when a, superwind developed, although the ejected matter would presumably be spatially distributed in a quite different manner (see articles by Livio and Icke in this volume). Iben and Tutukov (1993) estimate that all the known short-period binary PNNi are the result of early-AGB mass transfer. The implications of this finding are that these nuclei will live much longer than will PNNi resulting from mass loss during the later thermally pulsing AGB phase, and the (luminosity, core mass) relation will be different from the usual one assumed for AGB stars (see Iben and Tutukov 1993).

26

P. R. Wood: Evolution of AGB stars

References Abell, G.O. and Goldreich, P. 1966, PASP, 78, 232 Bedijn, P.J. 1988, AfcA, 205, 105 Abell, G.O. and Goldreich, P. 1966, PASP, 78, 232 Bessell, M.S., Brett, J.M., Scholz, M. and Wood, P.R. 1989, A&A, 213, 209 Bowen. G.H. 1988, ApJ, 329, 299 Bowen, G.H. and Willson, L.A. 1991, ApJ, 375, L53 Bowers, P.F. and Kerr, F.J. 1977, A&A, 57, 115 Chan, S.J. and Kwok, S. 1988, ApJ, 334, 362 Eder, J., Lewis, B.M. and Terzian, Y. 1988, ApJS, 66, 183 Feuchtinger, M.U., Dorfi, E.A. and Hofner, S. 1993, A&A, 273, 513 Groenewegen, M.A.T. and de Jong, T.1993, A&A, 267, 410 Habing, H.J. 1990, in From Miras to Planetary Nebulae: Which path stellar evolution, eds. M.O. Mennessier and A. Omont (Editions Frontieres: Gif sur Yvett), p.16 Heske, A., Forveille, T., Omont, A., van der Veen, W.E.C.J. and Habing, H. 1990, A&A, 239, 173 Holzer, T.E. and MacGregor, K.B. 1985, in Mass Loss from Red Giants, eds. M. Morris and B. Zuckerman (Reidel: Dordrecht), p.229 Iben, I. and Tutukov, A.V. 1993, ApJ, in press Jura and Kleinmann 1992, ApJS, 79, 105 Knapp, G.R. and Morris, M. 1985, ApJ, 292, 640 Knapp, G.R. 1986, ApJ, 311, 731 Netzer, N. and Elitzur, M. 1993, ApJ, 410, 701 Onaka, T., de Jong, T. and Willems, F.J. 1989, A&A, 218, 169 Olofsson, H., Carlstrom, E., Eriksson, K and Gustaffson, B. 1992, AfcA, 254, L17 Reimers, D. 1975, in Problems in Stellar Atmospheres and Envelopes, eds. B. Bascheck, W.H. Kegel and G. Traving (Springer: Berlin), p.229 Renzini, A. 1981, in Physical Processes in Red Giants, eds. I. Iben and A. Renzini (Reidel: Dordrecht), p.427 Sivagnanam, P., Le Squeren, A.M., Foy, F., and Tran Minh, F. 1989, A&A, 211, 341 Sweigert, A.V., Greggio, L. and Renzini, A. 1990, ApJS, 69, 911 te Lintel Hekkert, P. 1990, PhD thesis, Leiden van der Veen, W.E.C.J. and Habing, H. 1988, A&A, 195, 125 Vassiliadis, E. 1992, PhD Thesis, Australian National University. Vassiliadis, E. and Wood, P.R. 1992, Proc. Astr. Soc. Australia, 10, 30 Vassiliadis, E. and Wood, P.R. 1993, ApJ, 413, 641 Whitelock, P.A. 1990, ASP Conf. Series, 11, 365 Whitelock, P.A., Feast, M.W. and Catchpole, R. 1991, MNRAS, 248, 276 Willems, F.J. and de Jong, T. 1986, ApJ, 309, L39 Wood, P.R. 1979, ApJ, 227, 220 Wood, P.R. 1981, in Physical Processes in Red Giants, eds. 1. Iben and A. Renzini (Reidel: Dordrecht), p.135 Wood, P.R. 1990, in From Miras to Planetary Nebulae: Which path stellar evolution, eds. M.O. Mennessier and A. Omont (Editions Frontieres: Gif sur Yvett), p.67 Wood, P.R. and Cahn, J.H. 1977, ApJ, 211, 499 Wood, P.R. and Zarro, D.M. 1981, ApJ, 247, 247 Wood, P.R., Whiteoak, J.B., Hughes, S.M.G., Bessell. M.S., Gardner, F.F. and Hyland, A.R. 1992, ApJ, 397, 552 Zijlstra, A.A., Loup, C , Waters, L.B.F.M. and de Jong, T. 1992, A&A, 265, L5

The physical theory of winds from hot stars J. E. Drew Department of Physics, Nuclear Physics Laboratory, Keble Road, Oxford 0X1 3RH, U.K.

Abstract Radiation pressure driven wind theory as applied to OB and related stars is reviewed, beginning with the first detailed formulation of the theory by Castor, Abbott & Klein (1975). The character of the line acceleration term in the equation of motion is discussed. The main successes of the time-independent theory are noted, along with its failures which motivated the more recent development of time-dependent (shockedwind) theory.

1 Introduction An eaxly result from ultraviolet astronomy was that OB stars with bolometric magnitudes brighter than Mboi — - 6 suffer significant mass loss (Snow & Morton 1976). At about the same time the framework was laid down by Castor, Abbott & Klein (1975; hereafter CAK) for what, has proved since to be a very successful theory of mass loss from OB and other similarly highluminosity stars. In bare outline the physical model is a simple one in which the outward force able to overcome gravity is the pressure exerted by the hot star's radiation field on its own atmosphere. In practical application, complexity arises from the fact that it is overwhelmingly the scattering of radiation in spectral lines that, mediates the force (a point first appreciated by Lucy & Solomon 1970). Since its inception, the radiation-driven wind model has undergone some quantitatively significant refinements and has begun to be applied in recent years to the central stars of planetary nebulae. Furthermore, it is now set to overcome what for a decade or so has seemed a serious objection to its application to Wolf-Rayet stars. No major challenge to this model has yet been mounted, although its perceived difficulty with Wolf-Rayet winds has left room for suggestions that pulsation may lead to significant mass loss (see the article by Langer in this volume) or that, in such cases, radiation pressure may need the assistance of rapid rotation acting in tandem with a stellar magnetic field (Poe, Friend & Cassinelli 1989). In this brief review, the emphasis is upon the physical character of the line-driving force, its adaptive properties, and how it leads to predictions for the key observables, M and the terminal velocity w,x,. Successes and failures 27

28

J. E. Drew: Hot star winds

of the theory are then discussed and used to motivate its extension to allow for the inherent instability of radiation-driven winds. It shall be seen that while the instability and consequent wind-shocking mostly do not threaten the M predictions of modified CAK theory, they do provide an explanation for the otherwise surprisingly wide range of excitation apparent in 0 star spectra and for the quite abrupt cessation of mass loss at Mboi ~ —6.

2 The line-driving force The equation of motion describing radiation-driven outflow may be written: dv dr

I dp p dr

GM rz

. .

The last term on the right hand side is the radiation pressure term. It consists of the sum of the radiative acceleration due to electron scattering (
• are themselves functions In principle the line-force parameters, k, a and <> of radius and velocity - as are t (a measure of optical depth involving the reciprocal of the velocity gradient), ne (election density), W (the dilution factor) and CF (the so-called finite disk correction). In practice, the iteration strictly required between the calculated wind kinematics and excitation, which together determine the line force parameters, is usually avoided by adopting fixed values judged to be representative. Generally speaking k, which is related to the effective number of lines contributing to the driving, and Q, an index defining the mix of optically-thick and optically-thin lines, are the more important parameters. Smaller a implies a larger contribution from optically-thin transitions. Note that, a is more likely to be over- rather than underestimated as it is more of a, problem to be sure of completeness with regard to optically-thin, as opposed to optically-thick, line transitions. Among the key findings of the original CAK study are the following relations concerning the mass loss rate and terminal wind speed: a and

J. E. Drew: Hoi star winds

•2.-(if;) 4.

29

«>

In the above, uth is a representative thermal speed for the outflow, T is the ratio between the stellar bolometric luminosity L and the Eddington luminosity and vesc is the photospheric escape velocity. Calculations of the line-force imply values of a in the range 0.5-0.75. The dominant term in equation (3) is the luminosity dependence. A success of the original CAK theory is that this predicted dependence is borne out by observation - studies of large samples of OB stars have shown that M varies roughly as Lie (e.g. Garmany & Conti 1984), implying a ~ 0.6. However the original theory gave mass loss rates that were rather too high and terminal velocities that were too low. These discrepancies were all but eliminated when the finite disk correction, an allowance for the finite solid angle presented by the star to the inner wind, was inserted into the analysis (Friend & Abbott 1986; Pauldrach, Puls & Kudritzki 1986 - see equation 2). The effects of the finite disk correction were to (i) bring the outflow critical radius closer to the stellar photosphere, (ii) weaken the radiative force near the star and so lower the mass loss rate by a factor of ~2 (iii) allow the unchanged radiative force at larger radii to accelerate the reduced density wind up to terminal velocities around 3 times higher than those derived by CAK. Much of the physical insight on the dependence of the force-multiplier upon stellar parameters and chemical composition is due to Abbott (1982). The line list compiled by Abbott at that time for use in his calculation of the force-multiplier has continued to be applied and modified by Kudritzki's group in Munich. Abbott showed that the line-driving force is 'remarkably constant over the (effective) temperature range 50000 > Teff > 10000 K'. This has much to do with the way the photospheric radiation field controls the wind ionization state - in effect the line-blocking due to a hot-star wind shifts longward or shortward in wavelength in step with the shift in the wavelength of peak photospheric emission. Encouraged by this, Abbott provided a global fit to the force multiplier in which k — 0.28, a = 0.56 and 6 — 0.09 (cf. the later recommendations of Pauldrach et al. 1986, Pauldrach et al. 1990 which notably tend tdwards higher values of a ) . As this amounts to a massive simplification of a complex problem, the inclination to use an 'off-the-peg' parameterisation of the force multiplier understandably remains strong. It is often the option chosen. To digress briefly, another adaptive pattern of behaviour we can expect of hot star winds is in their thermal balance. Again the root cause is the close link between the photospheric radiation field and the ionization and

30

J. E. Drew: Hot star winds

excitation of the radiation-driven wind. Time-dependent modelling has not changed the expectation that most of a 'typical' hot-star wind is in the form of cool gas maintaining radiative equilibrium. Equilibrium wind temperature profiles, Te(r), derived by Drew (1989) scale in a strikingly simple way as a function of stellar effective temperature (see Figure 1 in Drew 1990). In fact the following numerical fit:

1

(5)

reproduces the model grid of calculated O-star wind temperature profiles to within 10 percent (the grid encompasses dwarfs through to supergiants and a range of velocity laws). Gayley & Owocki (199-1) have examined and discount additional kinematic heating processes, excepting the ion-drag frictional heating discussed by Springmann & Pauldrach (1992) which only interferes substantively with radiative equilibrium in the lowest-density winds (e.g. in r Sco, BO V). In contrast to the effective temperature insensitivity of the force multiplier, there is a marked dependence on chemical composition (Abbott 1982 in the mean). This follows primarily from the obtained M oc (Z/ZQ)0A4 fact that H and He contribute very little line acceleration. So while it can be acceptable to transfer force multipliers derived for one star or group of stars to another maintaining constant abundances, it is dangerous to do this if there is a significant abundance change. Kudritzki et al. (1987) give representative values of k, a and S for LMC, SMC and Galactic abundances.

3 The Wolf-Rayet wind momentum problem — a solution A longstanding problem in the study of Wolf-Rayet mass loss is the apparent shortfall in the available supply of momentum from the i-a.dia.tion field: it has been found that empirical estimates of the wind momentum rate Mv^ are often an order of magnitude higher than the radiant momentum rate L\>o\lc (Barlow, Smith & Willis 1981; and more recently - Schmutz, Hamann & Wessolowski 1989). If radiation pressure is to drive such winds, every stellar photon has to undergo repeated scatterings so that its momentum can be exploited repeatedly. Is this possible? A recent paper by Lucy and Abbott (1993) suggests quite plausibly that it is. The reason lies in another example of adaptive behaviour, like the effective temperature insensitivity of the force multiplier noted above. In contrast to the winds of normal OB stars, Wolf-Rayet winds are stratified such that ionization decreases outward. This is a response to what may be thought

J. E, Drew: Hot star winds

31

of as the radial cooling of the ambient radiation field (there is no classical photosphere in a WR star). Lucy & Abbott demonstrate how the ionization gradient causes the line-blocking to shift longward in wavelength with increasing radius (see their Figure 5). This creates the conditions needed for multiple-scattering to work as required, since the effect of many scatterings is a progressive redshift. In short, the wind opacity moves to where the photons are to be found. The specific calculation performed by Lucy & Abbott (1993) is for a WN5 star for which L/iMv^c) ~ 10 is achieved.

4 Successes and failures of modified CAK theory Since the inclusion of the finite disk correction, the time-independent theory of radiation driven winds has done a good job of matching the observed run of mass loss rate against bolometric luminosity for normal OB stars. Predicted and observed mass loss rates vary from around 10~ 8 M© yr" 1 for the latest 0 dwarfs up to as much as 10" 5 M© yr" 1 for the earliest 0 supergiants. However, nothing is perfect and some dispute continues (e.g. the recent reassessment by Lamers & Leitherer 1993). The current situation is not quite as satisfactory when it comes to the comparison between measured and predicted terminal velocities. On the basis of a relatively modest sample of OB stars, Groenewegen, Lamers & Pauldrach (1989) showed that predicted wind terminal velocities are on average too high by around 40 percent. This picture was substantially confirmed by Prinja, Barlow & Howarth (1990) whose sample included more than 200 stars. What is wrong? An obvious candidate for the blame is the force multiplier. Perhaps the recent drift toward a ~ 0.7 is inappropriate - much better agreement would be achieved for a between 0.5 and 0.6 (see equation 3). Another factor that may be at work has to do with the influence of wind-shocking on the kinematic structure (see below). Another way of testing the validity of modified CAK theory is to compare the predicted wind ionization with that implied by ultraviolet observations of wind-formed resonance line profiles. Cassinelli & Olson (1979) were the first to point out that 0 VI AA1O32,38 absorption, found to be widespread in O-star spectra obtained by the Copernicus satellite, was too strong to be the product of O 4 + valence shell ionization - they proposed instead Auger ionization of the dominant O 3 + ion by a modest flux of X-rays. Soon after, soft X-ray emission from OB stars was detected by the Einstein satellite (Lx ~ 10~7i/bob s e e Cassinelli 1985). However, Cassinelli k Olson's conjecture was apparently contradicted by the modelling of Pauldrach (1987) wherein sufficient O 5 + was obtained just from valence shell ionization. Mak-

32

J. E. Drew: Hot star winds

ing similar assumptions about the photospheric radiation field and wind density profile, Drew's (1989) modelling supported Cassinelli & Olson's interpretation. Drew suggested that Pauldrach's models might overestimate the 0 5 + fraction because of the use of elevated wind temperatures. Since then, Groenewegen & Lamers (1990) and MacFarlane et al. (1993) have identified Pauldrach's omission of dielectronic recombination from his models as a second factor leading to too high a predicted O 5 + abundance. The work of MacFarlane et al. shows for the case of £ Pup (04 If) that mixing X-ray emission of the magnitude observed inward to at least r ~ 2i?» solves the 0 vi problem. The observed strength of the Si iv A1397 P-Cygni feature also poses a problem - Si 3+ is easily destroyed by photoionization in 0 star winds and so the Si IV line should not be as prominent as it is observed to be. There are several examples in the literature of fits to UV spectra of hot stars that notably fail to match this line (e.g. Kudritzki et al. 1991). Drew (1989) suggested that, just as the wind-shocking invoked to explain the soft X-ray emission could solve the 0 vi problem, it could also make good the Si iv deficit. This is because wind-shocking provides both the source of X-rays needed for 0 5 + production and compressed zones within which the Si 4+ recombination rate can be enhanced by one or two orders of magnitude. However it should be noted that CAK theory does explain the trend that gives rise to the 'Si IV luminosity effect" (Walborn & Panek 1984: discussed by Drew 1989 and Pauldrach et al. 1990). It became apparent through the 1980s that line-driving is naturally unstable and that instabilities may amplify, eventually steepening into full-blown shocks (see Owocki, Castor & Rybicki 1988 and references therein). Thus the basic radiation-driven wind concept contains within it the physics required to explain both the observed X-rays and and also a range of otherwise troublesome spectroscopic phenomena, including those just described.

5 Wind-shocking and time-dependent theory In performing the first 1-D hydrodynamical simulation of the growth of instabilities in a radiation-driven wind, Owocki, Castor & Rybicki (1988) found that the resultant wind structure was one swept by a sequence of reverse shocks. In other words, the shocks were found to expand outward more slowly than the host fluid - the outflow stumbles over itself. This carries the implication that the thin post-shock shells of cool compressed gas, which will contain most of the integrated wind column density, travel at speeds below those predicted by the time-independent theory (see Fig-

J. E. Drew: Hot star winds

33

ure 1 of Puls, Owocki & Fullerton 1993). Simulation of unstable outflow is not yet advanced enough to reliably quantify this effect, but it may be a partial explanation for the remaining discrepancy between observed and predicted terminal velocities (discussed in section 3). So far no two- or three-dimensional simulations have been published, nor have the 1-D examples been continued on to the large radii sampled by radio emission. The details of the early scenario for wind-shocking proposed by Lucy (1982) have not been borne out by either the recent hydrodynamical simulations or by the shock temperatures deduced from fits to the observed X-ray emission (the greater sensitivity of ROSAT to very soft X-ray emission has not changed this: e.g. Cassinelli et al. 1991). Lucy envisaged a tightly-spaced train of soft shocks within hot star winds, radiating at a few times 105 K. Simulation and observation favour fewer, harder and more widely-spaced shocks radiating at a few times 10(:> K. However, Lucy (1984) has argued compellingly that the typically high optical depth presented by the inner wind to scattered line radiation must defer shock growth to larger radii. Hence, for more luminous stars with substantive winds, acceleration through the outflow critical point (whose properties fix M) will be untroubled by the kinematic and radiative effects of shocking of the outer wind. However, at around Snow & Morton's (1976) bolometric magnitude limit of ~ —6, the line-driving instability causes a qualitative change in the wind physics. For Mboi ~ - 6 , modified CAR" theory would suggest M ~ 10~ 8 M 0 yr" 1 . For mass loss rates of this order, the flow timescale begins to drop below the typical post-shock cooling time. Once heated to T^ilO 6 K by shocking, such winds stay hot and radiate accordingly. An outline discussion of the physical effects involved may be found in Castor (1987). We may picture lower-luminosity OB stars near the Snow & Morton limit as enclosed within hot tenuous envelopes of soft X-ray emitting gas created by windshocking. This inability to recover from shocking can explain the anomalously low terminal velocities associated with late-0 main sequence and lower luminosity stars (Abbott & Friend 1989). Their winds, once shocked, are too highly-ionized to yield the ultraviolet opacity needed for continued acceleration. Relatively extreme cases have now been identified in which windshocking occurs near enough to the critical radius to significantly irradiate it and so lower the mass loss rate: for the B stars, ft and e CMa, the derived mass loss rates are ~5 times below the predictions of modified CAK theory (Drew et al. 1994). This is apparently how radiation driving turns off at the lower bolometric luminosity limit - not gradually by simple extrapolation of the Lbol - M relation, but more suddenly as the consequences of the

34

J. E. Drew: Hot star winds

line-driving instability overwhelm. Qualitatively this amounts to another success of a very fruitful wind theory. Acknowledgements JED is presently in receipt of an Advanced Fellowship funded by the Science & Engineering Research Council of the United Kingdom. References Abbott, D. C , 1982, Astrophys. J., 259, 282 Abbott, M. J. k Friend, D. B., 1989, Astrophys. J., 345, 505 Barlow, M. J., Smith, L. J. k Willis, A. .]., 1981, Mon. Not. R. ash: Soc, 196, 101 Cassinelli, J. P., 1985, in The origin of nonradiative healing/momentum in hot stars, eds. A. Underhill k A. Miclialitsianos, NASA Conf.Publ. 2:158, 2 Cassinelli, J. P., Cohen, D. H., MacFarlane, J. J., Sanders. VV. T. k Welsh, B. Y., 1994, Astrophys. J., 421, 705. Cassinelli, J. P. k Olson G. L., 1979, Astrophys. J., 229, 304. Castor, J. I., Abbott, D. C. k Klein, R. I., 1975, Astrophys. J., 195, 157 (CAK) Castor, J. I., 1987, in Instabilities in Luminous Early-Type Stars, eds. H.J.G.L.M. Lamers k C.W.H. de Loore, D. Reidel, The Netherlands, pl59 Drew, J. E., 1989, Astrophys. J. Suppl., 71, 267 Drew, J. E., 1990, in Properties of hot luminous stars, ed. CD. Gannany, ASP Conf. Series, 7, 230 Drew, J. E., Denby, M. k Hoare, M. G., 1994, Mon. Not. R. ash: Soc, 266, 917. Friend, D. B. k Abbott, D. C , 1986, Astrophys. J., 311, 701 Garmany, C. D. fcConti, P. S., 1984, Astrophys. J., 284, 705 Gayley, K. G. k Owocki, S. P., 1993, Astrophys. J.. (in press). Groenewegen, M. A. T. k Lamers, H. J. G. L. M., 1990, Astron. Astrophys., 243, 429 Groenewegen, M. A. T., Lamers, H. J. G. L. M. k Panldrach, A. W. A., 1989, Astron. Astrophys., 221, 78 Kudritzki, R. P., et al., 1991, in Massive stars in starbursts, eds. C. Leitherer, N. R. Walborn, T. M. Heckman k C. A. Norman, CUP, Cambridge, p59 Kudritzki, R. P., Pauldrach, A. W. A. k Pnls, J., 1987, Astron. Astrophys., 173, 293 Lamers, H. J. G. L. M. k Leitherer, C , 1993, Astrophys. J., 412, 771. Lucy, L. B., 1982, Astrophys. J.,255, 286 Lucy, L. B., 1984, Astrophys. J., 284, 351 Lucy, L. B. k Abbott, D. C , 1993, Astrophys. J., 405, 738 Lucy, L. B. k Solomon, P. M., 1970, Astrophys. J.. 159, 879 MacFarlane, J. J., Waldron, W. L., Corcoran, M. F., Wolff". M. .1., Wang, P. k Cassinelli, J. P., Astrophys. J., 419, 813. Owocki, S. P., Castor, J. I. k Rybicki, G. B., 1988. Astrophys. J.. 335, 914 Pauldrach, A. W. A., 1987, Astron. Astrophys., 183, 295 Pauldrach, A. W. A., Kudritzki, R. P., Puts, .]. k Butler, K., 1990, Astron. Astrophys., 228, 125 Pauldrach, A. W. A., Puls, J. k Kudritzki, R. P., 1986, Astron. Astrophys., 164, 86 Poe, C. H., Friend, D. B. k Cassinelli, J. P., Astrophys. J., 337, 888 Prinja, R. K., Barlow, M. J. k Howarth, I. D., 1990, Astrophys. J., 361, 607 Puls, J., Owocki, S. P. k Fullerton, A. W., 1993, Astron. Astrophys., 279, 547. Schmutz, W., Hamann, W.-R. k Wessolowski, U., 1989, Astron. Astrophys., 210, 236 Snow, T. P. k Morton, D. C , 1976, Astrophys. J. Suppl., 32, 429 Springmann, U. k Pauldrach, A., 1992, Astron. Astrophys., 262, 515 Walborn, N. R. k Panek, R. J., 1984, Astrophys. J., 280, L27

Axisymmetric Outflows from Single and Binary Stars Mario Livio Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA and Dept. of Physics, Technion, Haifa 32000, Israel

Abstract Mechanisms that can produce axisymmetric planetary nebulae are reviewed. It is suggested that the "interacting winds" model, in the presence of a "density contrast" between the equatorial and polar directions, can produce most of the observed morphologies. Mechanisms that can produce a density contrast are examined and it is suggested that binary stellar companions and even brown dwarfs and massive planets may play an important role in the generation of such a contrast, either via common envelope evolution or by spinningup the envelope of the AGB star. It is shown that the statistics of planetary nebulae of different morphological types are consistent with models which rely on the presence of binary companions.

1 Introduction Axisymmetric outflows are associated with many nebulae (e.g. He 2-36, BI Cru, My Cn 18, IC 4406, K 3-72, Corradi k Schwarz 1993a, b, c; OH 17.72.0, La Bertre 1986; R Aquarii, Burgarella k Paresce 1992) and with Be stars. I will concentrate in the present review mainly on planetary nebulae (PNe). An examination of the catalogue of narrow band images of Schwarz, Corradi and Melnick (1992) and other images in the literature reveals a few interesting morphologies. In some cases, almost perfect rings are observed (e.g. ScWe 3, ScWe 2, Schwarz, Corradi k Melnick 1992; Hen 1357, Bobrowsky 1993). In others, a pronounced bipolar structure is extending from a ring (e.g. He 2-104, Schwarz et al, 1992; NGC 2346, Bond k Livio 1990). Some PNe exhibit an exact point symmetry (e.g. IC 4634, J 320, Schwarz et al. 1992). Finally, in some PNe highly col Unrated outflows ("jets") are observed (e.g. NGC 7009, Schwarz et al. 1992; Kl-2, Bond k Livio 1990). In the present work, I examine different potential mechanisms for the formation of axisymmetric outflows and the different morphologies that are observed. 35

36

M. Livio: Axisymmeiric

outflows from single and binary stars

2 The Shaping of Planetary Nebulae PNe are very probably bubbles, formed when a tenuous but fast (V ~ 1000 km s" 1 ) wind, emitted by the exposed hot nucleus, catches up with the slowly moving (V ~ 20 km s" 1 ) wind, ejected by the AGB star, and shocks it. This "interacting winds" model has been suggested for the shaping of PNe around single stars (Kwok 1982; Kahn 1982). Balick (1987) proposed a classification scheme, in which all PNe are divided into "spherical," "elliptical" and "butterfly," with the different morphology classes being the consequence of an increasing "density contrast" (in the slow wind) between the equatorial and polar directions. Balick suggested that "spherical" PNe form when there is no "contrast" in the AGB wind, while "butterfly" PNe ensue when material near the equator is considerably denser than material near the poles (thus making it much easier for the fast wind to penetrate in the polar direction). The exploratory work of Soker and Livio (1989) and the more recent, highresolution gas dynamics simulations of Icke, Balick and Frank (1992), Frank et al. (1993; and see Icke, these proceedings) have demonstrated that bipolar flows are indeed obtained by the "interacting winds" model, in the presence of a "density contrast." Furthermore, the work of Frank et al. (1993) has shown that a very wide range of PNe morphologies, ionization structures and kinematics can be reproduced by this model, when the inclination angle at which the nebula is viewed is taken into consideration. The main questions that need to be answered are therefore: (i) what is the mechanism that produces the density contrast in the slow wind? (assuming that this represents the correct model) or (ii) are there any other shaping mechanisms that can produce the observed axially symmetric morphologies?

3 Mechanisms that Can Produce a Density Contrast or Axially Symmetric Outflows I will now review different mechanisms that can, in principle at least, generate a density contrast or produce by themselves an axially symmetric outflow. Some of these mechanisms involve single stars and some binary systems.

3.1 Common Envelope Evolution The term common, envelope (CE) is usually used to describe a configuration in which a binary system (typically consisting of the core of a giant and a

M. Livio: Axisymmetric outflows from single and binary stars

37

secondary star) revolves inside an envelope which is not corotating with the binary and not even necessarily in hydrostatic equilibrium. The development of a CE usually accompanies the event of mass being transferred from one star to the other on a dynamical timescale (see e.g. Iben & Livio 1993 for a recent review), this being the consequence of the mass losing star being unable to contract as rapidly as its Roche lobe. Such mass transfer events can occur when mass is being transferred from a star which possesses a deep convective envelope (such as an AGB star; in which case it tends to expand upon mass loss) onto a less massive companion (this leads to a shrinkage of the orbit). A dynamical mass transfer event may ensue also as a consequence of the instability discussed by Lai, Rasio & Shapiro (1993a, b; this is probably most applicable for double degenerate systems). Since the timescale of mass transfer in these cases is much shorter than the timescale on which the secondary star is able to adjust thermally, the outer layers of the secondary start expanding, and the system rapidly evolves to a configuration in which the core of the mass losing star and the original secondary are embedded inside a common envelope. Due to frictional drag, the internal binary starts spiralling-in (e.g. Taam and Bodenheimer 1989, 1991; Livio and Soker 1988). The final outcome of the CE phase can be a dramatic reduction in the binary separation (or even a merger of the two components in some cases) accompanied by the ejection of the envelope due to the deposition of orbital energy (other sources of energy may contribute, see Iben & Livio (1993)). This course of CE evolution is responsible for the formation of PNe with close binary central stars (Bond & Livio 1990; Livio 1993; Yungelson et al. 1993), cataclysmic variables (Paczynski 1976), double degenerate systems (Iben & Tutukov 1984; Webbink 1984) and in fact all the binaries containing at least one compact component, with an orbital period of less than a few days. One of the most important results of multi-dimensional hydrodynamic studies of the CE phase, has been the demonstration that mass ejection (due to the spiralling-in binary) occurs preferentially in the orbital plane (Livio & Soker 1988; Taam & Bodenheimer 1989; Term an, Taam & Hernquist 1994, see Fig. 1). Thus, the generation of a density contrast between the equatorial (orbital plane) and polar directions is an almost inevitable consequence of CE evolution. A higher contrast can be expected for more massive secondaries (for a given primary configuration, see Livio & Soker 1988). It is interesting to note that two of the well known PNe which exhibit bipolar outflows (NGC 2346 and Kl-2, Bond & Livio 1990) are known to contain close binary nuclei.

38

M. Livio: Axisymmetric outflows from single and binary stars 1

'

I

I Tto.i. o. I

1

1

1

TfSmoTT

1

1

1

inm. ia,.»

Tlma - 1S0.4

Fig. 1. The particle distribution in the CE in the (x- z) plane. About 80% of the ejected matter was found to be confined to within 30 degrees from the equatorial plane. From the calculation of Terman, Taam & Hernquist (1994).

3.2 Magnetic Fields Magnetic fields are often invoked to explain the formation of configurations which possess axial symmetry. In the case of the shaping of planetary nebulae, magnetic fields could operate (in principle) in at least three different ways: (i) by centrifugally accelerating material in the equatorial plane (in the slow wind), to form a density contrast (to the best of my knowledge, this has actually not been previously suggested for PNe), (ii) by the evolution of a large scale azimuthal field that is embedded in an equatorial torus (in the

M. Livio: Axisymmetric outflows from single and binary stars

39

slow wind), and (iii) by exerting magnetic tension in the equatorial plane in the fast wind (and thus inducing a bipolar flow). In order to obtain a significant centrifugal acceleration of the slow wind, the minimal surface field strength that is needed is (e.g. Michel 1969; Blecher & MacGregor 1976)

M V

/ 2

/

VSw 15 km s

w

-

-l

UU

f

10

where "*" denotes surface values, Vsw iS the speed of the slow wind, VTOt is the equatorial rotation velocity and all other symbols have their usual meaning. Pascoli and co-workers (e.g. Pascoli 1990; Pascoli et al. 1992) suggested that the time evolution of a large-scale azimuthal magnetic field naturally leads to a bipolar morphology. These authors, however, have been somewhat unclear (at least to the present author) about the mechanism generating their initial configuration (in fact, Pascoli et al. 1992 assumed the ejection of an equatorial torus). In any case, for their model to work, Pascoli et al. require at the base of the convective zone and azimuthal field of at least Bv ~ 1000 G. From a dynamo model of fully convective stars (Tout k Pringle 1992) we can estimate the surface field to be

where rj ~ 3R*/lc (lc is the mixing length) and 7 measures the efficiency of the dynamo regeneration term. The azimuthal field component in the convective region can be estimated to be

*

77

L

\2/9

f R. y1*'* \AOORZ)

{

0.1

(3)

where Slcrit is the critical angular velocity, ilcru = (GM./ A comparison of eqs. (2) and (3) with the above requirements (eq. (1) and the required Bv) reveals that the minimum field strength required

40

M. Livio: Axisymmetric

outflows from single and binary stars

by the model of Pascoli et al. may be difficult to achieve in AGB stars. For centrifugal acceleration to have any effect it is required that the AGB star will rotate at (il*/£lCrit) > 0.02. This requirement is not easy to satisfy for single AGB stars, because of their large moments of inertia. Even when the increase in central condensation is taken into account, typically (n,/£lcrh)AGB ~ (0.01 - Q.l)(n,/ncrit)MS (see also Eriguchi et al. 1992), MS and for low mass stars ($l/SlCrit) ~ 0.01. However, the above requirement can be satisfied if the AGB star is spun-up by a companion (even of very low mass, see discussion on rotation in 3.4 below). In a recent work, Chevalier and Luo (these proceedings) suggest that the magnetic field may become important in the shocked wind bubble. They then obtain (under some simplifying assumptions) that a bipolar flow can be generated due to the presence of magnetic tension in the equatorial plane (and its absence in the polar direction). Chevalier and Luo show that a significant effect is obtained for a surface magnetic field strength of the central star larger than

1/2

/ RN r 1 ( y

m ( ) 2000kms-V U O " cm/ V 0.1 / ' where a is the ratio of the field energy density to the kinetic energy density in the wind, Rjy is the radius of the central star, Vp\y is the speed of the fast wind and Vrot is the equatorial surface rotational velocity of the central star. Again we notice (as in the case of centrifugal acceleration), that the magnetic tension can become more effective if the rotation velocity of the star is a non-negligible fraction of the wind velocity. 3.3 Rotation and Be Stars In a recent important work, Bjorkman and Cassinelli (1992, 1993) proposed that the winds from rapidly rotating (early type) stars are very effectively focused towards the equatorial plane. This is essentially a consequence of conservation of angular momentum in a case in which the gravitational force exceeds the radiation pressure force over a sufficiently large distance. If the velocity with which the gas arrives to the equator is supersonic, a wind compressed equatorial outflow ("disk") forms (see Fig. 2). Bjorkman and Cassinelli (1993) used simplified analytic solutions to show that the formation of the equatorially compressed outflow occurs for Vrot/Vcrit ~ 50% (for

M. Livio: Axisymmetric

outflows from single and binary stars

41

Equatorially Compressed Outflow

Fig. 2. Diagram of the compressed equatorial outflow produced by the wind from a star rotating at half the critical velocity. From the calculations of Bjorkman k Cassinelli (1993).

B2 stars; close to 90% for 0 stars). The analytic results of Bjorkman and Cassinelli were essentially confirmed (small differences exist) by the twodimensional hydrodynamic simulations of Owocki, Cranmer and Blondin (1994; and see also Blondin, these proceedings). An examination of the results of Bjorkman and Cassinelli (1993; see also Owocki et al. 1994) shows that the important parameter in determining whether or not an equatorially compressed outflow forms, is the ratio of the (equatorial) rotation velocity to the wind velocity. The threshold values of Vrot/V<x> (where V^ is the wind velocity at large distances), above which the compressed "disk" occurs are shown in Fig. 3. Since V^ rapidly increases for earlier spectral types than B2 (e.g. Prinja, Barlow & Howarth 1990; Bjorkman 1989), these stars must rotate faster in order to form an equatorial "disk." This may explain the decrease in the frequency of Be stars towards earlier spectral types than B2. Bjorkman and Cassinelli (1993) showed (using some simplifying assumptions) that when the dependence of the force multiplier on the ionization balance in the wind is included, a minimum

42

M. Livio: Axisymmetric

outflows from single and binary stars

Voo/V. esc

Fig. 3. The threshold values of KW^co (see text) above which an equatorially compressed outflow is obtained. Calculated on the basis of the results of Bjorkman & Cassinelli (1993). in the threshold (for "disk" formation) rotational velocity is obtained near spectral type B2. They suggested that this may explain why the maximum frequency of Be stars occurs near B2 (however, this problem is far from being settled, since there exists still a large discrepancy between the theoretical and observed terminal wind speeds at late spectral types). At any rate, the work of Bjorkman and Cassinelli (1993) and Owocki et al. (1994) has shown that in the case of radiatively driven winds in early type stars, once the ratio Ko«/Kx> exceeds some critical value (which depends on the ratio of V^ to the escape speed from the stellar surface, see Fig. 3), an equatorially compressed outflow forms.

3.4 Rotation and AGB Stars—Common Planets and Brown Dwarfs

Envelopes,

It is important to examine the question of whether an equatorially compressed outflow of the type obtained for early type stars (3.3 above) can also form in AGB stars, since this will clearly generate an equatorial to polar density contrast. The first thing to note is that unlike in the case of early type stars, for AGB winds, typically V^/Vesc^l (e.g. Reimers 1977; Weyman 1962; Likkel et al. 1992; Knapp & Morris 1985; Zuckerman & Dyck 1986). An examination of Fig. 3 then suggests that the threshold rotational velocity for the equatorial outflow to form is V ^ / V ^ ~ 0.25 (which corresponds approximately to ft*/ftcrtt ~ 0.35; or it could be as low as fi*/ficrtt ~ 0.2 if V^,IVesc ~ 0.5). Secondly, while the exact mechanism of wind acceleration in AGB stars is

M. Livio: Axisymmetric outflows from single and binary stars

43

not fully understood (radiation pressure on dust is probably important; see e.g. poster paper by Netzer and review by Wood, these proceedings), the following is very probably true. If the AGB star were to rotate sufficiently fast (for centrifugal support to be important), then the distance between the sonic point in the wind and the location of loss of centrifugal support can be expected to be sufficiently large, for gravity to focus the flow towards the equator. In the absence of direct observational material, the main question that needs to be answered of course is: can stars be expected to rotate as fast as ft*/Merit ~ 0.35 at least during some stage of their AGB phase"! As discussed in Section 3.3 above, such high rotation rates are highly unlikely for truly single AGB stars. However, it is interesting to examine the possibility of the star being spun-up by a companion, either tidally or via a CE phase. 3.4-1 Spin-up in a CE phase During the spiralling-in process, the internal binary deposits orbital angular momentum into the CE and it spins it up (see e.g. Iben and Livio 1993 for a review). For example, in a three-dimensional calculation of the CE phase of a 4.67M0 red giant with a 0.94MQ companion, Terman, Taam and Hernquist (1994) found that material in the vicinity of the double core has been spun-up to ~ 40% of the angular velocity of the two cores. Furthermore, we can use the fact that the density in the convective region of AGB configurations (which constitutes most of the envelope mass) behaves approximately as p ~ r~2 (e.g. Soker 1992). Tf we then assume that the convective region can be brought to nearly solid body rotation (due to strong turbulent viscosity coupling), then a secondary of mass A/2 depositing even only 20% of its angular momentum into the envelope (corresponding roughly to the minimum efficiency of energy deposition obtained in hydrodynamic calculations, e.g. Taam & Bodenheimer 1989; Livio k. Soker 1988; Terman et al. 1994), can generate an angular velocity of fi»/QCrtt ~ Q.%M2/Menv). Therefore, AGB stars with stellar companions that are sufficiently close to go through a CE phase can certainly be brought in many cases to the rotation rates that are necessary for the production of the equatorially compressed outflow (and thereby for the generation of a density contrast in the slow wind). 3.4-2 Spin-up by brown dicarfs and planets I shall now examine the question of the spin-up of a single AGB star by brown dwarfs or planetary companions. The potential role of planets and

44

M. Livio: Axisymmeiric

outflows from single and binary stars

brown dwarf is very relevant, since recent observations of many young stellar objects (e.g. in the Taurus-Auriga dark clouds, in the Orion nebula, in the L1641 molecular cloud) reveal that perhaps most (if not all) solar type stars are initially surrounded by disks (see e.g. Beckwith et al. 1990; O'Dell, Wen k Hu 1993; Strom, Strom k Merrill 1993 and references therein). While some recent searches (e.g. Murdoch, Hearnshaw k Clark 1993) find brown dwarfs to be rare in orbits closer than 10 AU, brown dwarfs (and high-mass planets) could still be found in wider orbits. In a series of recent works, Tassoul and Tassoul (1990,1992 and references therein) suggested that a hydrodynamic spin-down (up) mechanism can be much more efficient in synchronizing close binary systems than equilibrium tides (which are based on diffusive transport of angular momentum). The hydrodynamic mechanism operates by large-scale meridional transport of angular momentum, which is regulated by an Ekman-type suction layer (in a somewhat similar manner to the spin-down of a stirred cup of tea). On the basis of the results of Tassoul and Tassoul (1992) and Zahn (1977) it is easy to show (see also Soker 1994) that the orbital decay time of a companion around an AGB star is given approximately by Tdec ~

1.2 x 105 yr

MxV 1 / 8 M&)

\ Mi )

V300/

where 6 is the ratio of the synchronization time to the spin-down time, q = M2/M1 is the mass ratio, a is the initial separation and all other symbols have their usual meaning. I now make the following assumptions: 1. The luminosity on the AGB is given approximately by (Eggleton, private communication) x

1O 5 - 3 (A/ C /M 0 ) 6

_ f>*/

+ 10°-5(A/c/A/®)5'

where Mc is the core mass. 2. The radius of the AGB star is related to its luminosity and mass by (Eggen & Iben 1991) p

/ T

\ 0.68 / 1/

\ -0.1(5

A ~ (h.) (EL)

M. Livio: Axisymmetric outflows from single and binary stars

45

3. Mass loss rate is approximately given by a Reimers (1977) type formula

While this formula clearly does not represent the exact mass loss rate at all phases (see e.g. the review by Wood, these proceedings), it is adequate for our present limited purposes. 4. The relation between the total mass and the core mass is given by the results of Iben and Tutukov (1985, their Fig. 30). Under these assumptions, I calculated the evolution of a system consisting of an AGB star (with an initial mass on the AGB of A/° = IMQ and a core mass of Mc = 0.6MQ) and a low mass companion. The calculation showed that for an initial separation cio satisfying oo/R® ~ 5.5 (corresponding to do ~ 2020RQ), the orbit decayed completely during the AGB phase. For an initial separation cio ~ 2OOO.R0, the deposition of the orbital angular momentum into the envelope, could spin the envelope up to ftro
46

M. Livio: Axisymmetric

outflows from single and binary stars

4 Accretion Disks in Post Common Envelope Binaries As described in the Introduction, several PNe exhibit highly collimated outflows ("jets") and some possess an exact point symmetry (see also review by Schwarz, these proceedings), which is reminiscent of some radio maps of AGNs (e.g. Laing & Bridle 1987). This type of point symmetry in AGNs has been interpreted as possibly originating from a precessing jet. The kinematics of the blobs observed in IC 4634 are indeed consistent with a precessing jet interpretation (Corradi & Schwarz 1993 and Schwarz, these proceedings). Since the formation of jets (at least in AGNs and young stellar objects) is thought to be associated with accretion disks (which can provide the axial symmetry, the energy source and the magnetic field which is perhaps required for acceleration and collimation, e.g. Blandford 1993), it is interesting to examine the possibility of forming accretion disks around PNe central stars. This question has been recently examined by Soker and Livio (1994), who proposed the following scenario for the formation of accretion disks in some post CE binary central stars. Soker and Livio noted that the spirallingin timescale in the final stages of the CE is very short (not more than a few tens of years). Thus, the entropy profile in the secondary star remains frozen (Hjellming and Taam 1991). Typically, some secondaries emerging from the CE phase could retain a few hundredths of a solar mass of accreted high entropy CE material. Upon being exposed to the new boundary conditions (after the ejection of the CE), these secondaries (which depart significantly from thermal equilibrium) will start to expand adiabatically. Since the orbital separation is minimal at this stage, it is possible that some of these secondaries will start transferring mass to the PN core, thus forming accretion disks (and possibly jets). This phase should last no longer than the thermal relaxation time of the outer layers (~ 103-104 yrs). 5 Statistics It is of interest to calculate how many PNe nuclei are expected to undergo CE evolution (and thus to be able to produce a density contrast via CE ejection) and to compare this number with the statistics of elliptical and bipolar PNe. Yungelson, Tutukov and Livio (1993) assumed that the present birthrate of all binaries in the Galaxy can be written as dzv = 0:2M^25f(q)dM1d\ogAdq

yr" 1 ,

(9)

where A is the separation and f(q) gives the distribution in mass ratios, taken from Mazeh et al. (1992). They then followed the evolution of the

M. Livio: Axisymmetric I

I

|

1 1 "7

47

outflows from single and binary stars ]

1 1 1 1 1 1 1 —i—i—i—i—i—i—i—

14 A

12

0. M o

.1

08

m la XI

.06

r

/

r

\

/

-

04

.02

-

- y/ \ \ / J

n log(P/day)

,

y.-,

L_,VL

I

"

10

Fig. 4. The distribution of the birthrate of planetary nebula nuclei over orbital period. The solid line denotes CO nuclei accompanied by main sequence stars, the dotted line - He nuclei with main sequence companions and the dashed line - nuclei with white dwarf companions. From Yungelson, Tutukov k Livio (1993).

entire binary population by taking into consideration the evolution of the single stars, all mass and angular momentum exchange and loss episodes, CE phases etc. The total birthrate of PNe in the Galaxy that was obtained was 0.87 yr" 1 , in good agreement with observational estimates (e.g. Phillips 1989). The distribution of birthrate as a function of orbital period is shown in Fig. 4. Yungelson et al. found that ~ 22% of all PNe resulted in close binary nuclei or single nuclei that were obtained by mergers. Within the uncertainties that are involved both in the theory and in the observations, this fraction agrees reasonably well with expectations based on the observational search for multiplicity in 164 primaries of spectral type F7-G9 conducted by Duquennoy and Mayor (1991). The fraction of stars that were found by these authors to be binaries with sufficiently short separations to enable them to interact through a CE phase was ~ 17% (Livio 1993). Again within the uncertainties, this is also consistent with the statistics obtained by Corradi and Schwarz (poster, these proceedings) for bipolar (~ 13%),

48

M. Livio: Axisymmetric outflows from single and binary stars Table 1. Planetary Nebulae with Close Binary Nuclei PN

Central Star

A 41 DS 1 A 63 A 46 HFG 1 Kl-2 A 65 Ht Tr4 BE UMa Sp 1 NGC 2346

MTSer KV Vel = LSS 2018 UUSge V477 Lyr V664 Cas VW Pyx

Orbital Period (Days)

V651 Mon

0.113 0.357 0.465 0.472 0.582 0.676 1.00 1.71 2.291 2.91 15.991

Table 2. Planetary Nebulae with Binary Central Stars Which Show Photometric Variability (Which Probably Does Not Represent the Orbital Period) Planetary Nebula

Central Star

A 35 Lo Tr 1 Lo Tr 5

LW Hya = BD -22°3467 IN Com = HD 112313

Period of Optical Variability (days) 0.76 3.3 or 6.6 5.9

irregular (~ 17%) and elliptical (~ 65%) PNe (see also Schwarz, Corradi & Stanghellini 1993). The PNe with known close binary nuclei are presented in Table 1. The most recent addition is BE UMa, the PN around the previously known binary having recently been detected in 0 III and Ho by Tweedy, Liebert and Bond (1993). Iben and Livio (1993) attempted to reconstruct the evolutionary history of the system. Other central stars for which clear periodicities have been observed (which may, however, represent the rotation of the cool star; Bond, Ciardullo and Meakes, private communication) are the central stars of A 35, Lo Tr 1 and Lo Tr 5 (see Table 2; a suggestion that the nucleus of Lo Tr 5 is a triple system was made by Acker 1993). The central star of Sh2-71 exhibits light variations with a tentative period of 68^06 (Jurcsik

M. Livio: Axisymmetric outflows from single and binary stars

49

1993), however, more observations will be required to determine the exact nature of the system. Since all of these binary systems were found in a sample which involved ~ 110 objects (Drummond 1980; Drilling 1985; Bond, Ciardullo & Meakes 1992), again, a fraction of 10-15% of close binary nuclei is indicated. It should be noted, however, that it is extremely difficult to detect companions which are of a very low mass (see also Yungelson et al. 1993 for a discussion).

6 Suggestions It is perhaps premature to call the following points "conclusions," so they should be treated as suggestions: 1. The "interacting winds" model in the presence of a "density contrast" is responsible for PNe morphology. 2. I propose that equatorially compressed outflows that are similar to the ones generated in early-type stars can also form for (rapidly rotating) AGB stars and similar configurations. 3. Binary stellar companions or brown dwarfs and massive planets may play a crucial role in the generation of the density contrast, either directly (via common envelope ejection) or by spinning-up the envelope of the AGB star, leading to the formation of an equatorially compressed outflow. Mergers of the companions with the core are possible (possible examples are EGB 5 and PHL 932, Mendez et al. 1988a, b). 4. Accretion disks may form in some cases, following the common envelope phase. These disks may play a role in the production of highly collimated outflows and in the formation of point symmetric morphologies. 5. Planetary nebulae with binary nuclei provide us with the best source of information on common envelope physics, via.: • Comparison of observed and calculated distributions for the binary central stars (e.g. Yungelson et al. 1993). • Reconstruction of the evolutionary history of individual systems (e.g. Iben & Tutukov 1993; Iben & Livio 1993). • The observed morphologies. 6. The carbon star V Hya (with an observed Vsini ~ 10-20 km s" 1 ; Kahane et al. 1993) may represent the exciting possibility of a CE "in the act."

50

M. Livio: Axisymmetric

outfloivs from single and binary stars

7. PNe form a new addition to the classes of objects (formerly including AGNs, young stellar objects and some x-ray binaries) exhibiting "interesting" bipolar outflows and "jets." Acknowledgements I am grateful to Stan Owocki, Jean-Louis Tassoul and Lauren Likkel for very helpful discussions. This work has been supported in part by the Fund for Basic Research of the Israel Academy of Sciences at the Technion and by the Director's Research Fund at Space Telescope Science Institute. References Acker, A. (1993). in IAU Symp. 155, Planetary Nebulae, ed. A. Acker et al., in press. Balick, B. (1987). A. J., 94, 671. Beckwith, S. V. W., Sargent, A. I., Chini, R. S. k Glisten, R, (1990). A.J., 99, 924. Bjorkman, K. S. (1989). Ph.D. Thesis, University of Colorado. Bjorkman, J. E. &; Cassinelli, J. P. (1992). in Nonisotropic and Variable Outflows from Stars, eds. L. Drissen, C. Leitherer k A. Nota (San Francisco: ASP), p. 88. Bjorkman, J. E. k Cassinelli, J. P. (1993). Ap. J., 409, 429. Blandford, R. (1993). in Astrophysical Jets, eds. D. Burgarella, M. Livio k C. O'Dea (Cambridge: Cambridge University Press), in press. Belcher, J. W. k MacGregor, K. B. (1976). Ap. J., 210, 498. Bobrowsky, M. (1993). in preparation. Bond, H. E. k Livio, M. (1990). Ap. J., 355, 568. Bond, H. E., Ciardullo, R. & Meakes, M. G. (1992). in IAU Symp. 151, Evolutionary Processes in Interacting Binary Stars, ed. Y. Kondo et. al. (Dordrecht: Kluwer), p. 577. Burgarella, D. k Paresce, F. (1992). Ap. J., 370, 590. Corradi, R. k Schwarz, H. (1993a). A. & A., 269, 462. . (1993b). A. & A., 273, 247. . (1993c). A. & A., 268, 714. Drilling, J. S. 1985, Ap. J., 294, L107. Drummond, J. (1980). Ph.D. Thesis, New Mexico State University. Duquennoy, A. k Mayor, M. (1991). A. & A., 248, 485. Eggen, O. J. k Iben, I. Jr. 1991, A. J., 101, 1377. Eriguchi, Y., Yamaoka, H., Nomoto, K. k Hashimoto, M. 1992, Ap. J., 392, 243. Frank, A., Balick, B., Icke, V. k Mellema, G. (1993). Ap. J., 404, L25. Hjellming, M. S. k Taam, R. E. (1991). Ap. J., 370, 709. Iben, I. Jr. k Tutukov, A. V. (1984). Ap. J. Suppi, 54, 335. Iben, I. Jr. k Tutukov, A. V. (1985). Ap. J. Suppi, 58, 661 Iben, I. Jr. k Tutukov, A. V. (1993). Ap. J., 418, 343. Iben, I. Jr. k Livio, M. (1993). P.A.S.P., 105, 1373. Icke, V., Balick, B. k Frank, A. (1992). A. & A., 253, 224. Jurscik, J. (1993). in IAU Symp. 155, Planetary Nebulae, eds. R. Weinberger k A. Acker (Dordrecht: Kluwer), p. 399. Kahane, C , Audinos, P., Barnbaum, C , Morris, M., 1993, in Mass Loss on the AGB and Beyond, ed. H. Schwarz, ESO Conference and Workshop Proceedings No. 46, p.437 Kahn, F. D. (1982). in IAU Symp. 103, Planetary Nebulae, ed. D. R. Flower (Dordrecht: Reidel), p. 305.

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Knapp, J. k Morris, M. 1985, Ap. J., 292, 640. Kwok, S. (1982). Ap. J., 258, 280. La Bertre, T. (1986). The Messenger (ESO), No. 44, p. 6. Lai, D., Rasio, F. A. k Shapiro, S. L. (1993a). Ap. J., 406, L63. Lai, D., Rasio, F. A. k Shapiro, S. L. (1993b), Ap. J., submitted. Laing, R. A. k Bridle, A. H. (1987), MNRAS, 228, 557. Likkel, L., Morris, M. k Maddalena, R. J. (1992), A. & A., 256, 581. Livio, M. k Soker, N. (1988). Ap. J., 329, 764. Livio, M. (1993), in IAU Symp. 155, Planetary Nebulae, eds. R. Weinberger k A. Acker (Dordrecht: Kluwer), p. 279. Mazeh, T., Goldberg, D., Duquennoy, A. k Mayor, M. 1992, Ap. J., 401, 265. Mendez, R. H., Kudritzki, R. P., Herrero, A., Hasfeld, D. k Groth, H. G. (1988a), A. & A., 190, 113. Mendez, R. H., Groth, H. G., Hasfeld, D., Kudritzki, R. P. k Herrero, A. (1988b). A. & A 197 L25 Michel,'F. C. (1969). Ap. J., 158, 727. Murdoch, K. A., Hearnshaw, J. B. k Clark, M. (1993). Ap. J., 413, 349. O'Dell, C. R., Wen, Z. k Hu, X. (1993). Ap. J., 410, 696. Owocki, S. P., Cranmer, S. R. k Blondin, J. M. (1994). Ap. J., (in press). Paczynski, B. (1976). in The Structure and Evolution of Close Binary Systems, eds. P. Eggleton, S. Mitton and J. Whelan (Dordrecht: Reidel), p. 75. Pascoli, G. (1990). in From Miras to Planetary Nebulae: Which Path for Stellar Evolution?, eds. M. D. Mennessier k A. Omont (Gif sur Yvette: Editions Frontires), p. 136. Pascoli, G., Leclerco, J. k Poulain, B. (1992). P.A.S.P., 104, 182. Phillips, J. P. (1989). in IAU Symp. 131, Planetary Nebulae, ed. S. Torres-Peimbert, (Dordrecht: Kluwer), p. 425. Prinja, R. K., Barlow, M. J. k Howarth, I. D. 1990, Ap. J., 361, 607. Reimers, D. (1977). A. & A., 61, 217. Schwarz, H. E., Corradi, R. L. k Melnick, J. (1992). A. & A. Suppl., 96, 23. Schwarz, H. E., Corradi, R. L. M. k Stanghellini. L. (1993). in IAU Symp. 155, Planetary Nebulae, eds. A. Acker et al., p. 214. Soker, N. k Livio, M. (1989). Ap. J., 339, 268. Soker, N. k Livio, M. (1994). Ap. J., 421, 219. Soker, N. (1994). P.A.S.P., in press. Soker, N. (1992). Ap. J., 389, 628. Strom, K. M., Strom, S. E. k Merrill, K. M. (1993). preprint #790. Taam, R. E. & Bodenheimer, P. (1989). Ap. J., 337, 849. Taam, R. E. & Bodenheimer, P. (1991). Ap. J., 373, 246. Tassoul, J.-L. & Tassoul, M. (1990). Ap. J., 359, 155. Tassoul, J.-L. & Tassoul, M. (1992). Ap. J., 395, 259. Terman, J. L., Taam, R. E. & Hernquist, L. (1994). Ap. J., 422, 729. Tout, C. A. & Pringle, J. E. (1992). MNRAS, 256, 269. Tweedy, R. W., Liebert, J. W. k. Bond, H. E. (1993). private communication. Webbink, R. F. (1984). Ap. J., 277, 355. Weyman, R. (1962). Ap. J., 136, 476. Yungelson, L. R., Tutukov, A. V. k Livio, M. (1993). Ap. J., 418, 794. Zahn, J.-P. (1977). A. & A., 57, 383. Zuckerman, B. k Dyck, H. M. (1986). Ap. I., 304, 394.

Flows in Clumpy Circumstellar Media J. E. Dyson1 and T. W. Hartquist2 Department of Astronomy, The University of Manchester, Manchester MIS 9PL, England 2 Max Planck Institute for Extraterrestrial Physics, D-85740 Garching, Germany

Abstract Mass addition to flowing tenuous plasmas by the ablation of embedded clumps alters the flows and the observational characteristics of both flows and clumps. The boundary layers between the clumps and the flows are the sites of enhanced radiative losses. Flows which in smooth media would be driven by thermal pressure, are instead driven by momentum. There are many possible types of flows and we explore some of them in the context of Wolf-Rayet and planetary nebulae. Flows in which transsonic tenuous plasmas exit from mass-loading cores into smooth haloes are relevant for planetary nebulae. On intermediate scale lengths, the flow-clump interactions produce extended 'tails'. We give a general discussion of this and describe applications to the cometary tails behind globules in the planetary nebula NGC 7293 and to the tail of the galactic centre red supergiant IRS7. We finally briefly discuss diagnostics of the boundary layers themselves.

1 Introduction Practically all diffuse media of astrophysical significance are clumpy media which are responding to energy sources. The most important distinction between flows initiated in clumpy as opposed to homogeneous media, is that in the former, there is mass, momentum and energy interchange at clumptenuous plasma boundaries, i.e. in boundary layers. The consequences are major (Hartquist & Dyson 1993). This interchange reacts back on the dynamical, physical and even chemical state of the global tenuous plasma flow; conversely, the state of the global flow influences the interchange process. Structures on intermediate length scales are found where material has left the clump but has not yet lost its identity by assimilation into the global flow. Because of the directionality of the tenuous plasma flow, these intermediate scale structures are extended, i.e. they are 'tails1. Finally, the clumps themselves are affected by the loss of material. We discuss the effects of clumps on the large scale structures of winddriven flows with emphasis on three particular objects and classes of object: the clumpy Wolf-Rayet nebula RCW58; planetary nebulae (PNe); corehalo PNe with transsonic flows exitting from mass loading regions local52

J. E. Dyson & T. W. Hartquist: Flows in clumpy CSM

53

ized around the energy input sources, which then form high temperature giant haloes. We next describe the principles of intermediate scale structure formation and draw general conclusions from simple yet robust models We apply these principles to two specific problems; the long cometary tails visible behind the small globules in the PNe NGC 7293 and the tail of the galacto-centric M supergiant IRS7. Finally, we comment on diagnostic problems relevant to the boundary layers themselves and indicate optimum observational possibilities.

2 Effects on Large Scale Flows Pikelner's (1968) "two-shock" flow pattern has been the basis of most subsequent investigation of the effects of hypersonic winds on their surroundings. The expansion of the wind driven bubbles is governed by the thermodynamic behaviour of the shocked wind gas. If the cooling time in this gas is longer than the dynamical timescale, the bubble is driven by adiabatic expansion. If the cooling timescale is the shorter timescale, the bubble is driven by wind momentum. The most important parameter which determines which of these two regimes is operative is the wind speed (Dyson 1984). As a guide, wind velocities in excess of 200-300 km s" 1 result in pressure driven interactions, whereas velocities lower than this result in momentum driven flows. The transition from one regime to another is effectively discontinuous, and density gradients in the circumstellar material can cause flows to evolve from one regime to another (Dyson 1984). However, these conclusions are valid only if the surrounding medium is smooth. The presence of clumps completely changes the rules. Hartquist et al. (1986-H86) gave a phenomenological description of the ablation of clump material into the surrounding flow and argued that the mass loading rate is proportional to Mp, where Mp is the flow Mach number in the clump frame, and f3 = 0 if Mp > 1 and j3 = 4/3 if Mp < 1. We first make a general point regarding the effects of mass loading on flows in the context of steady 1-D flows, but one which turns out to have widespread implications. The continuity and momentum equations (cf. equations (5) of H86) give a schematic equation for the variation of the flow Mach number Mp with the distance coordinate r, dMF _ \ [TERMS] - [MASS - LOADING TERM] ]

~dT~\

(Ml - 1)

/•

()

In this equation, [TERMS] would include the effects of geometrical divergence, etc. Irrespective of whether a flow is subsonic or supersonic, mass

54

J. E. Dyson & T. W. Hartquist: Flows in clumpy CSM

loading always forces Mp towards unity. Thus H86 speculated that once mass loaded flows reach a Mach number of around unity, they would remain there. We argue on observational and theoretical grounds that this speculation has widespread validity.

3 The Dynamics of RCW58 The low (~ 5MQ) mass, clumpy structure and observed He and N enrichment imply that RCW58 is largely composed of stellar ejecta from the red supergiant phase preceding the present WR stellar phase (Smith et al. 1988— S88). Much less than 1% of the WR. wind kinetic energy appears to have been converted into kinetic energy of nebular material (Smith et al. 1984S84). 'Classical' wind-blown bubble theory predicts a. roughly 20% conversion efficiency for the pressure driven bubble which should have been produced by a 2000 km s" 1 WR wind (e.g. Dyson 1989). IUE absorption data on features formed in the bubble (S84) are clearly at variance with the 'classical' structures. The sense of the observed ionization potentialvelocity correlation is such that the only possible site for the formation of the absorption features is in the shocked wind zone. We assume that the flow behind the stellar wind shock is plane parallel and isobaric. In the observers' frame, the wind velocity and shock velocity are respectively —Vw and — V5. The observed ionic species with the highest and lowest ionization potentials are C 3 + and Fe + which exist respectively at temperatures T\ (~ 105K) and T2 (^ 104 K) and at corresponding densities p\ and p2.

Since the flow is isobaric, hP\T\

hp2T2

, pw(nv

=

.2 - Ks)

f0\ (^)

// //. where kf> is Boltzmann's constant, fi is the mean-mass per particle and pw is the immediate pre-shock density of the stellar wind. If we assume that the subsonic flow behind the shock mass-loads by a factor cf> between the post-shock wind gas and gas at Tj, and that the loading between gas at T\ and T2 is negligible, the continuity condition is -P1V1 - -P2V2 = -<j>Pw{V\v - Vs).

(3)

From equations (2) and (3), in the observers' frame, the velocity difference between the gas at T\ and T2 is kb

d> ^ ~

^

T

^

(4) (4)

J. E. Dyson & T. W. Hariquist: Floivs in clumpy CSM

55

Since Vw — Vs — 1800 km s" 1 (from the known wind speed), and the observed AV is —45 km s" 1 (S84), then cf> ~ 50, i.e. the post-shock flow is dominated by mass addition from clumps. H86 gave a phenomenological explanation of why takes the required value. They argued that the post-shock flow was driven to a Mach number of about unity (Section 2) and remained there until radiative cooling became important, and showed that under these circumstances, the ratio of entrained mass to wind mass is (25/8)(7V/2T,. ac |) 2 , where Tw is the initial post-shock wind temperature (w 4 107 K) and T r a d(« 105 K) is the temperature at which the radiative cooling timescale becomes shorter than the timescale for significant mass addition. Hence the mass ratio is about 44, in good agreement with that required. The absorption features are formed in a very thin layer (T ~ 105 K) in which negligible mass loading occurs. The strong radiative cooling removes most of the wind kinetic energy, thus accounting for the low conversion efficiency. Arthur, Dyson & Hartquist (1993) have constructed a spherically symmetric time dependent numerical model to describe R.CW58. A stellar wind mass loss rate of 1O~5M0 yr" 1 and wind velocity of 2000 km s" 1 were adopted. The mass loading rate was \0~34 Mp gm cm" 3 s" 1 . For a uniform ambient density distribution of 1 cm" 3 (the interclump medium) and an age of 13000 yr, they showed that the observed nebular radius and expansion velocity (S88) and, critically, the velocity spread in gas at temperatures between 105 K and 104 K, were well matched. However, the model temperature-velocity relationship is neither linear nor monatomic in that temperature range in contrast to the relationship inferred in S84 and H86 on the assumption that ionization equilibrium held. It is likely that this discrepancy can be resolved only if the non-equilibrium ionization structure is followed simultaneously with the hydrodynamics. This presents conceptual problems, in particular, the effect that the ablated material has on the ionization structure of the flow. In this flow, the interclump medium determines the position of the wind shock, implying that the clumps are neither too numerous nor massive. However, mass injection from the clumps totally dominates the shocked wind flow. As in the 'classical' case, the wind shock and that driven into the interclump medium move systematically outwards from the star. This model is probably too simplistic in its assumption of distributed mass loading sources. The morphology of the nebula shows that it contains a limited number of relatively large clumps. The mass loading zone may be located behind a stationary large scale bow shock around such a clump.

56

J. E. Dyson & T. W. Hartquisi: Flows in clumpy CSM

However, the overall energetics of RCW58 must be dominated by the effects of these localized zones. 4 Planetary Nebula Dynamics If clumps are numerous, massive and dense enough, the interclump medium plays little role in establishing the radius at which most of the wind is shocked. The wind is decelerated primarily in bow shocks around the clumps. Provided that the clumps are not substantially ablated, the positions at which the wind is shocked either do not move systematically outwards, or move more or less with the clump velocity, in clear distinction to the situation obtaining in RCW58. There is a substantial body of observational evidence showing that the AGB ejecta. wliich on pliotoionization give the major contribution to the mass of PNe, is clumpy clown to all resolvable scales. We assume therefore that clumps totally dominate the flow. This wind percolating through the clumps becomes mass-loaded and will be headed by a shock driven into the interclump medium. Initially, the mass-loaded wind also shocks and the flow consists of a supersonic mixture of ablated and wind gas which passes through the inner shock. The shocked mixture is separated from shocked interclump gas by a contact discontinuity. Implicit in this description is the assumption that the energy generated as the wind is decelerated in individual bow shocks can be lost radiatively, most likely in the interface regions. The temperature of the gas in theflowregions will be maintained at a temperature of about 104 K by pliotoionization induced by the radiation field of the central star. Provided the outer shock velocity is low enough, the supersonic interior flow is steady. At radii r > rjy = (SM./STT*/) 1 / 3 , the flow velocity v oc r~3 and the density p oc r4. M* is the stellar wind mass-loss rate and q the mass loading rate per unit volume. If the flow is isothermal, the Mach number is Mp — (K/c,)(rjvf/r)3 and approaches unity when r = r\ « (V./c,) 1 / 3 ^/ ~ 6rjv/ characteristically. Provided that the mass loading zone extends out further than radius r\, the flow should tend to a constant Mach number of around unity (H86). The inner shock becomes very weak or may vanish. Arthur, Dyson & Hartquist (1994) have simulated flows where the mass loading zone extends out further than r\. They assume that the gas behaves isothermally. They find that indeed the flow Mach number drops to a more or less constant value, but that the constant value is about 0.6-0.7. This is because of spherical divergence. An important application of these ideas is to the study of the extended haloes of PNe. PNe morphology supports the idea that the PNe clumps are

J. E. Dyson & T. W. Hartquist: Flows in dumpy CSM

57

localized around the central star as a result of the finite lifetime of the AGB ejecta phase. They may be surrounded by a lower density extended red giant envelope. The transsonic mass loaded wind eventually exits from the mass loading zone into a lower density region (which may itself be clumpy). The pressure in a transsonic flow is non-negligible, and its expansion is partially pressure driven. The flow can accelerate provided that its temperature is maintained by photoionization. The flow gently accelerates and the Mach number has a roughly logarithmic increase with radial distance. Hence Mach numbers of a few at most would occur in the halo regions.

Multiple shell PNe (Chu, Jacoby & Arendt 1987) are likely examples of bubbles with clumpy cores and gently accelerating haloes. Although few spectroscopic data for the haloes are available, we would expect them to show flow speeds corresponding to Mach numbers of about unity near their inner edges and up to a few further out. Indeed some. PNe seem to show this behaviour (Chu 1988).

Meaburn et al. (1991) made echelle studies of the halo of the PNe NGC 6543 and showed that the electron temperatures in an arc-like feature in the halo was about 15000 K, whereas the velocity dispersion in the gas was only 5-6 km s" 1 . They attributed this feature and its elevated temperature to a bow shock formed around a clump in the halo by the impact of a gently accelerating transsonic flow as described above. The gas flowing against the clump cannot have a speed much greater than 20-25 km s - 1 relative to the clump, otherwise the velocity dispersion would be too high. A shock decelerating gas at the core electron temperature of about 9000 K would heat the gas to the observed temperature of about 15000 K. The flow Mach number of 2-3 is in harmony with, and thus supports, the description above. Comparison of the required volume mass-loading rate here (Meaburn et al. 1991) with that required for RCW58 (H86) shows that the former is a factor of about 40 greater than the latter. This large difference in the mass loading rate is the reason why, in RCW58, the inner shock decelerating the wind is in the inner parts of the mass loading zone, whereas the wind in a PNe is decelerated in bow shocks around clumps distributed over a large part of the mass loading zone. Hot haloes have been observed in other PNe (Middlemass et al. 1990), but individual bow shocks have not been resolved, possibly due to too low an angular resolution. Their presence may still be responsible for the elevated halo temperatures (Dyson 1992).

58

J. E. Dyson & T. W. Hartquist: Flows in clumpy CSM

5 Intermediate Length-Scale Structures-Tails In general, clumps will be embedded in global flows with systematically directed velocity fields. Thus mass ablated from clumps will be preferentially accelerated in the radial direction and produce elongated structures or 'tails'. Dyson, Hartquist & Biro (1993) give simple analytic and semi-analytic models of tail shapes and appearances expected from the various categories of stream-obstacle systems resulting from the interaction of streams and sources. The advantage of these simple models is that they can be used to infer quite general yet robust conclusions about tail structures. In the following discussion, the stream can ordinarily be identified with a wind (mass loaded or not) and the source with the evaporative flow from an obstacle. (a) Supersonic stream interactions with sources. If the source is supersonic, the flow contains shocks in both the stream gas and the source gas. The post-shock regions are separated by a contact discontinuity, E, the shape of which can be calculated approximately from the balance of the normal components of momentum flux in the pre-shock flows. If the stream is plane parallel, the shape of E is invariant with respect to the relative momentum fluxes in stream and source (Dyson 1974). If both originate from point sources, S depends on these relative fluxes, but if, say, the stream dominates, either case in principle can give a, surface E which has a low width to length ratio. This does not imply that a long thin tail would be observed. If, for example, on S the area! emission power is proportional to the areal shock dissipation rate of mechanical energy near S, the tail will appear short and stubby because the areal dissipation rate strongly peaks towards the stagnation line. If the emission comes from gas which has cooled to a photoionization maintained equilibrium state, the areal emission power varies roughly as the gas pressure squared (% normal component of momentum flux on S) 2 . Again the fall-off is rapid. The slowest fall-off of areal emission power occurs if emission comes predominantly in a narrow temperature range considerably below the temperature to which gas on the stagnation line is heated. But even so, geometric divergence of wind flows always results in a rapid fall-off of areal emission away from the stagnation line and thus tails formed by two interacting supersonic winds appear to be short and stubby. If the source is subsonic, only the stream shocks and the shape of S is determined by the balance between the normal component of the wind momentum flux and the thermal pressure of the subsonic source gas. If the ratio of the pressure at some general position in the source gas to the stagnation pressure in the stream is (1 - e), where f < 1 for very subsonic

J. E. Dyson & T. W. Hariquist: Flows in clumpy CSM

59

flow, the tail width is about a factor e" 1 / 2 times the head width (Dyson et al. 1993), i.e. this interaction produces physically fat stubby tails. (b) Subsonic stream interactions with sources. If the source is subsonic with respect to S, the tail shape is determined by the balance of the thermal pressures of the stream and source gases. The tail width-to-head diameter ratio is about Z? 1 / 2 ^" 1 / 4 (Dyson et al. 1993) where the ratio of the downstream confining pressure to the stream stagnation pressure is (1 - 6), where 6 < 1 and /?(~ 1) is the Mach number at which the head empties. Since /3 w 61/2 can be expected, tail widths comparable to head diameters are possible even for very long tails. If the source gas is not highly supersonic and/or does not radiate well, the description is adequate and the effective source size is the characteristic shock radius. If the source gas is highly supersonic and the shock is nonspherical, the post-shock gas can remain supersonic and the shock shape is fixed by the balance of the normal component of the pre-shock wind momentum flux and the sum of the ambient gas pressure and the centrifugal correction due to flow along a curved surface. Canto (1980) calculated cavity shapes; in general they close in on themselves in the direction of decreasing ambient pressure and are crudely spherical abut some point coaxial with the source and displaced from it in the same direction. Long, thin cavities (i.e. long thin tails) cannot be produced by the small pressure gradients existing in subsonic streams. In summary, long thin tails are produced only when the wind is subsonic (i.e. the mass loss is a 'breeze' or the wind has been shocked far upstream) and the mass loss from the obstacle is subsonic.

6 Tails in the PNe NGC 7293 Spectacular high resolution images of the most prominent cometary globules in the highly evolved PNe NGC 7293 have been obtained by Meaburn et al. (1992). The clump gas may originate in SiO maser knots in the atmosphere of the parent red giant or AGB star (Dyson et al. 1989). These authors, on the basis of consideration of the survival times of clumps, predicted that they would be shown to possess molecular cores, as was later confirmed observationally. The tails behind the globules have lengths ~ 1016 cm, tail width-to-head diameter ratios of about unity, and tail length-to-width ratios of 5-10. Meaburn et al. (1992) suggested that the 'wind-swept' morphology was due to the impact of a supersonic wind from the central star on the globule. The discussion above shows that the impacting stream must be subsonic and the source must be no more than mildly supersonic. In

60

J. E. Dyson & T. W. Hariquist: Flows in clumpy CSM

any case, the central star is highly evolved with a low (as 600 LQ) bolometric luminosity and shows no signs of supersonic mass loss (Cerruti-Sola & Perrinoto 1985). A suitable subsonic stream could be produced if hot shocked stellar wind gas remains 'bottled up' in the interior cavity for a long enough period after fast wind activity from the nuclear star has ceased. Small pressure gradients would be expected in this gas as it gradually percolates outwards into the very clumpy PNe shell. Since the stream gas and globule gas will be approximately in pressure equilibrium, then (Dyson et al. 1989), nsTs ~ 107 cm" 3 K, where ns and Ts are the stream density and pressure. If the stream gas consists of hot shocked wind gas uniformly distributed throughout the central cavity, the mass of stream gas is Ms « 4 10~7 (7?L-/1017 cm) 3 (T s /10 8 K)" 1 MQ, where Rc(za 6 10 17 cm) is the radius of the central cavity and Ts « 108 K for a fast wind speed % 3000 km s" 1 . Hence Ms « lO" 4 A/ 0 , and since the fast wind mass loss rates of PNe central stars are ~ 10~' - 1 0 ~ 8 MQ yr" 1 , roughly a 103 - 104 yr supply of stellar wind gas needs to be bottled up. This is between 3-30% of the fast wind output for the dynamical age ( « 4 104 yr) of this PNe. We assume that the tails originate after the time when the fast wind has died away. The maximum tail length is C^ w viy:ite, where v^ is the tail gas velocity and te is some emptying time for the hot cavity gas. If the tail flow is isothermal at sound speed ct, application of Bernouilli's equation shows that its velocity is v^ « Mpct, where the hot shocked wind flows at Mach number M/r(
7 The Tail of the M Supergiant IRS7 Radio continuum observations revealed a cometary 'tail' of ionized gas originating at the M2I red supergiant IR.S7, located at a projected separation of about 6" from the galactic centre (Yusef-Zadeh & Morris 1991). The kinematics of the [Ne +] emission from the tail confirm the physical association of IRS7 and the tail (Serabyn, Lacy k Achtermann 1991). Although (YusefZadeh & Melia 1992) the cometary head can be well modelled by the impact of a hypersonic wind (originating probably from the blue cluster IR.S16) on the supergiant, it is much more difficult to argue that such an interaction

J. E. Dyson & T. W. Hariquisi: Flows in clumpy CSM can produce the long thin (length : width ratio ss 5:1) clumpy tail observed. However, the subsonic stream-subsonic source interaction which accounts so well for the globule tails of NGC 7293 cannot be operative here, since there is a hypersonic (w 50 km s"1) relative velocity between IR.S7 and its tail (Serabyn et al. 1991). An intriguing possibility (Dyson & Hartquist 1994) is that the tail can be produced by the filtering and mass loading of a hypersonic wind as it passes through a clumpy stellar atmosphere in a manner analogous to the PNe interactions described above. However, there are two important distinctions between the two cases. In the case of IRS7 the supersonic wind impacts externally on the stellar atmosphere. Secondly, the star is still in the process of ejecting its envelope. The basis of the model is that whereas with an impermeable source in a supersonic stream the ram pressure of the stream compresses the head but allows evaporation into a large solid angle on the source's downstream side, in a clumpy source, the distribution of clumps determines the head size. Also, the overlap of individual short fat tails around individual clumps with other short fat tails plus enhanced radiative losses produces a composite momentum conserving tail. The sideways spread of the tail is limited because the tail gas sound speed (« 10 km s"1 for T K, 104 K) is appreciably less than the tail velocity (« 50 km s" 1 relative to the star. These values are consistent with the observed tail length : width ratio of about 5:1. Dyson & Hartquist (1994) have shown that the assumption of a wind source with a mass output rate of M ss 310~3M(7) yr" 1 , velocity of 600 km s"1 at a distance of 0.35 pc from IRS7 (Yusef-Zadeh & Morris 1991) and a mass loading rate of ~ 1O~30 gm cm"3 s"1 gives a tail exit velocity of about 50 km s"1. This mass pickup rate is ~ 100 times that deduced for RCW58. However, the momentum flux experienced by clumps here is at least a factor ~ 100 times that experienced in clumps in RCW58 and this may account for the difference. Mass loading rates may also reflect the distribution of clump masses and densities. Clump characteristics can be estimated assuming they must be able to survive long enough to get a distance about equal to the radius of the cometary head of IRS7 (« 6 1016 cm). Estimated clump masses are w 0.1 MQ, appreciably larger than those in NGC 7293, but probably consistent with the larger clumps in RCW58. 8 Boundary Layer Diagnostics Smith et al. (1988) attempted to use optical emission line data to study the boundary layers between clumps and the wind of RCW58. Clumps were

61

62

J. E. Dyson & T. W. Hartquisl: Flows in clumpy CSM

found to have scales of 30", whereas the angular resolution of the data is ~ 2". Unfortunately, the [OIII] emission line showed no significant position dependence in its mean velocity or width along a clump. Its mean velocity and width also showed no significant differences from those of Ho and [Nil] features originating in the clump. Acceleration of ablated material in the boundary layer of an RCW58 clump must occur more slowly than the heating required to alter the ionization balance. Studies of circumstellar matter have yielded considerable insights into the global nature of mass loaded flows and into the intermediate scale flow structures around clumps. However, the best chances of diagnosing mixing, momentum transfer and heating in a clump-wind interface may lie in investigations of T-Tauri wind-molecular dark core interactions in low mass star forming regions (Charnley et al. 1990; Hartquist k Dyson 1993). The effects of magnetic pressure are currently being included in theoretical modelling of the chemistry in such boundary layers (Nejacl & Hartquist 1994). Emissions from CI, CH and the decay of high J-levels of CO are likely to provide diagnostics of the boundary layers. We would emphasize that resolution is as important as diagnostics and boundary layer phenomena in extragalactic sources are unlikely ever to meet this requirement. We are very grateful to our collaborators, Jane Arthur, Susana Biro, Max Pettini and Linda Smith for their manv contributions to the work described. References Arthur, S. J., Dyson, J. E. k. Hartquist, T. W. (1993). Mon. Not. R. astr. Soc, 261, 425. Arthur, S. J., Dyson, 3. E. & Hartquist, T. W. (1994). (submitted). Canto, J. (1980). A&A, 86, 327. Cerruti-Sola, M. & Perrinoto, M. (1985). Astrophys. ./., 291, 237. Charnley, S. B., Dyson, J. E., Hartquist, T. W. fc Williams, D. A. (1990). Mon. Not. R.

astr. Soc, 243, 405. Chu, Y.-H. (1988). In Planetary Nebtdae, IAU Symposium 131, eel. Torres-Peimbert, S., Kluwer Acad. Publ., Dordrecht, p. 105. Chu, Y.-H., Jacoby, G. H. & Arendt, R. (1987). Astropys. J. Suppl., 64, 529. Dyson, J. E. (1974). Ap. Space Set., 35, 299. Dyson, J. E. (1984). Ap. Space Sci., 106, 181. Dyson, J. E. (1989). In Structure and Dynamics of the Interstellar Medium, IAU Colloquium 121, eds. Tenorio-Tagle, G., Moles, M. fc Melnick, .]., Springer Verlag, Berlin, p. 149. Dyson, J. E. (1992). Mon. Not. R. astr. Soc, 255, 460. Dyson, J. E. & Hartquist, T. W. (1992). Ast. Lett. Comm., 28, 301. Dyson, J. E. & Hartquist, T. W. (1994). Mon. Not. R. astr. Soc, (in press). Dyson, J. E., Hartquist, T. W., Pettini, M. and Smith, L. J. (1989). Mon. Not. R. astr. Soc, 241, 625. Hartquist, T. W. & Dyson, J. E. (1993). Quart. J. R. astr. Soc, 34, 57. Hartquist, T. W., Dyson, J. E., Pettini, M. fc Smith, L. J. (1986). Mon. Not. R. astr. Soc, 221, 715 (H86).

J. E. Dyson & T. W. Hartquisi: Flows in clumpy CSM

63

Meaburn, J., Nicholson, R., Bryce, M., Dyson, J. E. fe Walsh, J. R. (1991). Mon. Not. R. astr. Soc, 252, 535. Meaburn, J., Walsh, J. R., Clegg, R. E. S., Walton, N. A., Taylor, D. &; Berry, D. S. (1992). Mon. Not. R. astr. Soc, 225, 177. Middlemass, D., Clegg, R. E. S., Walsh, J. R. k Harrington, J. P. (1990). In From Miras to Planetary Nebulae: Which Path for Stellar Evolution?, eds. Mennessier, M. O. & Omont, A., Editions Frontieres, Gif-sur-Yvette, p. 420. Nejad, L. A. M., & Hartquist, T. W. (1994). Mon. Not. R. astr. Soc, (in press). Pikelner, S. B. (1968). Astrophys. J. Lett., 2, 97. Serabyn, E., Lacy, J. H. fc Achtermann, J. M. (1991). Astrophys. J., 378, 557. Smith, L. J., Pettini, M., Dyson, J. E. & Hartquist, T. W. (1984). Mon. Not. R. astr. Soc, 211, 697 (S84). Smith, L. J., Pettini, M., Dyson, J. E. k. Hartquist, T. W. (1988). Mon. Not. R. astr. Soc, 234, 625 (S88). Yusef-Zadeh, F. & Melia, F. (1992). Astrophys. J., 385, L41. Yusef-Zadeh, F. & Morris, M. (1991). Astrophys. J.. 371. L59.

Ring Nebulae around LBVs and WR stars Linda J. Smith Department of Physics and Astronomy, University College London, Gower St., London WClE 6BT, U.K.

Abstract WR stars and their precursors, the LBVs, represent the late stages of evolution of hot, massive stars, and are often surrounded by ring nebulae. These are believed to be formed either by the action of the stellar wind, a past, episode of violent ejection from the star, or a combination of these two processes. The various research applications of LBV and WR nebulae are reviewed, particularly with regard to the information they provide on the central stars. Abundance studies show that N overabundances and O deficiencies are a general feature of ejecta around evolved massive stars. Observations of bipolarity in LBV nebulae provide valuable clues to wind asymmetries in the central stars. The nebulae can also be used to derive stellar effective temperatures through photoionization modelling. Finally, the connection between LBV and WR nebulae from an observational point of view is discussed.

1 Introduction The ring nebulae that are observed around Luminous Blue Variables (LBVs) and Wolf-Rayet (WR.) stars are examples of circumstellar media in the late stages of stellar evolution of hot, massive stars. These nebulae are excellent laboratories for studying the interaction of winds and ejecta with the interstellar medium (ISM). They also provide unique insights into the central stars, particularly from an evolutionary point, of view. In Sect. 2, LBVs and WR stars are introduced, and Sect. 3 discusses the formation and composition of their nebulae. Sect. 4 describes how LBV and WR nebulae can be used to study a variety of astrophysical problems, and in Sect. 5, the connection between LBV and WR. nebulae is discussed. The study of LBV nebulae is a fairly young area of research; LBVs were only defined as a distinct class of stars ten years ago (Conti 1984). Recent reviews on WR and LBV nebulae are those of Stahl (1989), Smith (1990) and Chu (1991).

2 LBVs and W R stars Just over a decade ago, it was recognised that the observed H-R diagram has a temperature-dependent upper luminosity boundary (Humphreys & Davidson 1979). Stars close to this boundary are unstable and are collec64

65

L. J. Smith: Ring nebulae around LBVs and WR stars

1

i

1

6.5

I

i

1

I

1

I

J " -

\ \ .60

\

127

1 1 1

R

" .120

6 00 O

I

•r) Car •y

J

1

AG Car

+

-

f He 3-519 .

5.5 -

I

\

: i

0 \* i 1

|

4.5

R 71

+ late I

I

I

,

log T,eft

|

I

I

WN

-

I

3.5

Fig. 1. The upper part of the H-R diagram showing the positions of some Galactic and LMC LBVs: the data have been taken from the review of Wolf (1992) for ij Car, R127, P Cyg, S Dor and R71; and from Smith et al. (1994) for AG Car and He 3-519. The positions of some Galactic late WN stars (Crowther et al. 1993) are also shown. The heavy line represents the Humphreys-Davidson limit, and the dashed line, the zero age main sequence with initial masses of 40-120 M© from Maeder (1990).

tively known as LBVs. In Fig. 1, the position of the Humphreys-Davidson limit is shown on the H-R. diagram, and the positions of several well-known LBVs in the Galaxy (rj Car, P Cygni and AG Car) and LMC (S Dor, R.127, R71) are indicated. LBVs are evolved, very luminous, unstable supergiants. They are characterised by irregular photometric (0.5-2mag) and spectral variations. The well-studied LBV AG Car appears to undertake excursions across the upper luminosity boundary approximately once a decade. At minimum, it has a hot supergiant spectrum and at maximum, it brightens in V by 2 magnitudes and resembles an A supergiant. In Fig. 1 the maximum and minimum states of some LBVs are indicated. They appear to move at constant bolometric luminosity and at least for AG Car, the mass loss rate also seems to be constant. The properties of LBVs have most recently been reviewed by Humphreys (1993) and Crowther k Willis (1993). In evolutionary terms, LBVs represent a short, unstable phase in the evolution of a massive 0 star (Mjnit.iai > 60 M 0 ) to a helium-burning WR

66

L. J. Smith: Ring nebulae around LBVs and WR stars Table 1. Parameters of some Galactic WR and LBV Nebula

Ml-67 RCW58 S308 NGC 2359 AG Car He 3-519

Central Star WN8 WN8 WN5 WN4 LBV LBV

Nebula Type E

E+W E+W W E E

nebulae

Diameter

Mass

(PC)

(M 0 )

Expansion Velocity (kms- 1 )

1.22 1.52 602 502 4.25 2.06

42 1 87 3 604 184 705 61 6

1

1.8 5.03 184 74 1.25 2.36

Notes: E - ejecta; W - wind-blown bubble. Refs: (1) Solf k Carsenty (1982); (2) Esteban el al. (1993); (3) Smith ei al. (1988); (4) Chu ei al. (1983); (5) Nota el al. (1992); (6) Smith et al. (1994).

star. In Fig. 1, the positions of some of the most luminous late VVN stars are shown (Crowther et al. 1993) — these are WR. stars which still have a small amount of hydrogen left and may be the recent descendants of LBVs. If the initial mass of the 0 star is « 25-60 M©, then evolution first proceeds to the red. The hydrogen-rich layers are lost in a. red supergiant (RSG) phase before the star turns back to the blue as a WR star (Chiosi & Maeder 1986). Approximately 160 WR stars are known in this Galaxy (van der Hucht et al. 1981). The atlas of Chu et al. (1983) lists 15 probable and 3 possible Galactic WR ring nebulae. More recently, CCD surveys have been completed for the northern (Miller & Chu 1993) and southern hemispheres (Chu et al. 1993). These new surveys have higher sensitivity and spatial resolution than the earlier photographic work, resulting in the identification of 3 new probable and 4 possible WR ring nebulae. Ring nebulae have also been discovered around WR. stars in the LMC (Lozinskaya, Dopita & Chu, these proceedings) and in M33 (Drissen et al. 1991). In Table 1, diameters, masses and expansion velocities are given for some representative Galactic WR nebulae. The diameters and masses cover a wide range of 2-20 pc and 1-60 M 0 . There are four well-studied Galactic LBVs: ;/ Car, AG Car, HR Car and P Cyg plus three recent additions: WRA 751 (Hu et al. 1990), HD 160529 (Sterken et al. 1991) and He 3-519 (Smith et al. 1994). Nebulae associated with P Cyg (Johnson et al. 1992), HR Car (Hutsemekers & Van Drom 1991a), and WRA 751 (Hutsemekers & Van Drom 1991b, de Winter

L. J. Smith: Ring nebulae around LBVs and WR. stars

67

Table 2. Mass loss rates (M), wind velocities (v^), wind kinetic energies (Ek), lifetimes (T), total kinetic energies (EIOT) and total mass return (M^OT) to the ISM for various evolutionary phases Phase

0 RSG LBV WR

M 210" 6

no- 55

410" 510"5

Ek

(km?"1) (ergs"1) 2500 20 250 2500

no

33

810 3 5 1 10 38

(yr) 4 106 310 5 310 4 510 5

i?TOT (erg)

MfOT (M 0 )

510 5 0

8 3 1 25

no 4 6

710 47 210 51

et al. 1992) have been discovered recently. Tims with the long-established nebulae surrounding ?? Car, AG Car and He 3-519, all the Galactic LBVs (except HD 160529) have associated nebulae. Typical diameters and masses for these LBV nebulae are 1-2pc and 2-4 Mg, (Table 1), much lower than most WR nebulae. These parameters, together with historical evidence for large eruptions in P Cyg and i] Car, suggest that LBV nebulae are formed primarily of stellar ejecta.

3 Formation of LBV and WR nebulae The formation and composition of the nebulae surrounding LBVs and WR stars are complex. Following the evolution of the central star as sketched above, the O star wind will initially evacuate the surrounding ISM. Then, depending on the initial mass of the O star, there are the denser winds of the RSG and LBV phases to consider, before finally, the powerful WR. wind interacts with the circumstellar material. In addition to this, bulk ejections of stellar material can occur during the LBV phase. In Table 2 typical mass loss rates, wind velocities and lifetimes of the various evolutionary phases are given. The lifetime of the LBV phase has been estimated by Humphreys (1991) from the LBV/WR ratio in the LMC and remains very uncertain. The parameters for RSGs are taken from Jura. (1991); the mass loss rate in Table 2 is only representative since it can vary by a factor of 100 between individual stars. Also included in Table 2 are the wind kinetic energies, the total kinetic energy and the total mass returned to the ISM by winds for each phase. It can be seen that the WR. phase is by far the most important in terms of the amount of material returned to the ISM through stellar winds. The bulk ejection of material during the LBV phase is not, however,

68

L. J. Smith: Ring nebulae around LBVs and WR stars

included in Table 2; Humphreys (1991) estimates that over the LBV lifetime > 5 MQ may be returned to the ISM through violent eruptions. The material returned to the ISM by winds and ejecta during the LBV, RSG and WR phases will be chemically enriched. Thus, from the above considerations, WR ring nebulae can be composed of (a) swept-up interstellar gas (wind-blown bubble); (b) a mixture of stellar ejecta (LBV) or swept-up RSG wind and swept-up interstellar gas; or (c) pure stellar ejecta. Chu et al. (1983) first sought to classify Galactic WR nebulae according to their morphology and internal motions. They identified ejecta-type and wind-blown bubbles as well as HII regions. To differentiate clearly between the three possibilities listed above requires abundance analyses. In an abundance study of 12 Galactic WR nebulae, Esteban et al. (1992) found that 4 were chemically enriched and formed almost entirely of stellar ejecta, 2 were a mixture of ejecta and interstellar gas, and the remaining 6 had IIII region abundances — of these 4 were classified as wind-blown bubbles and 2 as quiescent HII regions. Table 1 lists the parameters of some of these different types of WR nebulae. In the case of LBV ring nebulae, they are expected to be formed predominantly of stellar ejecta. The few abundance analyses of LBV nebulae that have been performed (e.g. AG Car: Mitra & Dufour 1990, de Freitas Pacheco et al. 1992; P Cyg: Johnson et al. 1992; 77 Car: Davidson et al. 1986) indicate anomalous abundances, consistent; with stellar ejecta (Sect. 4).

4 Applications LBV and WR ring nebulae are useful tools for the study of a variety of astrophysical problems. They can be used to investigate the dynamical interaction of winds and ejecta with the ISM, and in particular, to study clumps, radial filaments and shock structures. One example is the study of the WR nebula RCW 58 (Smith et al. 1988; Arthur et al. 1993) in which the hot, shocked WR wind interacts with cold clumps of stellar ejecta immersed in a wind-blown bubble. Smith et al. (1988) showed that the mass-loading of the stellar wind flow (by the ablation of the cold clumps) modifies the evolution of the wind-blown bubble as it will lead to enhanced cooling through the mixing-in of the clump material. Flows in clumpy circumstellar media are reviewed by Dyson (these proceedings) and Garcia-Segura & Mac Low (these proceedings) discuss hydrodynaniical models of RCW 58 . Dufour (these proceedings) presents deep narrow-band CCD images of some WR nebulae. Such images, particularly those in [0 III] tracing the shock struc-

L. J. Smith: Ring nebulae around LDVs and WR stars ture, can be used as direct probes of the interaction of the fast WR wind with the ambient ISM. Dynamical analyses can also provide information on the evolutionary history of the central star. For example, Robberto et al. (1993) have used a dynamical model with the observed parameters of the AG Car nebula to derive the mass loss rate and wind velocity of the pre-LBV phase which are consistent with that of a yellow supergiant. Viotti et al. (1993) have detected a faint HII halo around the AG Car nebula which they suggest is the wind from a previous evolutionary phase. LBV and WR nebulae can be used to determine physical parameters of the central stars. Since these nebulae are photoionised by their central stars, accurate stellar effective temperatures and luminosities can be obtained by modelling the observed ionisation structure of the nebulae. In a study of 8 WR nebulae, Esteban et al. (1993) have determined fundamental parameters using non-LTE WR model atmosphere flux distributions as input to the photoionisation models. Their temperatures agree well with those of stellar emission line analyses for the hotter WR stars but. are lower for the cooler late WN stars. They suggest that the reason for this discrepancy is the lack of line-blanketing in the WR model atmospheres. Abundance analyses of ejecta-type WR and LBV nebulae can be used to probe the chemical processes operating during the evolution of the central star. Esteban et al. (1992) have shown that ejecta-type WR nebulae have enhanced N and He, and 0 deficiencies (log N/0= +0.2-0.5 compared with the Orion value of —1.0). By comparing these abundances with the stellar evolutionary models of Maeder (1990), they find that they correspond to stars near the end of the RSG phase, with initial masses in the range 25-40MQ. Abundance studies of LBV nebulae are more difficult because the nebulae are often highly reddened and are of lower excitation, making electron temperature (Te) determinations uncertain. Mitra & Dufour (1990) analysed the AG Car nebula and found that 0 and S were overabundant while N was normal for Te = 9 000K. On the other hand, de Freitas Pacheco et al. (1992) found that N was overabundant by an order of magnitude and that 0 was deficient by at least a factor of 6 for Te = 12 400 K. For ejecta associated with i] Car, Davidson et al. (1986) find that most of the CNO material has been processed to nitrogen. Abundances studies of the nebulae surrounding P Cyg (Johnson et al. 1992), IIR. Car and WRA 751 (Hutsemekers & Van Drom 1991a,b) are uncertain because the crucial temperature diagnostic line [Nil] A5755 was not detected. Nevertheless, N/S ratios, which are relatively insensitive to Te, indicate N overabundances. In addition, nebular [01], [Oil] and [0 III] lines are not detected in these three

69

70

L. J. Smith: Ring nebulae around LBVs and WR stars

nebulae, indicating high 0 deficiencies. These abundance studies, taken overall, suggest that we are directly observing material that has been processed via the ON cycle (in which oxygen is directly converted to nitrogen) in massive stars. Another area of research where studies of LBV and WR nebulae can make an important contribution is the subject of asymmetric winds, through for example, the detection of bipolarity. Coronographic observations of the AG Car nebula (Paresce k Nota 1989; Nota et al. 1992) revealed a bipolar structure with a dusty jet-like feature extending from close to the central star to the nebula. Spectropolarimetry of AG Car itself (Schulte-Ladbeck et al. 1993a) shows that the wind is strongly asymmetric in a direction perpendicular to the "jet". Likewise, coronographic observations of the nebula surrounding the LMC LBV R.127 (Clampin et al. 1993; Nota, these proceedings) show that it has bipolar lobes. Again, spectropolarimetry (SchulteLadbeck et al. 1993b) reveals that the wind from R.127 is highly asymmetric with a position angle approximately perpendicular to the two bright nebular lobes. HST images of i] Car and the surrounding Homunculus nebula (Hester et al. 1991; Ebbets et al., these proceedings) reveal that the nebula, is a thin, clumpy shell, as well as showing that there is a stellar jet to the NE of 77 Car. Since major eruptions in LBVs (e.g. 1] Car in 1840) are rarely directly observable, it is important to obtain estimates of the masses and ages of the stellar ejecta clumps since they give valuable clues to the actual ejection process.

5 The connection between LBV and WR nebulae Since the most massive 0 stars evolve to WR stars via a short-lived LBV phase, it is quite possible that some LBV nebulae are the precursors of WR ejecta-type nebulae. Indeed, as mentioned in the previous section, these two types of nebulae share N overabundances and 0 deficiencies, although strictly comparable abundance analyses have yet to be performed. Miller & Chu (1993) have investigated the detectability of ejecta in terms of its emission measure and find that such nebulae are only visible for ~ 104yr which is small compared with the lifetime of the WR. phase (Table 2). Thus the remnants of stellar ejecta associated with the precursor LBV phase may have simply become too faint to detect. In this context, it is interesting to note that the two ejecta-type WR nebulae (Ml-67, RCW 58; Table 1) which have comparable sizes and masses to the AG Car nebula, have WN8 central stars which are relatively young WR stars with some atmospheric hydrogen.

L. J. Smith: Ring nebulae around LBVs and WR stars

71

The distinction between LBVs and WR stars (and thus their nebulae) has been recently questioned by Smith et al. (1994). They have carried out a detailed analysis of the LBV candidate He 3-519 and AG Car at minimum and find that they can best be described spectrally as cool WN11 stars. Moreover, an atmospheric analysis shows that He 3-519 and AG Car have low H/He ratios of 1.8 and 2.4 which are similar to some late WN stars. Smith et al. (1994) therefore suggest that these two LBVs have already evolved to the WN phase, and thus their nebulae should be properly classified as WR ejecta-type nebulae. Overall, the dynamics and abundances of LBV nebulae are less well known than those of WR, nebulae, mainly because most LBV nebulae have only been discovered over the last few years. Future work should therefore aim at improving this situation. A comparative abundance study of LBV and WR ejecta-type nebulae would be very worthwhile. It is also of interest to search for faint halos around LBV and WR nebulae; abundances and expansion velocities of these remnant winds have the potential of directly identifying the type of precursor star.

References Arthur, S.J., Dyson, J.E. k. Hartquist, T.W. (1993). Mon. Not. R. astr. Soc, 261, 425. Chiosi, C. & Maeder, A. (1986). Ann. Rev. Astr. Astrophys. 24, 329. Chu, Y.-H. (1991). In IAU Symp. No. 143: Wolf-Rayet Stars and Interrelations with Other Massive Stars in Galaxies, ed. K.A. van der Hiicht & B. Hidayat (Kluwer, Dordrecht), 349. Chu, Y.-H., Treffers, R.R. & Kwitter, K.B. (1983). Astrophys. J. Suppl, 53, 937. Chu. Y.-H, Garcia-Segura, G., Dopita, M.A., Bell, J.F., Lozinskaya, T.A., Marston, A.P. & Miller, G.J. (1993). In: Massive Stars: Their Lives in the Interstellar Medium, ed. J.P. Cassinelli & E.B. Churchwell, Astr. Soc. Pac. Con}. Ser. 35, 360. Clampin, M., Nota, A., Golimowski, D.A., Leitherer, C. fc Durrance, S.T. (1993). Astrophys. J., 410, L35. Conti, P.S. (1984). In IAU Symp. No. 105: Observational Tests of the Stellar Evolution Theory, ed. A. Maeder & A. Renzini (Kluwer, Dordrecht), 233. Crowther, P.A. &: Willis, A.J. (1993). In: Evolution of Massive Stars: A Confrontation between Theory and Observations, ed. D. Vanbeveren, W. Van Rensbergen fc C. de Loore, Space Sci. Rev., in press. Crowther, P.A., Smith, L.J. & Hillier, D.J. (1993). In: Evolution of Massive Stars: A Confrontation between Theory and Observations, ed. D. Vanbeveren, W. Van Rensbergen &; C. de Loore, Space Sci. Rev., in press. Davidson, K., Dufour, R.J., Walborn, N.R. k. Gull. T.R. (1986). Astrophys. J., 305, 867. Drissen, L., Shara, M.M. & Moffat, A.F.J. (1991). Astron. ./., 101, 1659. Esteban, C , Vflchez, J.M., Smith, L.J. & Clegg, R.E.S. (1992). Astr. Astrophys., 259, 629. Esteban, C , Smith, L.J., Vilchez, J.M. & Clegg, R.E.S. (1993). Astr. Astrophys., 272, 299. de Freitas Pacheco, J.A., Damineli Neto, A., Costa, R.D.D. & Viotti, R. (1992). Astr.

Astrophys., 266, 360.

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L. J. Smith: Ring nebulae around LDVs and WR stars

Hester, J.J., Light, R.M., Westphal, J.A., Currie, D.G., Groth, E.J., Holtzman, J.A., Lauer, T.R. k O'Neil, E.J. (1991). Astron. J., 102, 654. Hu, J.Y., de Winter, D., The, P.S. k Perez, M.R. (1990). Astr. Astrophys. 227, L17. van der Hucht, K.A., Conti, P.S., Lundstrom, I. k Stenholm, B. (1981). Space Sci. Rev., 28, 227. Humphreys, R.M. (1991). In IAU Symp. No. 143: Wolf-Rayet Stars and Interrelations with Other Massive Stars in Galaxies, ed. K.A. van der Hucht k B. Hidayat (Kluwer, Dordrecht), 485. Humphreys, R.M. (1993). In: Massive Stars: Their Lives in the Interstellar Medium, ed. J.P. Cassinelli k E.B. Churchwell, Astr. Soc. Pac. Conf. Ser. 35, 179. Humphreys, R.M. k Davidson, K. (1979). Astrophys. J., 232, 409. Hutsemekers, D. k Van Drom, E. (1991a). Astr. Astrophys., 248, 141. Hutsemekers, D. k Van Drom, E. (1991b). Astr. Astrophys., 251, 620. Johnson, D.R.H., Barlow, M.J., Drew, J.E. k Brinks, E. (1992). Mon. Not. R. astr. Soc, 255, 261. Jura, M. (1991). In IAU Symp. No. 143: Wolf-Rayet Stars and Interrelations with Other Massive Stars in Galaxies, ed. K.A. van der Hucht k B. Hidayat (Kluwer, Dordrecht), 341. Maeder, A. (1990). Astr. Astrophys. Suppl., 84, 139. Miller, G.J. k Chu, Y.-H. (1993). Astrophys. J. Suppl.. 85, 137. Mitra, P.M. k Dufour, R.J. (1990). Mon. Not. R. astr. Soc, 242, 98. Nota, A., Leitherer, C , Clampin, M., Greenfield, P. k Golimowski, D.A. (1992). Astrophys. J., 398, 621. Paresce, F. k Nota, A. (1989). Astrophys. J., 341, L83. Robberto, M., Ferrari, A., Nota, A. k Paresce, F. (1993). Astr. Astrophys., 269, 330. Schulte-Ladbeck, R.E., Clayton, G.C. k Meade, M.R. (1993a). In: Massive Stars: Their Lives in the Interstellar Medium, ed. J.P. Cassinelli k E.B. Churchwell, Astr. Soc. Pac. Conf. Ser. 35, 237. Schulte-Ladbeck, R.E., Leitherer, C , Clayton, G.C, Robert, C , Meade, M.R. (1993b). Astrophys. J., 407, 723. Smith, L.J. (1990). In: New Windows on the Universe, ed. F. Sanchez k M. Vazquez (Cambridge University Press, Cambridge), Vol. II, 453. Smith, L.J., Pettini, M., Dyson, J.E. k Hartquist, T.VV. (1988). Mon. Not. R. astr. Soc, 234, 625. Smith, L.J., Crowther, P.A. k Prinja, R.K. (1994). Astr. Astrophys., 281, 833. Solf, J. k Carsenty, U. (1982). Astr. Astrophys., 110, 54. Stahl, O. (1989). In IAU Coll. No. 113: Physics of Luminous Blue Variables, ed. K. Davidson, A.F.J. Moffat k H.J.G.L.M. Lamers (Kluwer, Dordrecht), 149. Sterken, C , Gosset, E., Jiittner, A., Stahl, O., Wolf, B. k Axer. M. (1991). Astr. Astrophys. 247, 383. Viotti, R., Polcaro, V.F. k Rossi, C. (1993). Astr. Astrophys.. 276, 432. de Winter, D., Perez, M.R., Hu, J.Y. k The, P.S., (1992). Astr. Astrophys., 257, 632. Wolf, B. (1992). In Nonisotropic and Variable Outflows from Stars, ed. L. Drissen, C. Leitherer k A. Nota, Astr. Soc. Pac. Conf. Ser. 22, 327.

The Interstellar Environment of Wolf-Rayet Stars in the LMC: New Survey and Statistics Tatiana A. Lozinskaya1, Michael A. Dopita 2 and You-Hua Chu 3 1

Sternberg Astronomical Institute, Moscow Lomonosov University, Russia Mt. Stromlo and Siding Spring Observatories, The Australian National University, Australia 3 Department of Astronomy, University of Illinois at Urbana-Champaign, USA 2

In order to understand the evolution of Wolf-Rayet (WR.) stars and their interaction with the surrounding circumstellar and interstellar gas we have undertaken an emission-line imaging survey of the (almost complete) WR star population in the Magellanic Clouds (Dopita. et of. 1994). Interference filter CCD images have been obtained in Ha and [0 III] A5007 for all WR stars in the LMC and the SMC. The survey was conducted using the 2.3 m telescope at the Siding Spring Observatory ANU. The field of view was 6'.7, and the pixel size was 0".65/pix. A total of 115 WR. stars in the LMC (Breysacher 1981; Lortet, 1991) and 9 WR stars in the SMC (Azzopardi & Breysacher, 1979; Morgan et al. 1991) were observed in this survey. This survey is the first complete survey of the ionized material around WR stars in the Magellanic Clouds, and indeed is the first complete survey in any galaxy. We have almost doubled the number of ring nebula known in the MCs, and have revealed a number of cases in which stellar ejecta has almost certainly been identified. As a consequence, we find that the incidence of ring nebulae around WR stars in the LMC is very similar to that in the solar neighborhood. (According to Lozinskaya, 1982; 1983; 1992 only 30-40% of WR. and Of stars in the distance-complete sample in the Galaxy are associated with ring nebulae; the nebula types of stellar ejecta. and wind-blown bubble are even more scarce, about. 10-15%. New deep surveys of the environments of WR stars in the Galaxy provided by Miller & Chu (1993) confirm theses statistics.) Since in terms of classical theory every WR star should be surrounded by at least two shells: a large "interstellar" shell swept up during the Main Sequence and a small bubble of circumstellar or interstellar gas swept up by the strong wind during the WR. stage, the following questions require answers. (i) Why don't we see ring nebula, around every WR star? In other words, are ring nebulae short-lived or seldom-created? 73

74

T. A. Lozinskaya et al.: WR stars in the LMC

(ii) Do all WR stars lose matter in the form of stellar ejecta? If not, then which spectral classes do and which do not? Several physical mechanisms that prevent the formation of a wind-blown bubble or decrease its lifetime have been considered so far. Among them there is a simple one: if the wind sweeps up not the interstellar gas, but stellar ejecta, then the lifetime of a shell may be shorter than the WR stage. The explanation seems reasonable since winds of the progenitor star and/or of the parent OB-association may evacuate the surrounding gas. We consider the matter using our new deep search and take into account the detectability of ejected and swept up shells. (Similar consideration have been recently provided by Miller and Chu (1993) for the Galaxy.) To evaluate the detectability of stellar ejecta. we take its typical mass and expansion velocity from observations of well studied ejecta in the Galaxy. For a limiting emission measure of 10 cm~f> pc, a 5 - 10 MQ ejected shell would fade away at a. radius of 4-7 pc if the shell's geometrical thickness A r / r ~ 0.1 - 0.2. For an expansion velocity of 50 - 100 km s" 1 the duration for detectability is only a few x 104 yr, which is about 10% of the WR. phase lifetime. For the same limiting emission measure and shell thickness, the minimal radius of a detectable swept interstellar shell corresponds to 100 - 200 pc for an ambient density 7?n = 0.1 cm" 3 or to 5 - 10 pc for no = 0.5 cm" 3 . Therefore, first, in a low-density area we should see ejecta-type nebulae around only about 10% of WR stars even if every one has ejected a shell in the course of its evolution. (In the dense ambient gas, the ejected and swept shells should merge.) And, second, we can try to discriminate ejected and swept up interstellar shells around WR. stars in low-density regions. To try to make this distinction we selected 31 WR stars in the LMC which appear to be located in low density ambient gas. The criteria we used for the selection were as follows: (i) (ii) (iii) (iv)

Low Hi-column density N(HI)< 1021crrr2, according to Rolilfs et al. (1984) Low color excess E(B-V) < 0.08 according to Breysacher (1986) Location outside bright CO clouds as delineated by Thaddens (1987) Location outside bright HI] complexes as delineated by Davies et al. (1986)

The results of the search for ring nebulae associated with these "selected" WR stars are as follows: Large faint ring nebulae of a typical size from 50 to 300-400 pc are found to be related to 64% of "selected" WR. stars. These big rings most probably display bubbles blown by the winds at Main Sequence and by the common winds of parent OB associations.

T. A. Lozinskaya et al.: WR stars in the IMC

75

Table 1. Statistics of big shells related to "selected" WR stars. "Selected" WRs N° "selected" WRs N° big rings around "selected" WRs Percentage of "selected" WRs with big rings

in LH ass.

nearby LH ass.

no OB ass.

10

4

17

31~~

10

2

8

20

100%

50%

47%

65%

total

Table 2. Statistics for all Br WR stars. All WR stars N° WRs N° big rings around WRs Percentage of WRs with big rings

in LH ass.

nearby LH ass.

no OB ass.

total

35 30

14 6

30 6

79 42

86%

43%

20%

53%

Winds of associations seem to dominate for the majority of big rings. Indeed, in Table 1 we show the number and percentage of "selected" WR stars with big shells separately for stars which belong to Lucke and Hodge (1970, hereafter LH) associations; for stars which are located nearby (less than 5' from the edge as delineated by LH) and for stars which neither belong to nor are in vicinity of any OB association. Table 2 shows the same statistics for all WR. stars, excluding those located in the overcrowded 30 Dor area. One can see that the majority of WR stars which belong to OB associations are related to big shells; the percentage of big shells around stars outside associations is less than 50%. Small ring nebulae are found to exist around eleven of the "selected" WR stars, all of them bar one inside big swept shells. The inner "small rings" might probably display circumstellar material. Tables 3 and 4 show the distribution of the small WR-rings between spectral types for "selected" WR. stars and for all WR stars in the LMC (excluding members of LH 100 in the core of 30 Dor).

76

T. A. Lozinskaya et ai: WR stars in the IMC Table 3. Statistics of the small WR-rings related to "selected" WR stars of different spectral types. spectral type

N(WNE)

N(WNL)

N(WN8)

N(WC)

19

5

2

4

7

1

1

2

37%

20%

50%

50%

N° "selected"WRs N° small rings around "selected" WRs Percentage of "selected" WRs with small rings

Table 4. Statistics of small WR-rings related to all WR stars of different spectral types. spectral type N° WRs N° small rings around WRs Percentage of WRs with small rings

N(WNE)

N(WNL)

N(\VN8)

N(WC)

58

18

3(4)

21

16

2

2

4

28%

11%

67 (50)%

19%

Though the sample is poor for WN8 stars and for "selected1' WC stars, three facts seem to be meaningful: (i) the scarcity of small rings around WNL stars, confirming data by Chu for the Galaxy, (ii) the low percentage of WC stars with small ring nebulae (again like in the Galaxy), (iii) the very high percentage of WN8 stars with small ring nebulae. The only WN8 nebula in the sample of "selected" stars is of Ejecta type. This third result confirms the previous data for the Galaxy and most probably indicates that first, these stars do differ from other WNL types and second, they do eject shells. The sizes of small ring nebulae related to "selected" stars of different spectral types are as follows. Ring nebulae around WNE stars have sizes in the range 3 to 30 pc; two WC ring nebulae are among the largest (sizes 30-45 pc); and two stars WN7 and WN8 are among the smallest (less then 10 pc). The statistics are poor (size strongly depends on the ambient density and we used only "selected" stars), though it is consistent with an evolutionary

T. A. Lozinskaya et a/.: WR stars in the LMC

77

sequence from the WNL classes to the WC types, as suggested by theory. The wide range of sizes of WNE-rings probably reflects different ways of evolution in the LMC: According to Breysacher (1986) both young and old stars belong to WNE types. Therefore, the statistical evidence for the LMC does not contradict the suggestion that all WR stars may eject a stellar shell of a typical mass about 5 - 1 0 M 0 at a velocity of about 50 to 100 km s"1; the strong stellar wind sweeps-up the ejected material. However, since all these estimates are very uncertain we'll refrain so far from making more definite statements. References Azzopardi, M. fc Breysacher, J. 1979, Astr. Astrophys., 75, 120. Breysacher, J., 1981, Astr. Astrophys. SuppL, 43, 20:5 . Breysacher J., 1986, Astr. Astrophys., 160, 185. Chu, Y.-H. 1982, Astrophys. J., 254, 578. Chu, Y.-H., Treffers, R. R. fc Kwitter. K. B. 1983. Astrophys. J. SuppL, 53, 937. Chu, Y.-H. 1981, Astrophys. J., 249, 195. Chu, Y.-H. 1991, in IAU Symposium 143, Wolj-Rayet Stars and Interrelations with Other Massive Stars in Galaxies, eds. K. A. van der Hiicht fc B. Hidayat, Dordrecht: Kluwer), p. 349. Davies, R.D., Elliott, K.H. fc Meaburn, J., 1976, Mem. R. Astron. Soc, 81, 89. Dopita, M.A., Bell J.F., Chu Y.-H. fc Lozinskaya T.A., 1994 Astrophys. J. SuppL, (in press). Lortet, M.-C. 1991, in IAU Symposium 143, Wolj-Rayet Stars and Interrelation with Other Massive Stars in Galaxies, eds. K.A. van der Hiicht. and B. Hidayat, Dordrecht: Kluwer, p.513. Lozinskaya, T.A., 1982, Astrophys. and Space Sci., 87, 313. Lozinskaya, T.A., 1983, Pis 'ma Astron. Zh., 9, 469. [Sov. Astron. J. Letters], 9, 247. Lozinskaya, T.A., 1992, "Supernovae and stellar wind in the Interstellar Medium", (Second Edition): American Institute of Physics, New York. Lucke, P.B. & Hodge, P.W., 1970, Astrophys. J., 75, 171. Miller G.J. k. Chu Y.-H., 1993, Astrophys. J. SuppL, 85, 137. Morgan, D.H., Vassiliadis, E. & Dopita, M.A. 1991, Mon. Not. R. astr. Soc, 251, 51P. Rohlfs K., Kreitschmann J., Siegman B.C. & Feitzinger J.V., 1984, Astr. Astrophys., 137, 343. Thaddeus P., 1987, "A CO survey of the LMC" Lecture Notes in Physics, 306 The other Galaxy, p. 241-246

Morphology & Physical Conditions in WR Shell Nebulae Reginald J. Dufour1'2 1

Department of Space Physics & Astronomy, Rice University, Houston, Texas, USA 77251-1892 2 Investigador Visitante, Instituto de Astronomia, UNAM, A.P. 70-264 Cd. Universitaria 04510 Mexico, D.F.

1 Introduction Wolf-Rayet Shell Nebulae (WRSN) provide a "quick look" at an intermediate stage of evolution of massive stars between the main sequence 0 stage and their ultimate demise as SNII. During this evolutionarily brief epoch, the 0 star develops a strong wind which affects the surrounding ISM, and can even have significant mass loss which enriches the ISM with H-burning products —specifically He and N (Maeder 1990). Therefore, studies of these objects are both interesting and important regarding the physics of windshock effects on the ISM and in the role they have in galactic chemical evolution. In this short contribution I will present some of the results of two recent students of mine who completed Ph.D. theses studying the morphology and spectra of the WRSN NGC 6888 (Mitra 1990) and NGC 2359 (Jernigan 1988). A more comprehensive review of the literature on WRSN is given by the fine paper by L. Smith in this volume. The theses studies incorporated CCD imagery mapping of the ionization structure of the nebulae in the emission lines of H/3, [OIII]A5007, Ha, [NII]A6583, & [SII]A6717+30; followed by spectroscopy of parts of the two nebulae that were of special interest from the imagery. Herein I will note some of the spectroscopic results regarding the hot wind-driven gas; the imagery mapping is available in their theses and moreso in a forthcoming Atlas of CCD Imagery of Galactic HII Regions (Hester et al. 1994).

2 NGC 6888 NGC 6888 is the prototype of the "wind blown" WRSN (Chu 1981). Our imagery and spectroscopy indicates that it also contains 5M e of ionized material, enriched in He and N, and therefore a composite "wind+ejecta" classification is more appropriate. This is supported by the two photographs on the next page: Figure 1 shows the nebula in Ha and Figure 2 shows the 78

R. Dufour: WR shell nebulae

79

Fig. 1. CCD image of NGC 6888 in Ha. Field diameter is 16'.

same field in [OIH]/Ha+[NII] (from Mitra 1990). For an adopted distance of 1.45kpc, the size of the nebula in Ha is 7.6x5.Ope, with the WN6 star (HD 192163) offcenter to the NW of the ellipsoid (in Ha). However, with respect to the boundaries of [OIII] (Fig. 2), the nebular perimeter is almost circular with the WN6 star centered. Therefore, while the nebula is bipolar in the "mass-loaded" stellar ejecta (best seen in Ha and [Nil]), it is spherical in the wind-driven bubble (as indicated from [OIII]). Spectra of 30 positions in NGC 6888 were taken in 1989 July & Sept. using the IRS on the KPNO 0.91m telescope. The spectra covered 36006800A with a resolution of ~12A and were taken simultaneously through two rectangular 35.6"x7.8" apertures separated by 61.2". Temperatures were derived from [OIII] & [Nil] lines for 12 knots in NGC 6888, as well as densities from [SII]. For the knots, the T e 's from [Nil] range from 650011,500 K and 11,900-29,100 K from [OIII] - with the [OIII] T e always higher for the knots where both ions could be observed. In addition, spectra were taken of several of the [OIII] rims or bubbles to the NE and N of the WN6 star which showed very high excitation and surprisingly high T e 's, with T e =55,000±20,000 K for the [OIII] rim (called NO3) on the NE edge of the nebula. Figure 3 shows the spectrum covering 3600-6800A. The ratio

80

R. Dufour: WR shell nebulae

Fig. 2. [OIII]A5007/Ha+[NII] ratio map of NGC 6888 showing the windblown bubble boundaries (dark = strong [OIII]). of [OIII]A5007/H/3 is ~20 and the line near 4350A is ~80% [OIII]A4363 (blended with H7) The spectra of NO3 and several other strong [OIII] features in NGC 6888 indicates that the entire nebula is filled with hot wind-shocked low density gas of T e ~50,000 K. While the existence of such a medium in WRSN comes as no surprise given theoretical expectations, the study of Mitra (1990) was the first successful measurement of pure wind-shock spectra in a WRSN. The ramifications of this medium is significant in interpreting the spectra of the stellar ejecta knots of NGC 6888 since the wind-shocked medium affects the observed spectra and derived physical properties of the higher ionization species. Using a "two-zone" model for the spectra of the knots seen in Ha and [Nil], Mitra demonstrated that the wind-shock medium contributes 1436% of the emission in [OIII]A5007 seen in the knots, as well as most of the [OIII]A4363 emission. When the wind-shock contributions are accounted for, the [OIII] T e of the knots become ~9000 K, comparable to the [Nil] T e observed and consistent with the knots being photoionized by the WN6 star. Even when the wind-shock emissions contaminating the knots' spectra are removed, abundances of He-N-0 for the knots in NGC 6888, derived

R. Dufour: WR shell nebulae

81

NGC6888-H03 5.00E-15

1

4.00E-15

-

1

•

-

3000s

-

1

1

1

1

1

3.00E-15

-

2.00E-15

-

I.00E-15

lib 1 .

1 lAdhiyjMiMilllnPl \ur

0

1 3500

1 4000

1 4500

5000

1

1

1

5500

6000

6500

Fig. 3. Spectrum of the shock on the NE rim of NGC 6888.

from both photoionization models and direct emission line diagnostics, are peculiar. For the eight best observed knots in NGC 6888 (those for which T e 's were derivable from both [OIII] & [Nil] and Ne from [SII]), average elemental abundances in terms of 12-f [X/H]±l
82

R. Dufour: WR shell nebulae Morphological Representation of NGC 2359

NGC 2359

im fcloniaaonSRUnen

i

«

Ouinpy. Photoiowied Mitenil

. Filunenuiv Emiiiicn

Expudinj Shock Front

7 kmiuiioo Shidowing

Till

Fig. 4. Cartoon of the structure of NGC 2359 (left) and a grey-scaled map of [OIII]A5007/Ha ratio (right, strong [OUT] is white). "ambient gas" showed no abundance abnormalities, which suggests to the author that this component is photoionized ISM outside the nebula - n o t the hot wind-shocked gas inside the nebula (which would likely have a very high velocity dispersion and broad [OIII] lines).

3 NGC 2359 NGC 2359 can be considered possibly as the prototype of "pure" wind-driven bubble nebula. It is formed by the wind of a WN4 star which had previously produced a blister HII region on the face of a molecular cloud which is seen almost edge-on to the east of the WRSN. The remarkable intricate structure and beauty of this wind-driven bubble and associated HII region was first demonstrated in image-tube emission-line photographs by Schneps et al. (1981) and subsequently quantified by the CCD imagery in the thesis of Jernigan (1988). Figure 4 shows a diagram of the structure of NGC 2359 and the [OIII]/Ha grey-scaled ratio image from Jernigan's thesis. Longslit spectra of the wind-driven bubble of NGC 2359 were obtained at CFHT with a 2.5"xl30" slit centered 10"N and 50"W of the WN4 star. The

83

R. Dufour: WR shell nebulae 1

E

1

1

1

S.OOE-16 -

-

5.00E-1S -

-

•1.00E-1S -

-

3.OOE-16 -

-

2.OOE-15 -

-

1.00E-16

-

0 "~

1 •1500

1 5000

5500

1 6000

1 6500

Wavelength (A) Fig. 5. Spectrum of the W rim of the wind-blown bubble of NGC 2359.

spectra cover the 4300-6800A wavelength range with a resolution of about 10A. Figure 5 shows the spectrum of a 12" wide extraction of the west rim of the bubble. It is characterized by very strong [OKI] lines, Ha, and H/3, with all other lines much weaker. Temperatures from [OIII] inside the bubble to its edge range from 21,000 K to 40,000 K. When compared to SN adiabatic shock models, the [OIII]/H/3 ratios for the bubble suggest shock velocities >>100 km s" 1 , but the [SII] lines are much weaker than predicted -due to the fact that the wind-shock is propagating through a preionized medium and the UV back illumination by the WN4 star prevents recombination behind the shock (i.e., an "incomplete" shock). This results in a higher [OIIIj/H/3 and weaker [SII]/Ha than that which corresponds to the rather slow expansion velocity measured for the bubble: Vs=30 km s" 1 . Abundances in the nebula has been studied by various investigations, most recently by Esteban et al. 1990, but most spectra have been obtained of the HII region to the E of the bubble which is contaminated in [OIII] by the wind. Possibly this has resulted in higher than appropriate Te's for the photoionized gas (~13,000 K) and lower O/H (12-f-[O/H]«8.2) abundances than "solar". There is currently no clear evidence for significant He- k N-enrichments in the bubble material.

84

R. Dufour: WR shell nebulae

4 Concluding Remarks While I have concentrated on the work of my two recent Ph.D. students on NGC 6888 and NGC 2359, the fine recent studies on these and other WRSN by colleagues in the UK, IAC, and USA have led to basic "appreciation" of the physical conditions and enriched He-N abundances in this significant evolutionary phase of massive 0 stars. Such provides important data and constraints for theoretical modeling of these phenomena, which is being developed in increased sophistication as noted by the papers of Dyson, Drew, Garcia-Segura, and others at this meeting. Acknowledgements I wish to thank Tammy Jernigan and Mila Mitra for permitting me to describe their thesis results in advance of publication. In addition, I note the many fine discussions and fruitful nights at the Palomar 1.5m telescope with Jeff Hester and Bob Parker. References Chu, Y.-H. (1981). ApJ, 249, 195. Esteban, C. & Vilchez, J. M. (1992). ApJ, 390, 536. Esteban, C , Vilchez, J. M., Manchado, A. & Edmunds, M. G. (1990). A&Ap, 227, 515. Esteban, C , Vilchez, J. M., Smith, L. J. & Clegg, R. E. S. (1992). A&Ap, 259, 629. Hester, J. J., Dufour, R. J., Parker, R. A. R. & Scowen, P. A. (1994). Atlas of CCD Imagery of Galactic HII Regions, NASA Research Publication, in preparation. Jernigan, T. E. (1988). Ph. D. Thesis, Rice University. Maeder, A. (1990). A&Ap. Sup., 84, 139. Mitra, P. (1990). Ph. D. Thesis, Rice University. Schneps, M. H., Haschick, A. D., Wright, E. L. & Barret, A. H. (1981). ApJ, 243, 184.

Three-Wind Model for Wolf-Rayet Bubbles Guillermo Garcia-Segura 1 ' 2 , Mordecai-Mark Mac Low3*1 1

University of Illinois at Urbana-Champaign Instituto de Astrofisica de Canarias 3 University of Chicago 2

1 Introduction Strong winds from massive stars can sweep up the ambient gas forming stellar wind bubbles, also called ring nebulae. Classically, ring nebulae around Wolf-Rayet (WR) stars have been modeled assuming a homogeneous interstellar medium (ISM), following Weaver et al. (1977). However, theory and observations have progressed to the point that, this simplification can no longer be justified. The evolution of massive stars has been studied by Maeder (1990). He shows tracks for 15-120 M{-0 in his plots (Figures 14). Main sequence (MS) stars between 25—40 M& evolve to WR stars after passing through a red supergiant (R.SG) phase. Observations of MS stars (Herrero et al. 1992) and WR stars (Willis 1991) reveal fast winds, as opposed to RSG stars (Humphreys 1991), where the winds are dense and slow (see also Chevalier & Liang 1989, Stencel et al. 1989). The above studies, suggest to us that the ISM initially encountered by a WR wind is far from homogeneous. This is the base of our three-wind model. In order to explain WR ring nebulae, we must take into account the history of the central stars, not just their interstellar environment. We have already presented a brief description of an analytic calculation of the dynamical behavior of the swept-up shell of RSG wind (Garcia-Segura. and Mac Low 1993). In this paper, we present numerical computations of the shell that follow it after instability sets in and it can no longer be modeled analytically.

2 Technique and simulation To study the the swept-up shell of RSG wind, we use the gas dynamical code ZEUS-3D, written by D. A. Clarke and M. L. Norman. This is the 3-D extension of the code ZEUS-2D, described by Stone & Norman (1992). ZEUS-3D is an explicit code that integrates the ideal gas equations in the absence of viscosity. We used a two-dimensional, spherical-polar, Eulerian grid, assuming azimuthal symmetry, with reflecting boundary conditions at the equator and the pole. Radiative, optically-thin cooling is used, with the plasma cooling curve given by MacDonald & Bailey (1981), an extension of 85

86

G. Garcia-Segura & M. Mac Low: Three-wind model for WR bubbles

the Raymond-Smith curve to lower temperature. The MS bubble and the RSG wind are set up analytically in the 2-D run, based on a 1-D computation, where the mass loss rate and wind velocity were time dependent. For the WR wind, we used the standard form for a thermally accelerated wind, adding energy and mass at each time step to a small source region at the origin of the grid. The RSG wind density falls off with a power law close to r~2, while the interior of the MS bubble has low, nearly constant density. The two functions are connected by a steep jump of two orders of magnitude at the boundary of the RSG wind. The RSG wind might form a shell of its own as it sweeps up the MS bubble interior. Such a shell would only enhance the instability of the WR. driven shell, so we do not include it so that we can determine whether the least unstable case is adequate to reproduce the observed morphology and dynamics. This approximation corresponds with assuming that the MS bubble has quickly cooled and so has a low-pressure interior. We found empirically that, in order to reproduce the observed blowouts, the RSG wind has to be slower at the equator, allowing the swept-up shell to arrive earlier to the equatorial boundary. To allow departure from sphericity in the absence of magnetic field, we have included an angular density variation in the RSG wind, proportional to sin.2O, where theta is the polar angle. We are interested in Rayleigh-Taylor (R-T) instabilities, or rather, to be formal, Richtmyer-Meshkov instabilities, since they are pressure-driven rather than gravitationally-driven. However, we do not perturb the sweptup shell with any particular eigenmode of the instability, since we are not studying the growth of a specific mode. Instead, we allow the perturbations to come from numerical noise and soundwaves in the bubble interior. Although dynamical instabilities are three dimensional, two dimensional runs can give us a qualitative description of their development. The evolution of the WR ring nebula is shown in Figure 1. The grid has 400 x 400 zones, corresponding to 5 pc in radius and 90 degrees in the polar direction. The boundary of the RSG wind occurs at 2 pc at the equator and at 3.6 pc at the pole. As soon as the WR wind is set up, an outer shock forms, with an expansion velocity of 80 km s"1. A contact discontinuity separates the swept-up shell from the hot interior. This interface is initially marginally R-T stable as it expands at constant velocity. The swept-up shell approaches a self-similar shape asymptotically, with some acceleration in the polar region, producing R-T instabilities. Even when the shell is marginally R-T stable, the shell can still have slow instabilities of the type described by Ryu & Vishniac (1991). The main dynamical effect occurs when the outer shock arrives at the boundary of the RSG wind. It is already somewhat

G. Garcia-Segura & M. Mac Loiv: Three-wind model for WR bubbles 87

lllllii Ilillllliill; iiiiillilfli?

••III llllllllllllllll

^^^^^^M ilNlillliliill

••IIII^^H illlllllliP isiiiiliilii

Fig. 1. Logarithm of density at 30000, 35000, 40000 and 45000 years. Levels from 102 cm~3 (black) to 10~3 cm" 3 (white). clumpy due to the above instabilities. The hot shocked wind out between the clumps, driving the local outer shock at high this point, most HQ emission comes from the fragmented shell RSG wind, while filamentary [0 III] emission comes from the bursting through the shell.

flows freely velocity. At of swept-up outer shock

3 Conclusions The three-wind model can explain the morphology of WR. ring nebulae: the shape of the bubble, clumpy or filamentary shells, and blowout features.

88

G. Garcia-Segura & M. Mac Low: Three-wind model for WR bubbles

The model can explain why the swept-up shell can show higher abundances than those seen in the ISM as reported by Esteban & Vilchez (1992), due to the fact that this is the processed gas from the RSG phase. It predicts blueshifted highly ionized gas like C IV and Si IV as reported by Smith et al. (1984), due to the high velocities in the blowouts. The model also predicts strong cooling in the hot shocked gas from two processes. First, conductive evaporation will be enhanced by the increased area of the interface. Second, the clumps of swept-up shell will be ablated by Kelvin-Helmholtz instabilities as gas blows out past them, directly mixing cold gas with hot shocked gas. This could be an explanation for the polar X-ray brightening observed in NGC 6888 (Bochkarev 1988), as a result of stronger instabilities in the polar direction (an analytic model of the X-ray brightness for stable interfaces can be found in Garcia-Segura & Mac Low 1993). Finally, this dynamical effect produces an ideal clumpy medium for supernova remnants. Acknowledgements We would like to thank Y.-H. Chu and A. Manchado for their encouragement and discussions. We thank M. Norman and the Laboratory for Computational Astrophysics for the use of ZEUS-3D. The calculations were performed on a Cray Y-MP C90 at the Pittsburgh Supercomputing Center, and visualized at the National Center for Supercomputing Applications. This work was partially supported by NASA grants NAG5-2245 and NAG5-1900. References Bochkarev, N.G. 1988, Nature, 332, 518 Chevalier, R.A., & Liang, E.P. 1989, Astrophys. J., 344, 332 Esteban, C. fc Vi'lchez, J.M. 1992, Astrophys. J., 390, 536 Garcia-Segura, G., &; Mac Low, M.-M. 1993, in A.S.P. Conf. Proc, Massive Stars: Their lives in the Interstellar Medium, eds. J.P. Cassinelli & E.B. Churchwell (San Francisco: ASP), 534 Herrero, A., Kudritzki, R.P., Vilchez, J.M., Kunze, D., Butler, K., fc Haser, S. 1992, Astron. Astrophys., 261, 209 Humphreys, R.M. 1991, IAU Symp.N°143, Wolf-Rayet Stars and Interrelations With Other Massive Stars in Galaxies, 485 MacDonald, J.,& Bailey, M.E. 1981, Mon. Not. R. astr. Soc, 107, 995 Maeder, A. 1990, Astron. Astrophys. Suppl., 84, 139 Ryu, D., & Vishniac, E.T. 1991, Astrophys. J., 368, .411 Smith, L.J., Pettini, M., Dyson, J.E., Hartquist, T.W. 1984, Mon. Not. R. astr. Soc, 211, 679 Stencel, R.E., Pesce, J.E. & Bauer, W.H. 1989, Aslron. J., 97. 1120 Stone, J.M.,& Norman, M.L. 1992, Astrophys. J. Suppl., 80, 753 Weaver, R., McCray, R., Castor, J., Shapiro, P. fc Moore, R. 1977, Astrophys. J., 218, 377 Willis, A.J. 1991, IAU Symp. N°143, Wolf-Rayet Stars and Interrelations With Other Massive Stars in Galaxies, 265

High Resolution Coronographic Imaging and Echelle Observations of Si 19: a new Luminous Blue Variable?1 Antonella Nota 2 ' 3 , Laurent Drissen2, Mark Clampin 2 , Claus Leitherer 2 , Anna Pasquali 4 , Carmelle Robert 2 , Francesco Paresce 2 ' 3 ' 5 , Massimo Robberto 5 . 1

Based on observations made at the European Southern Observatory, La Silla. Space Telescope Science Institute, 3100 San Martin Drive, Baltimore, MD. 3 Affiliated to ESA, Astrophysics Division, Space Science Department, of ESA. 4 Universita' di Firenze, Arcetri, Italy. Osservatorio di Torino, Pino Torinese, Italy. 2

1 Introduction The LMC star S119 is a member of the group of 0fpe/WN9 stars listed by Bohannan and Walborn (1989). The Ofpe/WN9 category, first defined by Walborn (1982), identifies peculiar supergiants whose spectra combine the typical Of characteristic emission lines of He II and N III with equally strong lower ionization emission features, such as those of He I and N II, and are believed to represent a transition phase in the evolution between massive 0 stars and WR stars. 2 High Resolution Echelle Observations We have observed Si 19 with the high resolution echelle spectrograph EMMI, coupled to the NTT, ESO La Silla, on September 18, 1991. The spectra cover the wavelength range 4100A- 7800A, with a spectral resolution of 0.089 A/pixel at 6563 A. The selected slit width was 1.5" x 5", with a plate scale on the detector of 0.345"/pixel. In the spectrum, previously undetected nebular lines of Ha, H/3, [Nil], [SII] appear strong and spatially extended, an indication that S119 is surrounded by a bright gaseous nebula. We detect clear splitting of all the observed nebular lines. In Figure 1 we show the radial velocity map obtained from the [Nil] 6583 A line profile. During the observation the slit was oriented EW, and the star was not centered in the aperture, so that only the eastern portion of the nebula lies completely within the slit, while the western region is marginally covered (~ 2"). The radial velocities are calculated at several positions along the nebula, spaced by 0.345" in direction EW and are referred to the LSR. The radial velocity map is consistent with the nebula being a hollow shell, expanding at a velocity of ~ 25 km/sec. 89

90

A. Nota et ai: SI 19: a neiv Luminous Blue. Variable? 180

T

170 160 -•

-5J 150 E 140 >. 130 -D

O

120 -

>

110 --

• •

D n

100 90 --

80 -5

-4

-3

-2

-1

Position from star in arcsec (East-West) Fig. 1. Radial velocity map of the SI 19 nebula obtained from the [Nil] 6583 A line, where in ordinate we report the nebular velocities, referred to the LSR and expressed in km/sec, and in abscissa the nebular positions, measured along the East (left) - West (right) direction. The zero point in the abscissa is the location of the star.

2.1 The Systemic Velocity Radial velocity measurements of the stellar lines of Si 19 indicate a systemic velocity VLSR between +100 and +140 km/sec. This is in excellent agreement with the value obtained for the surrounding ring nebula. These velocities are totally at odds with systemic velocities of LMC member stars, which normally range from +240 to +300 km/sec. The rectified spectrum of S119 in the region of the interstellar Nal 5890,96 doublet is shown in Figure 2. The Galactic absorption features are clearly visible (with VLSR = +4 km/s), as well as an intermediate velocity component (VLSR - +110 km/s), probably originating in the Galactic Halo. However, no components are detected (ew < 8 mA) at velocities corresponding to those expected for the LMC (between +240 and +300 km/s). For comparison purposes, Figure 2 shows the same spectral region in the spectra of two nearby LMC Ofpe/WN9 supergiants (BE 381 and BE 294), taken with an identical instrument configuration, where the LMC interstellar components are prominent (especially in BE 381). The peculiar systemic velocity and the absence of LMC interstellar features suggest that, unless SI 19 is not a member of the LMC (a very doubtful hypothesis in view of its mag-

A. Nota et al.: SU9: a new Luminous Blue Variable?

91

X

N •—* t—t cd

fi m 5890

5892

5894

5896

5898

5900

5902

Wavelength (Angstroms) Fig. 2. The rectified spectrum of S119 in the region of the interstellar Nal 5890,96 doublet is here shown. Notice the absence of absorption components at the LMC expected velocity (between +240 and + 300km/s). Compare the same spectral region in the spectra of two nearby LMC Ofpe/WN9 supergiants (BE 381 and BE 294), where the LMC interstellar component is prominent. nitude and spectrum, characteristic of an Of supergiant), it has probably been ejected from its birthplace and is now located outside the main body of the LMC. This ejection could have occurred as a result of a supernova explosion in a close binary or, more likely, after interactions with other stars in the early phases of a cluster formation. Assuming that SI 19 was ejected from an LMC cluster with an average velocity of 150 km/s towards the Sun shortly after its formation, it could have travelled a distance of 600-800 pc during the 4-5 Myrs required for a 40 M 0 star to reach the Of/WN stage. This distance represents a non-negligible fraction of the depth of the LMC, and is probably large enough to put the star in front of most LMC neutral gas clouds.

3 Coronographic Imaging Si 19 was observed with STScI NTT Coronograph mounted on the New Technology Telescope, ESO, La Silla. Images were obtained on February 12, 1993, in excellent seeing conditions (FWHM ~ 0.75") in the light of [Nil] (Ae/y = 6584.6 A, AA = 23.8 A) and Ha. With this technique we resolve,

92

A. Nota et al.: S119: a new Luminous Blue Variable?

Fig. 3. [Nil] image of the S119 nebula taken with the STScI NTT Coronograph. The nebula is 7" x 9" in size. The central star is not occulted in this image, still the contrast in the circumstellar region is enhanced by pupil apodization. North is up and East to the left.

for the first time, the circumstellar nebula which is shown in Figure 3. The nebula has an extension of 7" X 9" which translates into 1.9 x 2.1 parsecs at the LMC distance of 51.2 kpc (Panagia et al. 1992). The morphology of the nebula is clearly axisymmetric, with a very bright lobe extending towards the NE up to ~ 4" from the star. Towards the SW, more complex structures can be discerned. We find from the spectral data that the nebula is expanding with a velocity of ~ 25 km/sec; the linear size then implies a dynamical timescale of ~ 5 x 104 yrs for the nebula. The integrated, dereddened [Nil] 6584 A flux is derived from the images to be 1.7xlO~ 12 ergs cm" 2 s" 1 , in the assumption of E(g_y) — 0.16 (Bohannan and Walborn, 1989). We can calculate the ionized mass of the nebula from the integrated Ha emission luminosity. From the spectra, we obtain a ratio ricvnef,u;ar/[NII] ~ 1. We adopt an average electron density for S119 of 800 cm" 3 , obtained from the [SII] 6716/6731 A line ratio. Assuming a temperature T e ~ 7500 K, the nebular mass of S119 is ~ 1.7 M 0 .

A. Nota et al.: S119: a new Luminous Blue Variable? Table 1. Comparison Star

category

of nebular

AG Car R 127 S119

Gal. LBV LMC LBV LMCOf

properties

size (PC)

1.1 X 1.0 1 . 9 x 2 .2 1. 9 x 2 .1

93

(km/s) -70 -281 -30

t (yr) - 104 ~ 4 x 104 ~ 5 x 104

P (cm" 3 )

M (M 0 )

500 1000 800

-4.2 -3.1 - 1.7

1: Walborn, 1982.

4 Is S119 a new Luminous Blue Variable? The similarity with two well known Luminous Blue Variables, the galactic AG Carinae (Nota et al. 1992, Robberto et al. 1993) and R127 (Clampin et al. 1993) in the LMC, is remarkable. They both exhibit axisymmetric gaseous shells, with comparable properties in terms of size, expansion velocity (v), density (p) and mass (M) as we show in Table 1. Their nebulae are most probably the relic of a massive outburst, which occurred some 104 yrs ago. From a morphological point of view, their similarity would point to a common mechanism for the outburst (Figure 4). This finding would support the assumption that a close relationship exists between Ofpe/WN9 stars and Luminous Blue Variables. After all, one of the original members of the Ofpe/WN9 class, R.127 (Stahl et al. 1983, Walborn, 1984) became a Luminous Blue Variable (LBV) following an increase in brightness in 1982, as its spectrum evolved from an O-type to a B-type star and then to an A-type star (Wolf 1992). Similar spectral evolution had been noted for other LBVs which had been observed during their brightening phases [e.g., AG Carinae (Caputo & Viotti 1970), S Doradus (Leitherer et al. 1985; Wolf 1989), and R71 (Wolf cl. al. 1981)]. Throughout these brightness and spectral variations, the bolometric luminosity of R127 remained constant, a very well know characteristic of LBVs. The behaviour of R127 provides indications that there is a close relationship between the Ofpe/WN9 stars and the LBVs. In addition, Of features have been observed in the LBV AG Carinae during a light minimum phase, with the possible implication that the Ofpe/WN9 stars are LBVs in a quiescent state (Bohannan and Walborn, 1989). The similarity of the nebulae displayed by two classes of objects is striking, and suggests the possibility that a common mechanism has generated the outburst. More observational coverage of the currently known Ofpe/WN9 stars is required to investigate a) if light and spectral variations are present b) the structure of the circumstellar environ-

94

A. Nota et ai: S119: a new Luminous Blue Variable?

Fig. 4. Circumstellar nebulae around the two LBVs AG Carinae (galactic - left) and R127 in the LMC (right). Notice the similar morphology of the two axisymmetric shells. The scale is different, in the two images: the AG Car nebula is 32" x 36 " in size, while the R127 nebula is 8" x 9". North is up and East to the left.

ment, in order to establish on a statistically significant basis the relationship between 0fpe/WN9 stars and LBVs. References Bohannan, B., fc Walborn, N. 1989, P.A.S.P. 101, 520. Caputo, F., & Viotti, R. 1970, A. &A.,7, 266. Clampin, M., Nota, A., Golimowski, D., Leitherer, C , fc Durrance, S. 1993, Ap. J. Letters, 410, L35. Leitherer, C , Appenzellar, I., Klare, G., Lamers, H.J.G.L.M., Stahl, O., fe Waters, L.B.F.M. 1985, A. & A., 153, 168. Nota, A., Leitherer, C , Clampin, M., Greenfield, P. fc Golimowski, D. 1992, Ap.J., 398, 61. Panagia, N., Gilmozzi, R., Macchetto, F., Adorf, H.-M., & Kirshner, R.P. 1992, Ap.J. Letters, 380, L23. Robberto, M., Ferrari, A., Nota, A., & Paresce, F. 1993, A. & A., 269, 330. Stahl, O., Wolf, B., Klare, G., Cassatella, A., Krautter, J., Persi, P., & Ferrari-Toniolo, M. 1983, A. & A., 127, 49. Walborn, N. 1982, Ap.J., 256, 452. Walborn, N. 1984, in IAU Symp. 108, Structure and Evolution of the Magellanic Clouds, ed. S. van den Bergh and K.S. de Boer, p.239. Wolf, B. 1989, A. & A. Suppl., 217, 87. Wolf, B., Appenzeller, I., & Stahl, O. 1981, A. & A., 103, 94. Wolf, B., 1992, in Nonisotropic and variable outflows from stars , ed. L. Drissen, C. Leitherer and A. Nota, (A.S.P. Conf. series Vol. 22), p. 327.

New HST images of Eta Carinae and its surrounding nebulosity Dennis Ebbets 1 , Harry Garner 1 , Rick White 2 , Kris Davidson 3 , Eliot Malumuth 4 , and Nolan Walborn 2 1

Ball Aerospace and Communications Group Space Telescope Science Institute 3 University of Minnesota 4 Computer Sciences Corporation 2

1 Introduction Eta Carinae is a hot, massive, very luminous star which has erupted with episodes of greatly enhanced mass loss several times during the past few centuries. It is surrounded by an expanding shell of material, commonly known as the Homunculus, which was ejected during an outburst between 1830 - 1860. Fainter condensations are visible at greater distances. Their composition and radial proper motions suggest clumps of ejecta, some of which are contemporaries of the Homunculus, while others may have been expelled during earlier outbursts. Eta Carinae represents a dramatic although possibly brief phase of the interaction between a massive star in an advanced stage of evolution and its environment. The spatial and temporal characteristics of the nebulosity make it an attractive target for observations with the Hubble Space Telescope. The smallest structures currently known have sizes between 1/10 to 1/4 arc second, beyond the reach of ground based instruments, but well matched to the resolution of HST. Proper motions are typically 0.05 to 0.1 arc seconds per year. Displacements of several pixels per year will be seen with HST. Over the planned 15 year HST mission the kinematics of the expanding debris should be clearly revealed.

2 Observations and Data Processing Our observations were made on February 25, 1993 using the first generation Planetary Camera. Filter F336W isolated a violet spectral region which is free of bright emission lines in Eta Carinae. We avoided saturation problems caused by the large contrasts between the central regions and the nebulosity seen in nebular lines. We acquired one short and three longer exposures. The telescope was moved slightly between the subexposures to translate the images by approximately eight pixels on the camera. This sequence 95

96

D. Ebbets et ai: HST images of Eta Carinae

; ' • • - '

*

K*~,.

•

*

1

•

Fig. 1. A Planetary Camera image of the Ilomiinculus

of exposures was repeated with a nearby star to provide a. point spread function. As a reference for future archival users, the root names of our observations are W13M0101 - W13M0108 for Eta Carinae, and W13M0201 - W13M0208 for the PSF star. The data products provided by the STScI were combined to produce composite images of Eta Carinae and the PSF star. The individual images were shifted into registration and coadded, with cosmic ray events identified and withheld from the sums. The few saturated pixels near the core were "repaired" by scaling the data values in the short exposure and inserting them into the composite. An "adaptive histogram equalization" algorithm was applied to improve the visibility of low contrast features, and a "damped Lucy" algorithm was used to deconvolve the point spread function while controlling noise in regions of low signal (White, 1993). The processed images were digitally rotated to the correct orientation (north up, east left), and displayed with a linear or square root grey scale, inverted to produce a "negative" image.

D. Ebbets et al: HST images of Eta Carinac

97

3 Morphology of the Homunculus Figure 1 shows the Homunculus displayed using the AHE algorithm. Note that because there has been no deconvolution performed on this image the regions immediately surrounding the bright central star are dominated by the psf. The scale bar is 10 arc seconds long. As Hester et al. (1991) showed with the first HST observations of Eta. Carinae, the Homunculus has a bipolar appearance, rather than an "ovoid" or "ellipsoidal" shape suggested by ground-based pictures. Two nearly circular lobes extend to the southeast and northwest, each of which is eight arc seconds in diameter. They are tangent to each other at the central star. The centers of the lobes and the central star lie along a line. There is nothing obviously at the "centers" of either lobe. The southeast component is more conspicuous, and is presumably closer to us and less obscured. Both lobes are resolved into a large number of small knots, typically 1/4 arc second wide and up to one arc second long. Many of these knots can be associated with features identified photometrically by Burgarella and Paresce (1991). Their shapes and orientations give the appearance of clumps along the walls of a hollow spherical cavity (cf. Meaburn et al. 1993a). If the material which is visible today was ejected around 1830 the proper motions of the outermost knots must be approximately 0.05 arc seconds per year. HST observations over a decade or so will trace the trajectories of many of the resolved knots, and will help discriminate between possible physical models.

4 The Central Region Figure 2 shows the central 5 arc seconds of our image processed with the Lucy algorithm. It is dominated by a single bright source, but shows dozens of discrete bright spots. Most of these objects have sizes around 1/4 to 1/2 arc second, and show elongated or irregular shapes. They appear to be knots of diffuse matter rather than stars. A group of knots on the northwest side form a linear structure with a position angle near 316 degrees. They lie along the major axis of the Homunculus, and appear to project about four arc seconds into the NW lobe. A cluster of knots to the northeast form an arc that defines the northern rim of the SE lobe. There is a conspicuous dark lane on the west side of the central star which appears to be obscuring brighter regions behind it. It arcs from the north to the southwest, and has a small extension which gives the appearance of the Greek "lambda".

98

D. Ebbets et al.: HST images of FA a Carinae

Fig. 2. The region surrounding the bright central star

Fig. 3. The core of Eta Carinae in violet light

D. Ebbets et al.: HST images of El a Carinae

99

5 The Core Glimpses of the core were obtained using ground based speckle imaging by Weigelt and Ebersberger (1986) and by Hofmann and Weigelt (1988). Figure 3 shows our best image of the inner one arc second enlarged from the Lucy restored violet data. In addition to the dominant bright star, Eta Car A, there are numerous fainter features with diameters less than 1/4 arc second. The locations of speckle components B, C and D are indicated. Components C and D correspond to the brightest features in our image (after A). Their positions relative to A agree to within 0.05 arc seconds between 1985 and 1993. We see no evidence that these two components have proper motions comparable to those measured in the Homunculus and outer condensations. Speckle component B is not visible in our data. Either it is much fainter in our violet continuum than in the broad band red (which contained bright nebular emission lines), or it has moved or dissipated in the intervening eight years. One of the other components visible here is particularly interesting. The knot at a position angle 28 degrees (northeast) and 0.3 arc seconds from component A is aligned with the more distant, features which comprise the NN and NS condensations, and the intriguing jet like structure which they define (Hester et al. 1991, Meaburn et al. 1993b). The NN and NS condensations have proper motions of 0.1 arc seconds per year. If this new knot is kinematically related, and not just a coincidence in position angle, its motion away from the central star will be measurable with HST over periods of months, and questions about uniform vs decelerated motion, and its interaction with the ambient medium, should be fruitfully tackled with the HST. Acknowledgements This work is based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by AURA, Inc. under NASA contract NAS5-26555. References Burgarella, D. & Paresce, F. (1991). Astron. Astrophijs., 241. 5<).r>. Hester, J. et al. (1991). Astron. Journal, 102, 654. Hofmann, K. & Weigelt, G. (1988). Astron. Astrophys.. 203. L21. Meaburn. J., Walsh, J. & Wolstencroft, R. (1993a). Astron. Aslrophys., 2C8, 283. Meaburn. J. et al. (1993b). Astron. Astrophys., (preprint) Weigelt, G. k Ebersberger, J. (1986). Astron. Astrophys., 103, L5. White, R. (1993), Newsletter of the STScI Image Restoration Project, 1, 11.

Observations of Circumstellar Media Around Supernovae Bruno Leibundgutf Astronomy Department University of California Berkeley, CA 94720 USA

Abstract Some supernovae are visible for several years past explosion. The main energy source for this sustained emission conies from the supernova shock interacting with the remnant of the stellar wind of the progenitor star. We review the available evidence for this picture and exclude other power sources on the basis of the radiated energies. We also discuss a group of supernovae which display narrow emission lines with high fluxes in their spectra and very slowly declining optical light curves. These observations can most readily be explained as being clue to interaction with a very dense medium close to the supernova.

1 Introduction A variety of supernova interactions with circumstellar material (CSM) has been observed to date. The best, and most direct, example is the ring of emission around SN 1987A (Jakobsen et a.l. 1991). This material has been ionized by the UV and soft X-ray flash of the shock breakout at the surface of the supernova (Fransson et al. 1989, Lundqvist & Fransson 1989). The density enhancement in the ring is caused by the interaction of the fast blue supergiant wind colliding with the slow red supergiant wind of a. previous epoch (Blondin & Lundqvist 1993). In the case of SN 1993.1, the early detection of radio and X-ray emission, in combination with narrow emission lines in the UV and optical, are indicative of interaction with the CSM. Blue optical continua, X-ray detection at early phases, as well as the UV emission have been proposed as characteristics of a shock in the CSM around SN 1979C (Fransson 1984). Infrared light echos observed around SN 1979C and SN 1980K are indicators of cold CSM (Dwek 1983). The most frequently observed signatures of CSM around supernovae, however, are radio emission (Chevalier 1984, 1990) and optical detection several years past explosion. The structure of the stellar interior is exposed when the explosion becomes transparent and the circumstellar environment is excited by the supernova t Present Address: European Southern Observatory. Karl-Sdiwar/.schilcl-Stiasse 2, D-85748 Garching bei Miinchen, Germany

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B. Leibundgut: Supernovae and their circumstellar environment

101

Table 1. Supernovae interacting with circumstellar media SN 1957D 1961V 1970G 1978K 1979C 1980K 1981K 1983N 1984L 1986J 1987A 1988Z 1989R 1990B 1993J

Type

radio

Ref.

?

V V V V V V V V (V)

1 2 3-7

II II ?

II II Ib Ib II II II II II II

v/

V

8 9 10 6 6,11 6 12-15 16,17 18 — 19 20,21

optical Ref. (>3 years)

V v/

V V V V o —

V V V 0

—

22-24 25,2,26 27 8,28 29 30-32 — 22 — 31,32 33,34 35,36 37 38 still young

X-rays

Ref.

_ — —

_ — —

V

8 —

— —

(V) — — v/

39 — — — 40

41,42 — — —

43,44

Refei'ences: 1 Cowan k Branch 1985; 2 Cowan et al. 1988; 3 Gottesman et al. 1972; 4 Allen et al. 1976; 5 Brown k Marsclier 1978; 6 Weiler et al. 1986; 7 Cowan et al. 1991; 8 Ryder et al. 1992; 9 Weiler et al. 1991; 10 Weiler et al. 1992; 11 Panagia et al. 1986; 12 Rupen et al. 1987; 13 Weiler et al. 1990; 14 Sukumar k Allen 1989; 15 Bartel et. al. 1991; 16 Turtle et al. 1987; 17 Staveley-Smith et al. 1992; 18 Sramek et al. 1990; 19 Van Dyk et al. 1993a; 20 Pooley k Green 1993; 21 Van Dyk et al. 1993b; 22 Long et al. 1989; 23 Long et al. 1992; 24 Turatto et al. 1989; 25 Fesen 1985; 26 Goodrich et al. 1993; 27 Fesen 1993; 28 Dopita k Ryder 1990; 29 Fesen k Matonick 1993; 30 Fesen k Becker 1990; 31 Leibundgut et al. 1991; 32 Uomoto 1991; 33 Bouchet et al. 1991; 34 Suntzeff et al. 1991; 35 Stathakis k Sadler 1991; 36 Turatto et al. 1993; 37 this contribution; 38 A. C. Porter, private communication; 39 Canizares et al. 1982; 40 Bregman k Pildis 1992; 41 Dotani et al. 1987; 42 Sunyaev et al. 1987 43 Zimmermann et al. 1993; 44 Tanaka et al. 1993

shock. Investigation of the stellar properties, however, has to await fading of the radiative display from the explosion. In a few cases the CSM is dense enough to dominate the emission from the explosion. The list of supernovae (SNe) still optically observable several years after outburst includes now at least ten objects (Table \). With one exception, SN 1989R which has not been observed, all these objects have been detected at radio wavelengths. Conversely, all SNe II with a radio detection have been recovered optically. All SNe Ib/c with radio observations faded in the optical a few (<3) years past their discovery. Two supernovae observed in the radio

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B. Leibundgut: Supernovae and their circumstellar

environment

have no classification; SN 1957D has been recovered in the optical several decades after maximum, while SN 1981K has no optical detection so far. The radio light curve of SN 1987A is rather peculiar and is not necessarily due to strong interaction with the CSM (Staveley-Smith et al. 1992). Five supernovae have also been detected in X-rays at different epochs in their evolutions. While SN 1980K and SN 1993J exhibited X-rays shortly after outbreak, SN 1978K and SN 1986J had observable X-ray fluxes several years past outburst. The X-ray flux of SN 1987A is probably due to Comptonization of 7-rays and not necessarily related to CSM. The optical evolution of SNe II can be separated into several phases (e.g. Leibundgut 1994). Interaction with CSM can influence the general evolution by different amounts for individual supernovae. The contribution from CSM interaction can dominate the optical emission from the start (e.g. SN 1988Z; Turatto et al. 1993), influence the evolution near peak (SN 1993J; Gumming et al., this conference), or drive the emission several years after explosion (Chugai 1990, 1992, Chevalier & Fransson 1994). We will describe the evidence for shock interaction of a supernova with its CSM and present arguments that exclude other power sources. First we will review the available optical observations for supernovae older than a decade (§2). In section 3, we present the data on SNe II with strong and persistent narrow-line emission. The conclusions (§-•!) will try to justify the above distinction among SNe II.

2 Old Supernovae The investigation of supernovae at very late phases has still a rather short history. The first supernova for which an optical identification was made at an age >10 years is SN 1961V (Fesen 1985, Cowan et al. 1988), but its classification as a supernova has been questioned (Goodrich et al. 1989, Bower et al. 1993). Three objects were discovered by a dedicated search program of Robert Fesen, who recovered SN 1980K (Fesen & Becker 1990), SN 1979C (Fesen & Matonick 1993), and SN 1970G (Fesen 1993). At the same time a search for supernova remnants in nearby galaxies disclosed SN 1957D more than two decades after explosion (Long et al. 1989). SN 1978K and SN 1981K have very little optical information and will not be discussed here.

2.1 SN 1957D in M83 This supernova in the southern sky has no spectral classification. Nonthermal radio emission was found by Cowan k. Branch (1985). The same

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radio map also shows non-thermal sources near the sites of SN 1950B and SN 1983N, but neither has been recovered in the optical (Long et al. 1989). At optical wavelengths SN 1957D was first detected in the light of the [0 III] (AA 4959, 5007 A) doublet lines in 1987. The narrow-band imaging was verified by spectra obtained a year later (Long et al. 1989, Turatto et al. 1989). Strong emission in [0 III] but only weak lines of Ha and [O I] (AA 6300, 6364 A) were found. A remarkable drop in optical luminosity occurred between 1987 and 1991, when the supernova faded by a factor of ~5 (Long et al. 1992). This fading has been attributed to the supernova shock reaching the edge of the wind zone of the progenitor star. Another interpretation is that in 1987 an excited state was observed when the shock ran into a condensation. In this case the emission would vary on time scales of years. Careful monitoring of SN 1957D is expected to provide further clues as to the nature of this object.

2.2 SN 1910G in M101 The location of SN 1970G only 0.5 arcseconds away from a stellar-like source, which is probably part of the giant HTI complex NGC 5455 in M101 (NGC 5457), makes observations of this object exceedingly difficult (Fesen 1993). SN 1970G has been clearly recovered in [0 I] images and a spectrum has been secured. Broad emissions of [0 I] and Ha are detected, with narrow lines superposed. The latter are most likely due to the nearby HII region. The spectral identification of SN 1970G has profited greatly from comparison with SN 1980K, which shows similar line widths (Fesen 1993). While [O I] and Ha lines are observed, no [0 III] emission is detected in SN 1970G, a remarkable distinction from all other supernovae at similar ages. Fesen (1993) points out that SN 1970G exhibited a normal light curve and spectral evolution around maximum (Barbon et al. 1973, Kirsliner et al. 1973) and does not constitute a peculiar case, which might be argued for SN 1979C and SN 1980K (see below).

2.3 SN 1979C in M100 This supernova in the Virgo cluster of galaxies is the most distant in the sample of decade-old SNe II. It was exceptionally luminous at maximum (de Vaucouleurs et al. 1981, Branch et al. 1981) and its light curve did not display the characteristic plateau observed in many SNe II. The spectral evolution around maximum was also peculiar, exhibiting narrow Ha emission and weak absorption in the P Cygni profile of Ha (Branch et al. 1981).

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Interaction with CSM was inferred from the UV and blue spectral shape (Fransson 1984) and the strong radio emission, which turned on less than 13 months after explosion (Weiler et al. 1986). This supernova is relatively luminous even at these very late phases (Fesen 1993). The optical emission of SN 1979C is again mostly in [0 I] and Ha lines but also with a sizable [0 III] emission (Fesen & Matonick 1993). In addition, [0 II] or [Ca II] near A 7300 A is observed. The spectrum appears unchanged over the course of one year (Fesen & Matonick 1993). Narrow lines are probably due to a nearby HII region, which is situated less than 2 arcseconds from SN 1979C.

2.4 SN 1980K in NGC 6946 The most extensive data set is currently available for SN 1980K. It was recovered optically in 1987 (Fesen & Becker 1990) and has been monitored since (Leibundgut et al. 1991, 1993, Uomoto 1991, Fesen 1993). Around maximum this supernova is one of the best studied SN II. Its light curves declined steadily for the first 13 months (Barbon et al. 1982, Buta 1982). The spectral evolution was similar to SN 1979C (Barbon et al. 1982). It developed a nebular spectrum after about 8 months with strong Ha, [0 I], and [0 III] emissions (Uomoto & Kirshner 1986). Although the decline rate of Ha matched the half-life of 56 Co decay, Chugai (1990) showed that the emission had to be powered by an additional energy source. Note that this differs from the bolometric light curve (Schmidt et al. 1994), where the decline indeed reflects the characteristic lifetime of the radioactive source. More than a decade after maximum SN 1980K is still emitting in the same lines. The decline, however, has stopped completely and all lines are emitting at a constant flux level (Fig. 1). Typical line widths are ~5000 km s" 1 (Fesen & Becker 1990, Leibundgut. et al. 1991). Broad-band photometry established SN 1980K at V=22.8±0.2 mag. and R=21.9±0.1 mag. (Leibundgut et al. 1993). The total luminosity in the BVRI wavelength range amounts to 1.2-10"1'1 erg cm" 2 s" 1 . Given the distance to NGC 6946 the total emitted flux is then 8-1037 • (yyrs—) 2 erg s" 1 . Such high fluxes exclude long-lived radio isotopes as the power source since 10~2 M 0 of ''''Ti would be needed, in contradiction with nucleosynthesis yields in SNe II (Woosley et al. 1989). Since the line fluxes did not change subsequent to the recovery of SN 1980K, it is safe to integrate over this time span. The emitted energy is then ~ 10 46 • ( 7 5 ^ ) 2 erg and, if the flux had not changed since 1982 (the last observation of Uomoto & Kirshner), this number is doubled. This rules

B. Leibundgut: Supernovae and their circumstellar

10"

* io

12 -12

1

1 t

i

i

-13

-14

1

L-

10" 1 3 ^ CO

'

E

environment

1

i

1 ' i SN 1980K ^ —=

*

—=

2> 10 -15

^

10

•

i-

10-16 4000

105

A

O i

,

, 1

1

1

,

i

, -

6000 8000 JD (24440000+)

Fig. 1. Ha (filled symbols) and [O I] (open symbols) line fluxes for SN 1980K. The data are taking from the following references: Uomoto & Kirshner 1986 (circles), Fesen & Becker 1990 (triangles), Uomoto 1991 (hexagons), Leibundgut et al. 1991 (squares), and this contribution (diamonds).

out a light echo from the shock breakout as the power source, since at least 10% of the UV flash would have had to be converted into optical emission and emitted over a period of >10 years. Because the flash ionizes the volume just once, the number of Ha photons (3 • 10 5 ') would have required >3 MQ of hydrogen, assuming Case B recombination. The flux in SN 1980K is remarkably close to what, is expected from a powerful pulsar (like the Crab). Chevalier & Fransson (1992) have investigated the influences of powerful pulsars inside old SNe and young supernova remnants. They predict line velocities of a few hundred km s" 1 , which is not observed in SN 1980K. The observed line widths of ~5000 km s" 1 exclude this power source for SN 1980K. Shock interaction with CSM taps the kinetic energy pool (~10 51 erg) by converting a fraction of it into radiation. The forward shock ionizes the CSM and X-rays from the reverse shock produce emission from the ejecta (Chevalier & Fransson 1994). The observed energies and the line widths strongly favor circumstellar interaction. While the presence of the [O III] lines is probably due to ejecta heating, the [0 I] and Ha emissions more likely arise from the CSM (Chevalier & Fransson 1994).

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environment

Summary

While only a few supernovae at very late epochs have been detected and observations of them are still scarce, a trend seems to be emerging. All objects show emission lines of the same elements, namely hydrogen and oxygen with a possible contribution from calcium to the 7300 A feature. This fits nicely with the CSM model presented by Chevalier & Fransson (1994). The ionizations differ among the individual SNe: SN 1957D shows mainly [0 III] emission, while this is not detected in SN 1970G. The constant flux in SN 1980K, however, is not predicted by the models.

3 Strong Interactors In recent years a few SNe II have been observed displaying some distinct properties. A first example was presented by Dopita et al. (1984) with observations of SN 1984E at maximum. This object had narrow lines of hydrogen and helium on top of strong P Cygni profiles. The narrow emissions were attributed to an extended, optically thick, outflow. For the derived density a very strong wind phase had to bo invoked (Dopita et al. 1984). Schlegel (1990) proposed a subclass for SNe II consisting of events with narrow emission lines. His list included a few objects which are clearly influenced by nearby HII regions and the class as a whole has no distinct characteristics apart from the narrow lines. We will discuss here three objects with narrow emission lines and particularly slow decline rates, which are not affected by contaminating light.

3.1 SN 19863 in NGC 891 The longest-observed supernova of this group, SN 1986J, was first discovered in the radio in 1986, but is found on radio maps and optical images as far back as 1984 (Rupen et al. 1987). Estimates of the explosion date have large uncertainties, partly due to the peculiar nature of the supernova, but late 1982 seems generally accepted (Rupen et al. 1987, Weiler et al. 1990). The optical spectrum of SN 1986.1 is remarkable for its narrow, unresolved (Av<700 km s" 1 ) emission lines of II, He, N. and Fe, and broader 1 (AVR*2000 km s" ) lines of three ionization stages of oxygen (Leibundgut et al. 1991). The [0 I] doublet lines did not display the usual optically thin ratio, but rather were of about equal strength. A detailed non-LTE investigation indicated high density material (6 • 10s < UQO < 2 • 109 cm" 3 ) at low temperatures (3000 < T < 4000 K). The line strengths of the narrow components decreased steadily, but at a rate much slower than if they had

B. Leibundgui: Supernovae and their circumsiellar environment

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been powered by 56 Co decay. The absolute fluxes of the oxygen lines, however, have remained roughly constant (Leibundgut et al. 1991). If the [N II] (A5755 A) line indeed comes from the CSM, as suggested by the narrowness of the line, this implies that the wind is nitrogen-enriched as in SN 1987A (Fransson et al. 1989). SN 1986J has also been observed at X-ray wavelengths (Bregman & Pildis 1992). This emission could be either from the reverse shock or from dense clumps in the circumstellar medium excited by the outer shock (Chugai 1993).

3.2 SN 1988Z in MCG

03-28-022

The photometric and spectroscopic appearance of this supernova was also rather peculiar. Narrow lines dominated the spectrum, but broad and intermediate components of Ha could be resolved at early phases (Stathakis & Sadler 1991, Turatto et al. 1993, Filippenko 1991). The expansion deduced from the broad component was ~40000 km s" 1 and matches velocities in other SNe II. No P Cygni absorption was recorded for SN 1988Z in the photospheric phase. In addition to the narrow lines of II, lie, N, and Fe, broad emission from Ca was detected (Turatto et al. 1993). A broad bump around 5300 A observed half a. year after discovery was attributed to the blending of many Fe II lines (Chugai, this conference). The Ha luminosity evolution is also very slow with a peak about 400 days after discovery at around 1041 erg s~1 and a. decline to ~3-10 40 erg s" 1 within 1200 days (Turatto et al. 1993). The shapes of the L(Hcv) curves, and the spectra at about 3-4 years after explosion (Leibundgut 1991) are rather similar for SN 1988Z and SN 1986J.

3.3 SN 1989R in UGC 2912 This is a new representative of SNe II with narrow lines and slowly decaying light curves. The first spectrum 2 days after discovery displays narrow emission lines (Kirshner et al. 1989) on a rather featureless continuum. There is only a hint of broad emission visible at Ha. Later observations show considerable reddening of the spectrum. The TR. triplet of Ca II is possibly detected 2 years past discovery, but should be confirmed by a more careful analysis. The R light curve of SN 1989R has a. break at around 250 days past discovery, when the slope changes from 0.025 ± 0.001 to (1.9 ± 0.2) • 10~ 3 mag/day. The SN was still detectable 3 years after discovery.

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Summary

We have presented three cases of SNe with narrow emission lines of H, He, and N in their spectra and small decline rates in their R light curves. A detailed comparison of the spectrum of SN 1986J four years past explosion with that of SN 1988Z three years past discovery reveals a great similarity. The same set of lines is present in both SNe. They were also prolific radio emitters and are easily distinguished from the other radio supernovae (Van Dyk, this conference). SN 1989R is a new and fainter example of this group. Scaled to the distance of SN 1988Z it is ~2.3 times (0.9 mag) fainter at discovery in R. Note that none of the SNe has been discovered on the rise. The light curve of SN 1988Z is decaying very slowly (~0.003 mag/day) and the SN has faded by <4 magnitudes in 3 years. SN 1986.1 and SN 1989R declined in a comparably slow manner. All properties indicate an interaction of the supernova with dense CSM from the very onset of the explosion.

4 Conclusions Circumstellar material around supernovae is evidenced in several ways. The UV flash from the explosion can excite this matter and is observable as spatially resolved emission around SN 1987A, as narrow emission lines superposed on the supernova UV and optical spectrum (Fransson et al. 1989; Sonneborn, this conference, Gumming et al., this conference), or as an IR echo (Dwek 1983). Direct observations of X-rays and a blue continuum from down-scattered X-ray photons can indicate a shock in the CSM (Fransson 1984). The most obvious signature of interaction with the CSM so far is the radio emission (Van Dyk, this conference). All SNe with sustained optical radiation are radio supernovae. The supernovae discussed here have been separated into two groups. To clarify the main distinction, we summarize the fluxes and energies emitted by all objects in Table 2. The measured fluxes are all rather small and constant for most old (t>10 years) SNe, while large but decreasing for the younger ones. All narrow-line SNe have been observed in rather distant galaxies (H O =75 km s" 1 Mpc" 1 is assumed for SN 1988Z and SN 1989R; all other distances are either from cepheids or planetary nebulae luminosity functions). The luminosity in the Ha line is two to three orders of magnitude larger in the narrow-line SNe. Note that we did not correct for extinction. An absorption of Ay = 1.5 for SN 1986J (Leibundgut et al. 1991) increases the flux and luminosity of this object by a. factor of ~-1. Thus, the narrow-line

B. Leibundgut: Supernovae and their circumsiellar environment Table 2. Ha Luminosities SN

Galaxy

1957D 1970G 1979C 1980K 1986J 1988Z 1989R

M83

a

15

e

1

in 10~ erg s" cm~ in 1U

c d

M101 M100 NGC 6946 NGC 891 MCG+03-28-22 UGC 2912

109

of old supernovae

D (Mpc)

F(Ho) a

L(Ha)*

E(IIa) c

Period4 (years)

4 7 16 7.5 10 90

12-2e

20-4 e

1.8 2.5 1.7

>10 >80 >10

>le >10 >10 >10

4 22 11

100-10 100-10

2.3

75

20-1

>100 >1000 >100

>10 3 -10 2 >10 5 -10 4 >10 4 -10 3

10 5 2

2

erg s

in 1045 erg Period of integration for the total emitted energy Fluxes from the [0 III] lines

supernovae are much more luminous than regular old SNe. This indicates a much stronger interaction with a very dense medium as was inferred for SN 1986J and SN 1988Z (Filippenko 1991, Turatto ct al. 1993, Chugai & Danziger 1994), and is supported by the X-ray detection of SN 1986J (Bregman k Pildis 1992, Clvugai 1993). The detection of [N II] indicates nitrogen-enriched CSM, possibly from the dredge up in the convection near the surface of a rather massive star. The total energies emitted in Ha are remarkably uniform at >10 4 6 erg (over a decade time span) for the old SNe II. The other group has energies of about one to two orders of magnitude higher over much shorter periods. Other distinctions include the different radio luminosities, radio light curve shapes (Van Dyk, this conference) and decline rates of the optical light curves. A further distinction lies in the widths of the emission lines. The narrow lines arise from H, He, N, and Fe, while O appears in extended lines. For the old supernovae only broad 0 and Ikv lines are observed. Narrow-line supernovae display all signatures of strong interaction with their circumstellar environment. The optical emission is dominated by the CSM at all times, while for regular supernovae the CSM interaction becomes visible only a few years past explosion. Observations of CSM around supernovae have become available only recently and the sample is still small. Nevertheless, the shock interaction with the remnant of the stellar wind seems clearly confirmed in all cases. Although we are not yet able to derive detailed parameters of the CSM from

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these observations, they reinforce our picture of SNe II resulting from core collapse in massive stars that suffered appreciable mass loss during their lives. Eventually, we should be able to map such mass loss histories as the supernova shock moves through the CSM. The variations in SN 1957D might be a hint of clumping in the wind or pulsations of the progenitor. Likewise, flux variations might disclose the presence of companion stars modulating the mass loss as proposed from radio data for SN 1979C (Weiler et al. 1992).

Acknowledgements I would like to thank W. Vacca for helpful comments on the manuscript. Financial support by the Swiss National Science Foundation and through NSF Grant AST-9115174 is gratefully acknowledged.

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Radio Supernovae as Probes of Progenitor Winds Schuyler D. Van Dyk1-2, Kurt W. Weiler1, Nino Panagia 3 , and Richard A. Sramek4 1

Remote Sensing Division, Code 7215, Naval Research Laboratory, Washington, DC, 20375-5351, USA 2 Naval Research Laboratory/National Research Council Cooperative Research Associate Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA 4 National Radio Astronomy Observatory, P. O. Box O, Socorro, NM 87801, USA

Abstract Radio supernovae (RSNe) are an excellent, means of probing the circumstellar matter around, and therefore the winds from, supernova (SN) progenitor stars or stellar systems. The observed radio synchrotron emission is best described by the Chevalier model, which involves the generation of relativistic electrons and enhanced magnetic fields through the SN shock interacting with a relatively high-density circumstellar envelope, which is presumed to have been established through mass loss in the late stages of stellar evolution. From the Chevalier model, modified to include a mixed, internal, nonthermal emission/thermal absorption component, we can use the observed radio emission from these SNe to derive physical properties, including the ratio of the mass-loss rate to the stellar wind speed, which determines the circumstellar matter density. Assuming a value for the wind speed then allows us to determine a mass-loss rate for the star responsible for the circumstellar matter and to estimate its mass. For Type II RSNe, this mass loss is assumed to originate from the presupernova star itself, while for Type Ib/c RSNe, the stellar wind is assumed to be from the binary companion to the SN progenitor. Extreme examples of progenitor winds are found for unusual Type II RSNe, for which radio properties indicate that the matter around these SNe resulted from very high mass-loss rates in the late stages of the evolution of very massive stars. Additionally, we have observed deviations from the standard model, probably providing evidence for inhomogeneities in the circumstellar matter density and possibly indicating the presence of stellar pulsations or an interacting binary companion.

1 Introduction SN 1970G was the first instance of a SN with detectable radio emission (Gottesman et al. 1972), and SN 1979C was the first SN to be observed in detail over a lengthy time span (Weiler et al. 1986, 1991). Since the detection of radio emission from SN 1979C in April 1980, we have carried out an 112

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ongoing monitoring program using the Very Large Array (VLA). We have observed about 50 SNe for radio emission (see Weiler et al. 1989, Van Dyk et al. 1994), and of these, eight examples of RSNe have been detected and studied in detail: SN 1980K, SN 1981K, SN 1983N, SN 1984L, SN 1986J, SN 1988Z, SN 1990B, and SN 1993J. In addition, Ryder et al. (1993) have recently discovered radio emission from the unusual SN 1978K, which appears to have radio properties similar to SN 1986J and SN 1988Z. Another unusual RSN is the very luminous Mk 297A recently discovered by Yin & Heeschen (1991). We will not discuss here the SNe 1957D in M 101 (Cowan & Branch 1985) and 1987A in the LMC (Turtle et al. 1987; Staveley-Smith et al. 1992), which have also been detected in the radio, since these are, respectively, poorly studied and of a very unusual SN type. However, we do draw the reader's attention to an interpretation by Ball & Kirk (1993) of SN 1987A's current radio emission. We find that most RSNe share the following common properties: 1) nonthermal emission with high brightness temperature; 2) light curve "turn-on" at shorter wavelengths first and longer wavelengths later; 3) a rapid increase in flux density with time at each wavelength, with a power-law decline, with index /3, after maximum; and, 4) a decreasing spectral index between two wavelengths, as the longer wavelength emission goes from being optically thick to thin, with the spectral index, a, asymptotically approaching an optically thin, nonthermal, constant negative value (see Weiler et al. 1986). Over the years we have also noted several systematic trends in radio emission from SNe (see Weiler et al. 1986, 1989): Type la SNe are not radio emitters to the sensitivity limit of the VLA; Type Ib/c SNe (e. g., SN 1983N, 1984L, 1990B) have a steep spectral index, a rapid "turn-on" at 6 cm before optical maximum, a rapid decline after maximum, and homogeneity in all radio properties, particularly spectral luminosity: and, Type II SNe (e. g., SN 1979C, 1980K, 1981K) show more diversity, a flatter spectral index, a. slower "turn-on" at 6 cm after optical maximum (1 month to more than 1 year), and a slow decline after maximum . Extreme examples of Type II RSNe are SN 1986J (Weiler et al. 1990), SN 1988Z (Van Dyk et al. 1993b), and, possibly, SN 1978K (Ryder et al. 1993). In Figure 1 we illustrate the difference in behavior between the three groups of RSNe by showing the radio light curves for the Type Ic RSN 1990B (Van Dyk et al. 1993a), the "normal" Type II RSN 1979C (Weiler et al. 1991), and the peculiar Type II RSN 1988Z (Van Dyk et al. 1993b). The radio emission from SNe can best be described by the Chevalier (1981, 1982, 1984) "mini-shell" model, with the radiation interpreted as synchrotron emission from relativistic electrons/positrons and enhanced mag-

1 .5

:' I — -

ii I, 1.8

2

2.2

i

| in

2.4

MI

|

m

-

-—

rI

[//

2.6

1.7

5. D. Van Dyk tt al.: Radio supernovae & progenitor winds

r

114

I J haft

b)

ml

IMI

2.8

3

Mil I I I I L

3.2

3.4

3.6

log (Days Since Explosion)

log (Days Since Explosion)

.4

2.8

2.9

3

3.1

3.2

3.3

log (Days Since Explosion)

Fig. 1. The radio light curves for: a) SN 1990B, b) SN 1979C, c) SN 1988Z. Data at 20 cm are represented by open triangles; data at 6 cm, by open squares; data at 3.6 cm, by filled triangles; and, data at 2 cm, by filled squares.

netic fields external to the SN photosphere produced through interaction of the SN shock wave with high-density circumstellar matter arising from mass loss in a red supergiant stellar wind before explosion. It is through this interpretation that RSNe are an excellent means of probing the circumstellar matter, and therefore the winds from, massive presupernova stars or star systems.

2 Modelling the Radio Emission Using the Chevalier model for the radio emission from SNe, we can derive model radio light curves to fit the observed data. The formulation for this model is based on several assumptions. First, the external circumstellar medium is assumed to have a uniform density PcWcum = (M/iv)/4nr2 from

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a red supergiant wind of constant speed w and mass-loss rate M. Second, the SN shock radius evolves with time during the interaction with the circumstellar medium as R oc tm (Chevalier 1981, 1982). Third, the SN ejecta have a density profile of dejecta oc r~n, where m and n are related by m = (n — 3)/(n — 2). And, fourth, the nonthermal radio emission from circumstellar interaction can be formulated as 5 oc i/ a ^-(T+ 5 - 67n )/ 2 = vatP, where 7 = —2a + 1 is the index of the power law of the relativistic electron energy spectrum and a is the optically thin radio spectral index (Chevalier 1984). The progression in "turn-on" time of the radio light curves with wavelength, as well as the very steep initial rise of flux density, are due to the absorption of the synchrotron radiation by the intervening circumstellar medium. We assume that the medium is completely ionized, such that the absorption is purely thermal, free-free, with frequency dependence v~2A. For free-free absorption external to the emitting source, r a R~3, so that r oc i~ 3m = ts. In addition to the external absorbing matter, it has been found that for some RSNe for which the circumstellar medium is very dense, e. g., SN 1986J (Weiler et al. 1990), the addition of an internal, expanding, mixed nonthermal emitting/thermal absorbing medium is necessary. Assuming constant ejected mass, constant temperature, and homologous expansion, the electron density for this internal emitting/absorbing region is ne oc R~3 oc t~3m, and for absorption by internal matter, r' oc EM oc n2eR oc

r 5 m = ts>.

Combining these components, the generalized formulation for the model flux density, S, the external absorption, r, and the internal absorption, r', for a RSN is:

where -2.1

and -2.1

/.i

* \ S'

(The mixed absorber attenuation term (1 - C~T')T'~X is appropriate for a plane-parallel approximation.) The scaling parameters Ii\, A'2, and A'3 formally correspond to the flux density, external absorption, and internal absorption, respectively, at 5 GHz, one day after explosion. We solve for

116

5. D. Van Dyk et al.: Radio supernovae & progenitor winds Table 1. Model Parameters

for Radio

SN

a

A2

Type Ib/c RSNe: 1983N 4.4 x 10 3 -1.03 1984L 2.8 x 1O2 -1.01 1990B 2.0 x 102 -1.12

-1.59 -1.48 -1.27

5.3 x 102 6.9 x 102 1.5 x 10"

Type II RSNe: 1970G 1.5 x 105 1979C 1.5 x 103 1980K 7.4 x 101 1981K 1.9 x 101 1986J 6.7 x 105 1988Z 1.5 x 105

4 -1.58 8.7 x 10 -0.78 3.7 x 107 -0.66 3.4 x 105 5 -0.50 < 1 . 5 x 10 -1.18 3.0 x 105 4.0 x 1012 -1.50 5.8 x 104 1.0 x 10 12

A'i

a

-0.81 -0.74 -0.52 -0.74 -0.67 -0.80

Supernovae A:3

to 1983 Jun 29 1984 Aug 12 1989 Dec 15

1979 Apr 1980 Oct 1981 Jul 1982 Sep 1988 Jan

04 01 31 13 23

the "best" values for the fitting parameters a, ft, I\\, A'2, A'3, and to ( t h e date of explosion) from each RSN's dataset using a minimum reduced x2 procedure. To reduce the number of fitting parameters we assume from t h e Chevalier model t h a t 6 = a - (3 - 3 and 6' = 5tf/3 (see VVeiler et al. 1990). The model provides a good fit t o the d a t a for all RSNe, indicating t h a t t h e RSNe generally behave in a regular manner (although deviations from this model are observed for some RSNe; see §5). T h e fitting parameters for t h e well-studied RSNe are listed in Table 1. We give t h e calculated values for 6 and 6', as well as the indices m and n for the SN, in Table 2. T h e fitting parameters and calculated properties for SN 1993J are not well determined at this time (see Weiler et al., this volume).

3 Mass-Loss Rates from the Progenitor Star Systems We can use the fitting parameters to derive M / w , which is essentially a measure of the circumstellar density pcircnm, using the following expression appropriate for free-free absorption (Equation 16 from Weiler et al. 1986): M

- P ~ 3.02 x 10-6 r?^ Hz m"1-5 [——^

w/10 km sec"- 1

1.5

A

\10 4 kmsec" 1 / (4) 1.5 / f \1.5ra /

rp

\ 0.68

where for / = 1 day, we assume r 5 GHZ(1 da.v) = A'2. If we further assume

S. D. Van Dyk et ah: Radio supernovae & progenitor winds

117

Table 2. Calculated Properties for Radio Supernovae m

n

( = -
("»= S=§)

M (M o yr- 1 )

Type Ib/c RSNe: 1983N -2.44 1984L -2.53 1990B -2.85

0.81 0.84 0.95

7 8 22

8 x 10~6 7 x 10"6 2 x lO- 5

Type II 1970G 1979C 1980K 1981K 1986J 1988Z

0.74 0.99 0.95 1.08 0.83 0.77

6 78 23

2 x 10"4 6 x 10"4 7 x lO- 5 <2 x 10~6 2 x 10"4 1 x 10- 4

SN

6

8'

RSNe: -2.23 -2.96 -2.86 -3.24 -2.49 - 4 . 1 5 -2.30 - 3 . 8 3

8 6

that V{ ~ 1.5 X 104 km s x on day U = "15, that T ~ 105 K for the ionized wind, and that w — 10 km s" 1 , we arrive at the estimates for the massloss rates responsible for the circumstellar matter around the RSNe given in Table 2.

4 Classes of Radio Supernovae RSNe appear to fall into three main classes, based on their observed radio properties, which possibly follow a progression in circumstellar matter density, p c i r c u m (oc M/w). One of these properties is the observed radio luminosity, which is consistent with the theoretical expectation that L oc ( M / w ) ^ - 7 + 1 2 m ) / 4 (Chevalier 1982). We discuss the nature of the radio emission in light of the dependence on /t>cii-cuni-

4.1 Type la RSNe The lack of radio emission is almost certainly due to a lack of circumstellar matter originating in a red supergiant wind. This is consistent with the currently-popular model of a deflagrating/detonating white dwarf in a lowmass interacting binary system.

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S. D. Van Dyk et al.: Radio supernovae & progenitor winds 4.2 Type Ib/c

RSNe

The rough constancy in the 6 cm spectral luminosities of Type Ib/c RSNe implies similar environments for all detected events. Although three objects does not constitute a statistically significant sample, this apparent constancy also implies that they may be standard radio candles. The radio emission is assumed to be generated in circumstellar matter which originated in the wind from a hydrogen-poor red supergiant RSG (mass M £ 4 - 5 MQ) binary companion to the SN progenitor, rather than from the progenitor itself, as is assumed to be the case for Type II RSNe (see §4.3). The SN ejecta density profile of approximately /)cjecta ex r~7 indicates that Type Ib/c explosions may be similar to Type la explosions, with the difference being a higher-density presupernova stellar wind environment for Type Ib/c SNe. However, core collapse for Type Ib/c progenitors cannot be ruled out. 4.3 Type II

RSNe

The Type II RSNe are assumed to arise from single red supergiant progenitors and appear for the moment to bifurcate into "normal" (e. g., SN 1979C and SN 1980K) and "peculiar" (e. g., SN 1986.1 and SN 1988Z) types based on their radio properties. The peculiar types appear to be the result of the SN shock interacting with a very dense circumstellar medium, as is the case for the normal types, but with a dominant contribution of internal absorption, A'3. The very high luminosity for peculiar Type II RSNe can be explained by very high Mjxv for this RSN subclass. Furthermore, the smaller values of m (i. e, R <x tm, where S = — 3m) for peculiar Type II RSNe follows from the SN shock rapidly losing energy while being decelerated through interaction with AM £ 2 M& of circumstellar matter. We suspect, however, that, as more data becomes available on Type II RSNe, this bifurcation in properties may become a continuum, probably reflecting the relation of M/w with MZAMS, with SMr.j (SN 1980K) Z MZAMS < 3OM0 (SN 1986J).

5 Variations on a Central Theme As more objects become available for study and as objects are observed over long time baselines and at more frequencies, significant deviations from the simple standard model in equations (1) through (3) are apparent. Weiler et al. (1992a) did a detailed study of the deviations of the observed radio light curves for SN 1979C from the standard model and found that the deviations were quasi-periodic. They concluded that the modulations in the radio emission must arise from modulations in the circumstellar density,

5. D. Van Dyk et al.: Radio supernovae & progenitor winds

119

which implies that the mass-loss rate of the progenitor star was not constant during the red supergiant phase. Mechanisms which can produce this nonuniform wind include along-period, high-eccentricity, non-interacting binary system, in which the companion star gravitationally affects the wind of the progenitor, or possible He/C shell flashes within the progenitor. Weiler et al. (1992b) find other possible wind density inhomogeneities for SN 1986J. The recent SN 1993J (see Green & Pooley and Weiler et al., this volume) has become the most studied RSN in history, and because of the extensive coverage, both in time and radio frequency, it appears that deviations from standard models may exist in the data for this supernova as well. Acknowledgements Basic research in Radio Interferometry at the Naval Research Laboratory is supported by the Office of Naval Research through funding document number N00014-93-WX-35012, under NRL Work Unit, 2567. References Ball, L., k Kirk, J.G. 1992, ApJ, 396, L39 Chevalier, R.A. 1981, ApJ, 251, 259 Chevalier, R.A. 1982, ApJ, 259, 302 Chevalier, R.A. 1984, Ann. N. Y. Acad. Sci., 422, 215 Cowan, J.J., k Branch, D. 1985, ApJ, 293, 400 Gottesman, S.T., Broderick, J.J., Brown, R.L., Balick, B., k Palmer, P. 1972, ApJ, 174, 383 Ryder, S., Staveley-Smith, L., Dopita, M., Petre, R., Colbert, E., Malin, D., k Schlegel, E. 1993, ApJ, 416, 167 Staveley-Smith, L. Manchester, R.N., Kesteven, M.J., Campbell-Wilson, D., Crawford, D.F., Turtle, A.J., Reynolds, J.E., Tzioumis, A.K., Killeen, N.E.B., k Jauncey, D.L. 1992, Nature, 355, 147 Turtle, A.J., Campbell-Wilson, D., Bunton, J.D., Jauncey, D.L., Kesteven, M.J., Manchester, R.N., Norris, R.P., Storey, M.C., k Reynolds. J.E. 1987, Nature, 327, 38 Van Dyk, S.D., Sramek, R.A., Weiler, K.W., k Panagia, N. 1993a, ApJ, 409, 162 Van Dyk, S.D., Weiler, K.W., Sramek, R.A., k Panagia, N. 1993b, ApJ Letters, 419, L69 Van Dyk, S.D., Weiler, K.W., Sramek, R.A., Panagia, N., fc Murata. K. P. 1994, in preparation Weiler, K.W., Sramek, R.A., Panagia, N., van der Hulst. J.M., & Salvati, M. 1986, ApJ, 301, 790 Weiler, K . W , Panagia, N., Sramek, R.A., van der Hulst, J.M., Roberts, M.S., k Nguyen, L. 1989, ApJ, 336, 421 Weiler, K.W., Panagia, N., k Sramek, R.A. 1990, Ap.I, 364, 611 Weiler, K.W., Van Dyk, S.D., Panagia, N., Sramek, R.A.. k Disccniia, J.L. 1991, ApJ, 380, 161 Weiler, K.W. Van Dyk, S.D., Pringle, J.E., k Panagia, N. 1992a, ApJ, 399, 195 Weiler, K. W., Van Dyk, S. D., Sramek, R. A., & Panagia. N. 1992b, Bull. Astr. Soc. Amer. (BAAS), 24, 1293 Yin, Q.F., k Heeschen, D.S. 1991, Nature, 354, 130

Circumstellar Interaction in Supernovae Claes Fransson Stockholm Observatory, S-133 36 Saltsjobaclen, Sweden

Abstract The observational evidence for circumstellar interaction from radio, optical, UV and X-rays are briefly summarized. The basic hydrodynamical and radiative processes are reviewed and applied to the early and late phases of Type II supernovae. Particular emphasis is put on the recent SN 1993J.

1 Introduction Circumstellar interaction has turned out to be of crucial importance for the interpretation of observations of supernovae at both early and late stages. Much of the progress in this field is a result, of the combination of radio, optical, UV and X-ray observations. Here I review the basic evidence for circumstellar interaction, some of the most important physical processes, and specific examples at both early and late, stages in the supernova evolution. For a complementary review see especially the excellent review by Chevalier (1990). 2 Observational evidence for circumstellar interaction The first evidence that circumstellar interaction is important for supernovae came from observations of SN 1979C. UV observations during the first few weeks showed a number of emission lines, interpreted as a result of circumstellar interaction (§4.3). Unambiguous evidence for circumstellar interaction came from radio observations more than a year later, showing a wavelength-dependent turn-on of the radio emission (Sramek & Weiler 1990; Van Dyk, this volume). Emission was first seen at short wavelengths, and later at longer. This behavior is interpreted as a result of decreasing freefree absorption by the ionized gas in a circumstellar medium around the supernova (Chevalier 1982b). As the supernova expands, the optical depth Tjf = J^° Kfjneriidr, from the radio emitting region close to the shock through the circumstellar medium decreases, explaining the radio turn-on. The most natural origin for the circumstellar medium is the wind from the supergiant progenitor (or possibly companion star). Mass loss rates estimated for galactic supergiants are highly uncertain. Compilations by Jura 120

C. Fransson: Circumstellar interaction in supernovae

121

&; Kleinmann (1990) and Knapp & Woodhams (1993) give wind velocities in the range uw = 10-30 km s"1 and mass loss rates 3x 10~ 7 -5x 1O~5M0 yr" 1 . Assuming a fully ionized wind with constant mass loss rate and velocity, p = M/4nuwr^. The free-free optical depth at wavelength A is r,,(A) « 7.1 x 102A2 ( ^ ) 2 T 5 - 3 / 2 K r 3 ^

(1)

where M-5 is the mass loss rate in units of 10~5 MQ yr" 1 , uw\ the wind velocity in units of 10 km s"1, and T$ the temperature of the circumstellar gas in 105 K. From the radio light curve, or spectrum, the epoch of TJJ = 1 can be estimated for a given wavelength, and from the line widths in the optical spectrum the maximum expansion velocity, V, can be obtained. Because the effects of the radiation from the supernova, has to be estimated from models of the circumstellar medium, the temperature in the gas is uncertain. Lundqvist & Fransson (1988) find that initially the radiation heats the gas to Te « 105 K. Te then decreases with time, and after a year Te « (1.5 — 3) X 104 K. In addition, the medium may recombine, which further decreases the free-free absorption. From i(T(\)jj = 1) the ratio of M/uw can be calculated. Because M/uw a Te x~^, errors in Te and xe may lead to large errors in M. If the medium is dumpy this may lead to an overestimate of M/uw. The mass loss rates for 'normal' Type II supernova.? obtained in this way are between 3 x 10~5 and 2 x 10"'M 0 yr" 1 . These are at the high end of those determined for red supergiants, but given their sensitivity to the temperature and clumping of the medium, it is premature to draw any strong conclusions from this. SN 1987A showed an initial weak burst of radio emission (Turtle et al. 1987). From the light curve Chevalier & Fransson (1987) estimated M/xiw « 1O~8M0 yr" 1 km"1 s. The low value of M/uw is probably a. direct consequence of the blue B3Ia progenitor, which had a. wind with uw ss 550 km s"1, much higher than for a. red supergiant. The resulting mass loss rate is then rather normal, ~ 3 x 1O~6M0 yr" 1 . No Type la's have been detected in radio (Sramek & Weiler 1990). If Type la's occur in binary systems with a late type companion, radio emission from the companion wind might be expected. The absence of radio emission argues against this. Mass loss rates for the Type Ib's are in general lower by up to an order of magnitude, compared to the Type IPs (Van Dyk et. al. 1993a). This is consistent with the fact that these are believed to originate from explosions

122

C. Fransson: Circumstellar interaction in snpernovae

of WR-like stars, probably in binary systems (e.g., Yamaoka & Nomoto 1991). The winds from the progenitors are therefore likely to be ~ 1000 — 2000 km s" 1 . Additional signatures of circumstellar interaction are discussed in §4.

3 Physical processes 3.1 Hydrodynamics The interaction of the ejecta, expanding with velocity > 104 km s" 1 , and the nearly stationary circumstellar medium results in one reverse shock wave propagating inwards (in mass), and one outgoing circumstellar shock. The density in the circumstellar gas is pcs = M/4iruwr2, if the mass loss rate and wind velocity are constant. After the first few days the outer parts of the ejecta expand with constant velocity, V(m) a r for each mass element,TO,SO that r(m) = V(m)t and p(m) — po(m){t0/t)3. Hydrodynamical calculations show that to a good approximation the ejecta density can be described by a power law in radius, pej = po{t I to)~3{Vot I r)n •> with n « 7 — 12. A useful similarity solution can then be found (Chevalier 1982a, b; Nadyozhin 1985). Here we sketch a simple derivation. More details can be found in these papers, as well as in the review by Chevalier (1990). Assume that the shocked gas can be treated as a thin shell with mass Ms, velocity Vs and radius Rs. Balancing the ram pressure from the circumstellar gas and the impacting ejecta, the momentum equation for the shocked shell of circumstellar and ejecta is Ms^-

= 4nR2s(pejV2ev -

PcsV

2

)

= 4rrR][pcj(V - Vs)2 -

2 PcsV }.

(2)

Here, Ms is the sum of the mass of the shocked ejecta. and circumstellar gas. The swept up mass behind the circumstellar shock is MC3 = MRa/uw, and that behind the reverse shock Mrev = 4ir t^V"[t/Ra)n~3/(?i — 3), assuming that Rs >> Rp, the radius of the progenitor. With V = Rs/t we obtain M D„ , 4vpot3o Von r - 3 l d2Rs 1 : uw s (n (71 - 3) 6) Rs Ks " J dt* df . [n. li3 VVn l f"-3

,2 Po o

o

J

/

Rsft t

This equation has the power law solution l/(n-2) 7^TT n - T~ 1/ " II

R.(t) = l(n - 4)(n - 3) M\

V2 )

,11) \ 2 dR s' dt

\,f /,}/? (dR M s 2 4TT uw R \ dt

\ 2"|

(3)

C. Fransson: Circumstellar interaction in supernovae

123

The form of this similarity solution can be written down directly by dimensional analysis from the only two independent quantities available, p0 t\ Von and M/uw. This solution applies after a few expansion times, when the initial radius has been 'forgotten'. To justify the neglect of the initial conditions, n > 5. More accurate similarity solutions, taking also the structure within the shell into account, are given by Chevalier (1982a) and Nadyozhin (1985). In general, these solutions differ by less than ~ 30% from the thin shell approximation. The maximum ejecta velocity close to the reverse shock depends on time The velocity of the circumstellar shock, dRs/dt, as V = Rs/t oc rlHn-2\ in terms of V is Vs — V(n — 3)/(rc — 2) and the reverse shock velocity, Vrev = V — Vs = V/(n — 2). Assuming equipartition between ions and electrons, the temperature of the shocked circumstellar gas is

and the reverse shock T — "(n-3)»The time scale for equipartition between electrons and ions is ~ 2.5 X 107 (T e /10 9 K)1-5 (n e /10 7 c m " 3 ) - 1 s. One finds that the reverse shock is marginally in equipartition, unless the temperature is Z 5 X 108 K. The ion temperature behind the circumstellar shock is ^ 6 x 109 K for V4 j£ 1.5, and the density a factor ^ 4 lower than behind the reverse shock. Ion-electron collisions are therefore ineffective, and Te << Twn, unless efficient plasma instabilities heat the electrons collisionlessly (fig. 1). The electron temperature of the two shocks will therefore be highly different, ~ (1 - 3) x 109 K for the circumstellar shock and 107 - 5 x 108 K for the reverse shock, depending on n. The radiation from the reverse shock is therefore mainly below ~ 20 keV, while that from the circumstellar shock is above ~ 50 keV. The density behind the reverse shock is prev _ 4TT uwp0t3oVontn-3 pcs

n

AI R ~

_ (n - 4)(n - 3) £

and is therefore for n £ 7 much higher than behind the circumstellar shock. In addition, the similarity solution by Chevalier (1982a) shows that there is a strong density gradient from the reverse shock to the contact discontinuity between the shocked ejecta and circumstellar gas (fig. 1). This reflects the higher ejecta density at the time when these mass elements

124

C. Fransson: Circumstellar interaction in supernovae

2.8x10"

2.8x10

3X10"

Radius (cm)

Fig. 1. Density and temperature structure of the reverse and circumstellar shocks for n — 1 and a velocity of 2.5 x 10'1 km s" 1 at 10 days. Note that because of the slow Coulomb equipartition the electron temperature (dotted line) is much lower than the ion temperature (full line) behind the circumstellar shock.

were shocked. However, this structure is hydrodynamically unstable. The negative density gradient, in combination with a positive pressure gradient between the contact discontinuity and circumstellar shock, leads to a Rayleigh-Taylor instability. Chevalier, Blondin & Emmering (1992) have calculated the structure using a 2-D PPM hydro code. They indeed find that instabilities develop, with dense, shocked ejecta gas penetrating into the hotter, low density shocked circumstellar gas. The instability mainly distorts the contact surface, and does not affect the general dynamics seriously. The calculation assumed that cooling is not important. If the gas cools efficiently one can imagine dense, cool blobs penetrating the unshocked circumstellar gas in front of the shock.

3.2 Radiative emission During the first month radiation from the supernova photosphere is strong enough for Compton scattering to be the main cooling process for the circumstellar shock. The photospheric photons have energies ~ 3 kTejj « 1-10 eV. The optical depth to electron scattering behind the circumstellar shock is

C. Fransson: Circumstellar interaction in supernovae

125

re = 0.18 M.5 u~\ V'1 Jt^ys.

(8)

A fraction r^ of the photospheric photons will scatter N times in the hot gas. In each scattering the photon increases its energy by a factor Av/v « 4 kTe/mec2 ^ 1. The multiple scattering creates a power law continuum that may reach as far up in energy as the X-ray regime. If relativistic effects can be ignored (Te ^ 109 K), the spectral index is

a=

{i ~ Tif ln[ I ( 0 - 9 2 2 8 " I n Te)]}12" \

(9)

(Fransson 1982). Typically, 1 < a < 3. For Te > 109 K relativistic effects become important and increases the cooling considerably (Lundqvist & Fransson 1988). One can estimate the free-free luminosities of the two shocks from

U = t* J MTMtdTHJssWtj}^.

. (10)

The density behind the circumstellar shock is pcs = 4 po = M/{•KUWR2S). The swept up mass behind the circumstellar shock is Mcs = MRs/uw and that behind the reverse shock is Mrev = (n — 4)M c s /2. With jjj — 2.4 x r e 0 5 , we get

£,-« 3.0 X 1039 gfJ Cn [^lYf-L-Y

ergs"1,

(11)

where gjj is the free-free Gaunt factor, including relativistic effects. For the reverse shock Cn = (n — 3)(??. — 4)2/4(??. — 2), and for the circumstellar shock Cn — 1. This assumes electron-ion cqnipa.rtit.ion. which is highly questionable for the circumstellar shock. Because of occultation by the ejecta only half of the above luminosity escapes outward. At Te ^ 2 x 107 K line emission increases the cooling rate, j w 3.4 X 23 10~ T~70'67 erg s~ 1 cm 3 . If the reverse shock temperature falls below 2 X 107 K a thermal instability may occur and the gas cool to ^ 104 K, where photoelectric heating from the shocks balances the cooling. Most important, the cool gas may absorb most of the emission from the reverse shock. If cooling, the total energy emitted from the reverse shock is Lrev

=

47Ti( s -2 P eiVrr,, ^3'rev-

=

(n - 3)(n - 4) 4

(

n

_ 2TT^ ) 3

J

3 7MV

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Because V oc /- 1 /( n ~ 2 ) ) Lrev a ^~3/(n~2) in the cooling case. For high M/uw the luminosity from the reverse shock may contribute appreciably, or even dominate, the bolometric luminosity. 4 Early circumstellar interaction The earliest form of circumstellar interaction occurs at shock breakout. As the shock approaches the surface, radiation leaks out on a time scale of less than an hour. The color temperature of the radiation is ~ (1 — 5) X 105 K and the energy ~ (1 - 10) x 1048 ergs (Falk 1978; Klein & Chevalier 1978; Ensman & Burrows 1992). This burst of EUV and soft X-rays ionizes and heats the circumstellar medium on a. time scale of a few hours. In addition, the momentum of the radiation may accelerate the circumstellar gas to a high velocity. Most of the emission at energies £ 100 eV is emitted during the first few hours, and after 24 hours little ionizing energy remains. After about one expansion time (~ Rp/V) the reverse and circumstellar shocks are fully developed, and the radiation from these will dominate the properties of the circumstellar gas. The fraction of this emission going inward is absorbed by the ejecta and there re-emitted as optical and UV radiation (Fransson 1982; 1984). 4.1 SN 1981A The primary example of the influence of the soft X-ray burst is the echo from the ring of SN 1987A. The properties of the circumstellar medium and the emission from the ring are discussed in detail in the contribution by Lundqvist in this volume. The UV light curves observed with IUE up to day 2000 are presented in Fransson & Sonneborn (1994). General reviews of SN 1987A can be found in Chevalier (1992), Lundqvist (1992) and McCray (1993). Here only the basic scenario is summarized. Ensman & Burrows (1992) have calculated the burst in detail, and find a peak effective temperature of ~ 6 x 105 K and a color temperature of ~ 1 X 106 K. In total ~ 7 X 1046 ergs were emitted above 13.6 eV during the first day. Photons above ~ 100 eV, capable of producing e.g. N V, are emitted during less than ~ 500 sees. At shock breakout the radiation ionized the pre-existing ring on a time scale of seconds (Lundqvist & Fransson 1991). The temperature in the gas was then ~ 105 K. With a ring geometry the gas is optically thick in the continuum, and distinct ionization zones develops. Subsequently, the gas cooled and recombined. After a year the temperature was ~ 5 X 104 K, in agreement with temperatures determined

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from the [0 III] ratio (Wampler & Richichi 1989). The distance of the ring from the supernova is R — 6.3 x 10 1 ' cm, and its inclination i as 43° (Panagia et al. 1991). Therefore, after 2 sin i R/c w 410 days the full ring will be seen. This marks the maximum in the light curves. The recombination/cooling scenario, together with the light echo effects, explain the shape of the observed light curves. By modelling the state of ionization Fransson & Lundqvist (1989) found that the radiation temperature of the burst should be in the range (3 - 6) x 105 K. The ejecta are predicted to collide with the ring in the year ~ 2000 ± 3, with mainifestations in radio, IR, optical, UV and X-rays (e.g., Luo, McCray & Slavin 1994).

4.2 SN 1993J The recent SN 1993J is a good example where emission from the shock breakout has observable consequences. The first detection of the supernova was on March 28.30 at magnitude 13.6 (Merlin & Neely 1993). Hydrodynamical models (Shigeyama et al. 1994) show that shock breakout should have occurred at approximately March 28.0. The first spectra taken with IUE (Wamsteker et al. 1993; Sonneborn et al. 1993; Fransson & Sonneborn 1994; Sonneborn this volume) on March 30.21, gave a photospheric temperature of 22,500 K. One day later on March 31.2 the temperature had decreased to 14,500 K. Such rapid cooling is expected just after shock breakout, and agrees well with light curve calculations by Shigeyama et al. (1994). Shigeyama el al. 's model shows that immediately after breakout T e // « 3 x 105 K, decreasing in 24 hours to ~ 4 x 104 K. The total energy during the first day was ~ 6 X 1018 ergs, most coming out as UV and soft X-rays. The optical spectrum displayed a very rapid evolution (cf. Meikle et al. 1993; Filippenko, Matheson, & Ho 1993). From having a nearly featureless blue continuum, very broad lines became apparent after ~ 15 days. On April 11, expansion velocities of at least 19,000 km s" 1 can be seen for Ha (Meikle et al. 1993). In early May the spectrum changed character from an H-dominated Type II spectrum to a spectrum dominated by lines of He, resembling a Type Ib. From this it was apparent that only a small fraction of the hydrogen envelope remained on the progenitor at the time of the explosion. Based on the bolometric light curve several groups have concluded that the mass of the hydrogen envelope was ;$ IMQ (Nomoto et al. 1993; Podsiadlowski et al. 1993; Utrobin 1993; Woosley et al. 1994). Most of these models invoke a. binary scenario, although Hoflich, Langer, &

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Duschinger (1993) propose that the mass loss is by a stellar wind from a single star. With a ZAMS mass of ~ 15M 0 , about 1OM0 must have been lost from the progenitor, either in the form of a wind, to a companion star, or by outflow in an excretion disk. Besides the hot continuum, the most obvious feature in the UV spectrum during the first week was the N V AA 1238.8 - 1242.8 doublet in emission. High resolution observations showed a line width ~ 103 km s" 1 , implying a circumstellar origin. The lines are very strong in the first spectra, but weaken rapidly with time. On March 30.2 the luminosity was ~ 2x 1040 erg s" 1 . Between March 30.2 and April 4.5 the N V flux decreased to ~ 5 x 10 38 erg s" 1 . By April 22 the N V emission was undetectable. The presence of the N V line shows that dining the first few days there was a high flux of photons above 77 eV. Models by Fransson, Lundqvist & Chevalier (1994) (henceforth FLC94) show that the burst from the shock breakout ionizes the circumstellar gas nearly completely. Because of the high density, gas inside ~ 1015 cm has time to recombine before the shock hits the gas, explaining the presence of the line. Once the reverse and circumstellar shocks are formed, radiation from these re-ionize the gas completely, effectively cutting off the emission. The N V line, as well as Ha (Cumming et al. , this volume), shows evidence for circumstellar velocities of ~ 1000 km s" 1 . Fransson & Sonneborn (1994) and FLC94 propose that this can be explained as a result of pre-acceleration of the wind by the strong UV burst at shock break-out. The UV and optical were the first wavelength ranges to show evidence for a circumstellar medium. A few days later radio emission was also detected from SN 1993J (Weiler et al. 1993; Pooley & Green 1993), as well as X-rays (Zimmermann et al. 1993a; Tanaka 1993). The radio emission showed the same characteristic pattern as had previously been observed for other Type II and Type Ib supernovae. First, emission was seen at short wavelengths, and later at longer. The radio flux at a given wavelength increased roughly linearly with time up to a maximum, and remained nearly constant up to at least 100 days (Pooley & Green 1993). From the radio spectrum of Van Dyk et al. (1993b) and Phillips & Kulkarni (1993) on April 22.5 FLC94 estimate that TJJ = 0.3 at 2 cm on day 25.5. Because of the heating of the circumstellar medium by the outburst of the radiation, as well as by the shocks formed by the interaction of the ejecta and the circumstellar medium, FLC94 find from photoionization models of the circumstellar medium that the temperature was ~ 3 x 105 K. The expansion velocity on day 25.5 was ~ 20,000 km s" 1 . From equation (1) one finds M/uw « 3 x 1O" 6 M 0 yr" x ( km s" 1 )" 1 , or M = 3 x 1O~ 5 M 0 yr" 1 if

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uw = 10 km s" 1 , within approximately a factor two. This is of the same order as derived for other 'normal' Type II supernovae (§2). Although qualitatively correct, the form of the observed light curves do not agree with the expectations (Lundqvist et al. in Wheeler & Filippenko 1994; FLC94; Lundqvist, this volume). In particular, the rising part is nearly linear, while models predict a more abrupt turn-on. The rising part of the light curves depends mainly on the density distribution of the circumstellar medium. FLC94 find that while a standard p oc r" 2 wind does not fit the observations by Pooley & Green, a better fit is obtained for p oc r~ 1 5 . Exploiting this idea, Weiler et al. (this volume) find that this also fits the observations at other wavelengths. The departure from p oc r~2 may be a result of variations in the mass loss rate during the last ~ 1000 years before the explosion, or due to the influence of the binary companion. SN 1993J has been observed by ROSAT (Zimmermann et al. 1993a,b) in the 0.1 - 2.4 keV band and by ASCA (Tanaka et al. 1993) between 1 10 keV. The luminosity at the first observations, ~ 7 days after explosion, was 1.6 X 1039 erg s" 1 between 0.1 - 2.4 keV, and 5 x 1039 erg s" 1 between 1 - 1 0 keV. There is evidence for a decrease of ~ 45% from April 3.4 to May 4 in the ROSAT flux (Zimmermann et. al. 1993c). On November 2 - 3 1993, on day 220, ROSAT observed a remarkably soft flux with kTx « 0.5 keV (Zimmermann et al. 1993d). At the same time the column density had increased by a factor ~ 6. The November observation strongly indicates that the flux was then coming from the reverse shock. Equation (6) can be written as k,Trev — 117 F 4 2 / ( n - 2 ) 2 keV. With V4 > 1.5 the density gradient must have n > 25. This implies that the reverse shock is radiative, and therefore has a dense cool shell, absorbing most of the X-rays from the reverse shock. The X-rays during the first months must therefore come from the circumstellar shock, explaining the much higher temperature, kTx ^ 10 keV, at these epochs. Without collisionless heating Te « 2 x 109 K. Using equation (11) with gjj « 2 the outgoing luminosity at 10 days is ~ 3 x 10 39 (A/_. 5 /v u ,i) 2 erg s" 1 , with a fraction <7//(5 keV) AE/gjf kT\ ~ 0.1 in the ASCA band. Therefore, we find that M w 4 x 10~5 MQ yr" 1 is indicated by the X-ray observations. The ASCA/ROSAT ratio is ~ (9.0/2.3) [ > 10 keV. In FLC94 a more detailed estimate of the luminosity is given, taking into account departures from electron-ion equipartition and the hydrodynamic structure. A similar discussion, assuming the X-rays come instead from the reverse shock, is given by Suzuki et al. (1993). The decrease in luminosity with time is rather slow compared to the expected, L oc Te°cS /-< n - 3 )/( n - 2 ). Again, a p a r~1>5 distribution provides a

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better fit to the observations. As the column density of the cool, absorbing gas becomes thinner, N(H) oc Rj'1, the soft flux from the reverse shock is expected to increase, as long as the shock is radiative. Gamma-rays from SN 1993J have also been detected by OSSE on the Compton Observatory (Leising 1993). At ~ 12 days the 40 - 200 keV flux corresponded to ~ 7 x 1040 erg s" 1 . FLC94 find that to explain the OSSE/ASCA flux ratio, Te > 2 X 109 K, indicating that collisionless heating of the electrons is probably necessary. A possible problem is that the high temperature gives a Comptonized flux in the ROSAT and ASCA ranges higher than observed, and with too steep a spectrum. The total mass swept up during this first half-year by the supernova is only M t V/uw « 0.02 MQ. Emission from the circumstellar interaction may therefore be observable for a long time, as is the case for a number of other Type II supernovae (§5).

4-3 Other

supernovae

SN 1979C was the first supernova for which circumstellar interaction was found to be important. IUE spectra showed strong emission lines from He II, C III-IV, N III-V, and O III (Panagia et al, 1980; Fransson et al. 1984). In addition, a strong UV excess was found. Fransson (1982; 1984) interpreted the UV excess as Comptonized emission of the photospheric radiation by the circumstellar shock. The lines were explained as a result of the inwardgoing Compton emission in the EUV, photoionizing the outer ejecta, and possibly the cooling, shocked ejecta. The decay of the UV lines followed the decrease in the photospheric flux feeding the Comptonization, in agreement with the observations. The modelling of the radio observations implied a mass loss rate of ~ 1.5 x 10~4MQ yr" 1 (Lundqvist & Fransson 1988). In spite of this high mass loss rate, no X-rays were observed (Palumbo et al. 1981), probably because most of the X-ray emission was absorbed by the cool shell between the reverse shock and contact surface. The Type II-L SN 1980K was the first supernova to be detected in X-rays (Canizares, Kriss, & Feigelson 1982). About 40 days after outburst it was detected by Einstein, with a luminosity of ~ 2 x 1039 erg s" 1 between 0.2 - 4 keV. A month later it was again observed but only marginally detected with less than half the previous luminosity. This supernova was also well observed in radio (Weiler et al. 1992), and from the radio turn-on a mass loss rate of ~ 3 x 1O~5M0 yr" 1 has been estimated (Lundqvist & Fransson 1988). The X-ray luminosity derived for this mass loss rate agrees well with that observed.

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IR emission has been observed from several Type II supernovae. Both SN 1979C and SN 1980K were detected (Bode & Evans 1980; Dwek 1983). Dust is assumed to be the main accelerating component in late type winds, and dust emission is therefore not surprising. The radiation from the supernova will evaporate the dust inside a radius ReVap ~ 10 17 — 10 18 cm, depending on the luminosity of the supernova and the evaporation temperature, 1000 — 2000 K (Dwek 1983). The heated dust grains outside Revap produce a light echo of a duration ~ 2 Revap/c. The mass loss rates inferred from the models agree well with those obtained from the radio. Graham & Meikle (1986) modelled the IR observations of the Type la SN 1982E (Graham et al. 1983) with a dust echo and find that a good fit to the data can be obtained for a dusty wind with M ss 1O~5M0 yr" 1 , and suggest that the binary companion was a low mass AGB star with heavy mass loss. A narrow Ha component in SN 1984E was observed by Dopita et al. (1984), who estimated a mass loss rate as high as ~ 3 x 1O~5M0 yr" 1 . This estimate is, however, uncertain, since it assumes pure recombination for the Ha line. This is clearly not the case, as is indicated by the P-Cygni profile. Also this object showed evidence for pre-acceleration.

4-4 Abundances From the UV lines it has, in two cases, been possible to determine the relative abundances of the CNO elements. In SN 1979C the ratios of the broad N III] A 1750 and C III] A 1909 lines gave a relative abundance ratio N/C « 8 by number (Fransson et al. 1984). For the O/N ratio only an upper limit of ~ 0.5 could be determined. The relative fluxes of the narrow lines from the ring of SN 1987A resulted in N/C « 7.8 and N/O « 1.6 (Fransson et. al. 1989). The nitrogen enrichments implied by these observations clearly indicate that the observed gas has undergone CNO-processing. This in turn, implies that substantial mixing must have occurred in the progenitor. The large nitrogen enhancement found for SN 1987A has turned out to provide one of the main constraints on the evolution of the progenitor. For reviews of various models see e.g. Podsiadlowski (1992).

5 Very late emission from supernovae The persistent radio emission shows that many supernovae, even after several years, are interacting with a circumstellar medium. The radio flux of SN 1979C has, for example, only decayed by a factor ~ 2 over the more than

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ten years it has been observed (Weiler et al. 1991). Recent observations of several supernovae, reviewed by Leibundgut in this volume, show that circumstellar interaction is also important for the optical emission at late stages. Here I only mention some of the implications, discussed in Chevalier k Fransson (1994), henceforth CF94. As the supernova ejecta expand they sweep up the circumstellar medium. At the same time the reverse shock propagates further into the ejecta. The maximum ejecta velocity therefore decreases continuously. Hydrodynamical calculations show that the density profile flattens considerably for ejecta velocities less than ~ 5000 km s" 1 . As a first approximation we may use a simple power law model, but with an n that may be different from that inferred from the early epochs. The maximum ejecta velocity, i.e. the widths of the line profiles, is then /

v = Vn

\WT)

\ V(i-2)

^/{n~2)

km8 1

"'

where V6 = 12.1 x 103, V8 = 9.61 x 103, Vi0 = 7.44 x 10 3 , and V12 = 6.51 X 103 km s" 1 . Here po-\e is a reference density at one year and at 5000 km s" 1 in units of 10~ 16 g cm" 3 . Hydrodynamical models show that n » 8 and p o _ 1 6 » 1. With M_ 5 = 5.0, we find V % 5000 km s" 1 at 10 years, in agreement with the observed velocities. The shock emission is dominated by the reverse shock with velocity ~ 5000/(n—2) « 700—1500 km s" 1 . The temperature behind the reverse shock is therefore 1 0 6 - 107 K. Below ~ 4 x 106 K the shock spectrum is dominated by far-UV and X-ray lines below ~ 100 cV. At higher temperatures Xrays between 0.5 - 1 . 0 keV dominate. The low temperature makes cooling important for the reverse shock. This can lead to a thermal instability, cooling the gas to ^ 104 K, where photoelectric heating may balance the cooling. The result is a cool, dense shell between the reverse shock and contact discontinuity, opaque to all X-rays below 0.2 — 1 keV. Therefore, most emission from the reverse shock will be absorbed, half in the ejecta and the other half in the cool shell (fig. 2). Photoionization models show that the energy absorbed by the ejecta is transformed into Lycv, Ha and emission from highly ionized species, like [0 III]AA 4959 - 5007, C III] A 1909, and C IV AA 1550. The emission from the cool shell is dominated by Lycv, Ho and Mg II A 2800. The temperature in the ionized ejecta is ~ (15 - 20) X 103 K, while that of the cool shell is only ~ 8000 K. The ionized zone in the ejecta is confined to the region just inside the reverse shock with AR/R a 0.1.

C. Fransson: Circumstellar interaction in supernovae

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The most interesting objects to compare the model to are SN 1979C and SN 1980K. Both are strong radio supernovae, and are fairly typical Type II's. SN 1980K was the first to be recovered optically, and among the 'normal' cases is the best observed (Fesen & Becker 1990; Leibundgut et al. 1991). The optical spectrum was dominated by Ha, [0 I] AA 6300 — 64, and [0 III] A A 4959 — 5007. These lines are also the strongest in the models. In addition, both in the UV and IR the models show several lines with similar or larger fluxes. Observations in these ranges should therefore be rewarding. The [0 III] AA 4959 — 5007/Ha ratio shows an interesting evolution (fig. 3). Up to ~ 2 years the [O III] lines are collisionally de-excited. As the ejecta density decreases the lines become increasingly strong. At an age of ten years or more they dominate the cooling of the ejecta. This is consistent with observations of SN 1980K, SN 1979C (Fesen & Matonick 1993) and the 30-year-old SN 1957D. In the last of these, the [O Til] lines are by far the strongest (Turatto et al. 1989; Long et al. 1989). The decline of the Ha light curve of SN 1980K was consistent with radioactive decay from 56 Co during the first year or so. At the time of the last observation by Uomoto & Kirshner (1986) at 670 days, Chugai (1988) found a clear excess, and proposed that the origin of the excess was circumstellar interaction, based on a suggestion by Chevalier & Fransson (1985). Subsequent observations at 7.8 years and later show a nearly constant Ha flux. This indicates that the energy source is nearly constant. As we saw in §3.2, this is expected if the reverse shock is radiative, Lrev oc i~ 3 /( n ~ 2 ). Whether the reverse shock is radiative or not is sensitive to both the density gradient, n, and the mass loss rate. Most emission originates either in the cooling shell or from the outer parts of the ejecta close to the reverse shock, and the lines are expected to have flat, box-like profiles, with extent V « (4 - 5) x 103 km s" 1 , as is seen in SN 1980K and SN 1979C. The observed profiles, however, tend to be asymmetric with a stronger blue edge. A possible reason for this may be dust absorption in the ejecta. While SN 1957D, SN 1979C and SN 1980K are fairly well behaved, and fit well into the scenario discussed above, there is a number of objects which have a more peculiar behavior. These include SN 1978K, SN 1986J, and SN 1988Z, and are discussed by Leibundgut (this volume). This class probably only represents a small fraction of the total number of Type II supernovae. Most of them have been discovered through a variety of peculiarities; SN 1978K and SN 1986J due to their extraordinary radio emission, and SN 1988Z because of its peculiar optical spectrum (in spite of cz ~ 6000 km s" 1 !).

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C. Fransson: Circumstellar interaction in supernovae

shocked ejecta T * 1CT K cool shell

Lya, Ha, Mg U ionized wind

shocked wind T * 3x1 Q8K

Fig. 2. Schematic figure showing the different regions of the supernova - circumstellar medium interaction, and the most prominent lines from the various regions (from CF94).

J

5

0

1

'

1

' ' "1

4959-5007

IS

10

1

-

-

J

y .

i

, ,,

5 t (years)

-

11

10

i 20

,-

Fig. 3. Evolution of the [ 0 III] A A 4959 —5007/ITo ratio as a function of time (from CF94).

C. Fransson: Circumsiellar interaction in supernovae

135

SN 1986J was discovered as a radio source, with a radio luminosity of ~ 3000 times Cas A. Modelling the light curve Weiler, Panagia & Sramek (1990) estimate the explosion date to have been September 1982. VLBI observations (Bartel et al. 1991) show the source as an irregular shell, whose velocity is ~ 8000 km s" 1 , with 'protrusions' reaching out to ~ 18,000 km s" 1 . The optical spectrum shows no evidence of lines with these kinds of velocities. Instead, two systems of lines are present, with Ha having a FWHM velocity of ~ 530 km s" 1 , and [0 I] - [0 III] lines having higher velocities, 1000 - 2 0 0 0 km s" 1 . A likely explanation for this apparent inconsistency is that the low velocity H lines originate in shocked circumstellar gas, while the high velocities inferred from the radio may be due to the expanding blast wave. These aspects can be combined if the circumstellar medium is very clumpy, with high density clouds immersed in alow density wind (Chugai 1993). Alternatively, the circumstellar gas may be asymmetrically distributed, with low density in the polar direction and high density in the equatorial plane, possibly in the form of a circumstellar disk. In either case, the II lines originate from shocked circumstellar gas. To obtain the low velocities, the circumstellar gas density has to be higher than the ejecta by a factor pCs/Pej ~ [V /Vcs)2 * 100. The high 0 I density estimated by Leibundgut, el al. (1991), ~ 109 cm" 3 , makes it likely that the 0 lines originate in the oxygen-rich regions of the ejecta. SN 1986J has also been detected as an X-ray source with ROSAT, with a luminosity of (1.6 - 7) x 1040 erg s" 1 between 0.1 - 2.5 keV, and kTx = 1.0 - 3.9 keV (Bregman & Pildis 1992). The high column density found is consistent with absorption by cooling gas. SN 1988Z showed similar characteristics: strong radio emission (Sramek, Weiler & Panagia 1990), and narrow optical lines at late stages (Stathakis & Sadler 1991; Turatto et al. 1993). This supernova was discovered shortly after explosion. During the first weeks SN 1988Z showed a broad Ha emission with a velocity width of ~ 20,000 km s" 1 , as well as a narrow component with FWHM ;$ 200 km s" 1 . The broad component faded rapidly, and instead an intermediate component with FWHM % 1500 km s" 1 appeared. Stathakis & Sadler (1991) find that the bolometric luminosity was higher than for normal Type IPs, and suggest that this is due to the circumstellar interaction. Chugai (1992) found that the light curve could be fitted for a mass loss rate M/uw\ « 3 x 10~3Mi) yr" 1 . This is much higher than is likely, given the turn-on time of the radio flux. In the same way as for SN 1986J, the high mass loss rate, the broad Ha and the narrow components can, however, be explained qualitatively by an aspherical circumstellar medium.

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Recently, SN 1978K has been added to the same category (Ryder et al. 1993), having a strong radio flux, ~ 300 times Cas A, an X-ray luminosity of ~ 1040 erg s"1 with kTx ~ 0.5 keV, and narrow emission lines, FWHM ~ 450 km s~l. In January 1980 an Einstein observation gave an upper limit of 0.9 x 1039 erg s" 1 , indicating strong X-ray absorption. The mass loss is estimated to be ~ 4 x 1O~4M0 yr" 1 . An alternative scenario for the late emission is excitation by a central pulsar (Chevalier & Fransson 1992). A fundamental difference between these two models is that while the pulsar scenario predicts line profiles with a maximum width of < 1500 km s"1, increasing with time, the circumstellar interaction scenario predicts line velocities of ~ 5000 km s"1, decreasing with time. Chevalier (1987) proposed that SN 1986.J was powered by pulsar excitation. However, the line widths, the presence of strong radio emission, and X-rays all indicate circumstellar interaction in this case also. In addition, VLBI observations show that the emission originates in a shell. Acknowledgements The modeling of the observations of SN 1993J has been done in collaboration with Roger Chevalier and Peter Lundqvist. I am also grateful to Robert Cumming and Peter Lundqvist for comments. This work is supported by the Goran Gustafsson Foundation for Research in Natural Sciences and Medicine. References Bartel, N., Rupen, M. P., Shapiro, I. I., Preston, R. A., k. Rius, A. 1991, Nature, 350, 212. Bode, M.F. & Evans, A. 1980, AW 193, 21 P. Bregman, J. N., & Pildis, R. A. 1992, ApJ, 398, L107. Canizares, C. R., Kriss, G. A., & Feigelson, E. D. 1982 ApJ, 253, LIT. Chevalier, R. A. 1982a, ApJ, 258, 790. Chevalier, R. A. 1982b, ApJ, 259, 302. Chevalier, R. A. 1987, Nature, 329, 611. Chevalier, R. A. 1990, in Supernovae, ed. A. G. Petschek, Berlin, Springer, p. 91. Chevalier, R. A. 1992, Nature, 355, 691. Chevalier, R. A., & Fransson, C 1985, Supernovae as Distance Indicators, ed. N. Bartel,Berlin, Springer, p. 123. Chevalier, R. A., & Fransson, C. 1987, Natwe, 328, 44. Chevalier, R. A., & Fransson, C. 1992 ApJ, 395, 540. Chevalier, R. A., & Fransson, C. 1994, ApJ, 420, 268. (CF94) Chevalier, R. A., Blondin, J. M., & Emmering, R. T. 1992, ApJ, 392, 118. Chugai, N. N. 1988, Ap&SS, 146, 375. Chugai, N. N. 1992, Sov. Astr., 36, 63. Chugai, N. N. 1993, ApJ, 414, L101. Dopita, M., Evans, R., Cohen, M., & Schwartz, R. 1984, ApJ, 287, L69.

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2D Hydrodynamic Models of Supernova Progenitor Winds John M. Blondin Department of Physics North Carolina State University Raleigh NC 27695-8202

Abstract The conventional wisdom that a Type II supernova explosion occurs inside a spherical stellar wind bubble blown by the wind of the red snpergiant progenitor misses two important points: the progenitor wind may be time-dependent, and it may be asymmetric. These two features of SN progenitor winds have been well illustrated by the ring observed around SN 1987A. The existence of this circumstellar shell directly implies a time-dependence in the wind on time scales less than about 10,000 years. Also, the shell is undeniably asymmetric, implying some form of asymmetry in the progenitor wind(s). Some of the theories for an asymmetric circumstellar medium include gravitational focussing in a wide binary, rotationally deformed wind, colliding winds in a binary system, and asymmetric mass ejection in a common envelope or accretion phase of a close binary system. The wind dynamics of these various theories will be reviewed with an eye toward understanding the true history of Sk -69°202 .

1 Introduction The standard picture of a Type II SN progenitor star is a red supergiant (RSG) that has evolved from a massive star with an initial main-sequence mass above ~ I O M Q . These RSGs are observed to have very massive, slow winds with terminal speeds in the range of 10 — 50 km s" 1 , and mass loss rates in the range of 10~7 - lO~ 5 M0yr" 1 . These slow winds will gradually blow a stellar wind bubble of RSG wind into the relic main-sequence stellar wind bubble, building up a shell of shocked RSG wind at the edge of the expanding bubble. Given the typical lifetimes for this RSG stage, this RSG wind bubble is expected to reach a radius on the order of 10 pc. If one assumes that the RSG wind remained at constant velocity and mass loss rate throughout its lifetime, the blastwave from the SN explosion will therefore propagate through a circumstellar medium with p oc r~2 for roughly 5,000 years. This simple picture is not always appropriate: The progenitor stellar wind may be time-dependent, asymmetric, or both. The resulting circumstellar 139

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medium (CSM) would then be non-spherical, and/or would not follow the p oc r~2 radial profile. The clearest evidence against a steady, spherical wind is the CSM around SN 1987A. Optical observations of SN 1987A have revealed a narrow ring surrounding the SN (Jakobsen et al. 1991), and slightly more extended lobes above and below this ring (Wampler et al. 1990). The inner ring has a radius of ~ 6.3 x 10 17 cm, a. radial velocity (either expansion or contraction) of 10.3 km s" 1 (Crotts & Heathcote 1991), and an inferred density of approximately 2 x 10~ 4 cm~ 3 (Lundqvist & Fransson 1991). The radius and low velocity of the ring implies a change in wind (and hence stellar) properties only some 104 years before the SN explosion. The nonspherical shape of the ring implies some sort of asymmetry in the stellar wind from the progenitor star in at least one phase of it's late evolution. In fact, a circumstellar shell at roughly the radius of the ring was anticipated by Chevalier and Fransson (1987). This conclusion was based on two observations. First, the pre-SN star was a blue supergiant (BSG) rather than a RSG as expected. Second, optical observations of the light echo revealed a circumstellar shell at roughly 3-4 pc (Crotts & Kunkel 1991). If this shell were interpreted as the edge of a stellar wind bubble created by a RSG, then one would be lead to the conclusion that the progenitor evolved from a RSG into a BSG before it's fateful explosion as a SN. This evolutionary scenario is consistent with at least some theoretical models for the evolution of the progenitor of SN 1987A (e.g., Arnett 1991). In this scenario, the fast wind of the BSG must have overtaken the slow wind of the RSG, sweeping it up into a shell. The surprise is not that there is a. circumstellar shell at a distance of ~ 10 18 cm, but that it is not spherical.

2 Interacting Winds Model Models for the formation of aspherical circumstellar structures (in planetary nebulae) in fact predate SN 1987A (Kwok 1982; Kahn 1982). In this "interacting winds" model, a fast, low density wind from a proto-white dwarf blows into the relic slow wind emitted from the previous asymptotic giant branch (AGB) stage of evolution. Because the later wind is faster than the previous AGB wind, it overruns and sweeps up the gas of the AGB wind forming a dense, expanding shell. If the AGB wind was not spherically symmetric, that asymmetry would be reflected in the asymmetry of the swept up shell. This interacting winds model was applied to the circumstellar shell of SN 1987A by Luo & McCray (1991) and Wang k Ma.zza.li (1992), and later by Blondin & Lundqvist (1993). In this case the shell was formed as the

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fast wind of the BSG progenitor overtook the slow wind of the RSG phase. In order to calculate the evolution of the shell, Wang & Mazzali and Luo & McCray used two assumptions: that the shell of swept up RSG wind was infinitely thin, and that the shocked BSG wind was isobaric. The first assumption is valid as long as the shocked RSG wind has time to radiatively cool, collapsing to high density and very narrow width. The second assumption is valid in the limit that the sound speed in the shocked BSG wind is much larger than the expansion velocity of the bubble, cs ^ vexp. Luo & McCray (1991) pointed out that under this assumption any asymmetry in the BSG wind would be wiped out by subsonic motions in the shocked BSG wind, and would not affect the expansion of the pressure-driven shell. Hence the asymmetry of the shell must relate to an asymmetry in the RSG wind. While these models produced a nice explanation for the ring around SN 1987A, there remained a few problems. One of these dilemmas was the inconsistent conclusions: Wang & Mazzali (1992) claimed a. very small asymmetry was needed, while Luo & McCray (1991) employed a wind asymmetry of 5. Second, both of these models predated the reported velocity mapping of the shell by Crotts & Heathcote (1991), which found a very small expansion of the ring of only 10.3 km s" 1 . And third, the assumption of an isobaric interior turns out not to be a valid approximation for the parameters used in their models. This third point was illustrated by Blondin & Lundqvist using 2D hydrodynamical simulations. For the standard parameters used by previous authors, they estimated a ratio of cs/vexp ss 4. Although this ratio exceeds unity, it does not do so by much. Hydrodynamic simulations using these standard parameters showed that the interior pressure varied by more than two orders of magnitude, with a high pressure region formed on the equator at the point where the previous authors found the formation of a cusp. On the other hand, by the time the velocity ratio cs/vfrp reaches a value of 10, the interior pressure is relatively uniform. Note that a direct result of dropping the isobaric approximation is that an asymmetry in the BSG wind could give rise to an asymmetric shell. The principle conclusions of the hydrodynamical models of Blondin & Lundqvist (1993) were that a strong cusp was not formed as seen in the thin shell models, and that a very strong density asymmetry in the RSG wind was required to produce the observed ring/lobe structure around SN 1987A. A reasonably good model was constructed in this work by assuming some rather extreme parameters: a very slow, dense RSG wind with v = 5 km s" 1 and M = 2 x lO~ 5 M 0 yr~ 1 , a density in the equatorial wind of the RSG some 20 times higher than in the polar wind, and a very steep density

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Fig. 1. A hydrodynamic model for the circumstellar shell surrounding SN 1987A (from Blondin & Lundqvist 1993). The density contours are spaced logarithmically. The highest density gas is shaded in black to highlight the location of the equatorial ring. profile in the RSG wind such that roughly 50% of the RSG wind mass was collimated within 10° from the equatorial plane. The circumstellar bubble resulting from the interaction of a spherical BSG wind with this aspherical RSG wind is shown in Figure 1.

3 Wind Asymmetry The hydrodynamical modeling of the ring around SN 1987A suggests that the interacting winds model can produce the observed structures if there is a significant asymmetry in the progenitor stellar wind. Why might the stellar wind have been asymmetric? At least three possibilities come to mind; rotation of a single star, interaction in a close binary system, or dynamically significant stellar magnetic fields. Due to lack of space I will

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not review the possible role of magnetic fields, but simply refer the reader to the contributions in these proceedings by Livio and by Chevalier & Luo.

3.1 Stellar Rotation Let us consider the simplest case first: rapid rotation of a single star. Eriguchi et al. (1992) pointed out that although the radius of the final BSG progenitor was larger than the radius of the original main sequence star, the moment of inertia of the BSG was smaller than that of the main sequence progenitor due to the increased mass concentration toward the center of the star. Taking into account mass and angular momentum loss at each evolutionary stage, their work suggests that a main sequence star rotating at 15% of its critical rotation rate would end up as a BSG rotating at 35% of its critical rotation rate. This may be sufficient to produce a significant asymmetry in the BSG wind. Bjorkman & Cassinelli (1993) have recently presented a simple but powerful explanation for Be stars, of how a radiatively-driven wind from a rapidly rotating star will be strongly focussed into the equatorial plane. This effect is particularly strong in B-type stars where the terminal wind velocity is only of order the escape velocity from the surface of the star, as is likely for the immediate progenitor to SN 1987A. Although this model has yet to be directly applied to the progenitor of SN 1987A, 2D hydrodynamic simulations of a B2.5 star show that a rotation rate of 40% the critical rotation produces a density contrast from equator to pole of 2.5 (Owocki, Cranmer, & Blondin 1994). This density contrast rises sharply as the rotation rate nears the critical rotation rate. This theory of rotating winds may also apply to late-type stars as suggested by Livio (these proceedings). While the direct application of this theory to RSG winds must await a better understanding of such winds, it is at least interesting to speculate that this theory may suggest the existence of strongly asymmetric winds from rotating RSGs. Given that the winds from late-type stars are relatively slow compared to the escape velocity from the stellar surface, one might even speculate that this dynamical effect will be more pronounced in evolved stars than in early-type stars. Any confirmation of such speculation must await a theory for the origin of winds from late-type stars.

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Binary

There are a variety of possible ways in which the interaction of a binary companion could have affected the CSM of the supernova progenitor star. To simplify matters, we can separate these processes into two categories, interacting binaries and wide binaries, where here we use the term wide binary to refer to systems in which the binary separation is sufficiently wide, and the secondary mass sufficiently small that no mass is transferred between the stars and they remain two separate stars throughout the evolution. In the former category of interacting stars there are several processes that might lead to an asymmetric CSM shortly before the SN explosion. In the absence of an observed massive companion, these scenarios all come down to one common final stage in the binary evolution: a common envelope (CE). During the CE evolution angular momentum is transferred from the binary orbit to the rotation of the stellar envelope. The final outcome of the CE evolution depends critically on how much energy is transferred from the binary orbit into the CE. If sufficient energy is transferred to the CE that the envelope will be ejected from the system, the binary will end up with the companion star in a compact orbit around a BSG. The BSG will presumably have gained angular momentum during this process, but may not be in synchronous rotation with the binary orbit. If this is the case, tidal torques will continue to increase the rotation rate of the BSG star. In the absence of sufficient energy transfer during the CE evolution the binary companion will merge with the core of the SN progenitor while the CE is still partially intact. The final state of the system would then be a rapidly rotating RSG. Subsequent evolution must then produce a. BSG star, as constrained by observations. Conservation of angular momentum then implies that this BSG progenitor would be rotating near the critical rotation rate. Thus any CE evolution is likely to end with a rapidly-rotating BSG. If the initial binary separation of the progenitor system is larger than ~ 7 X 1013 cm the system does not evolve into a common envelope phase (Chevalier & Soker 1989) and we are left with our latter category of wide binary systems. In this case we would expect the companion star to still be in the vicinity of the supernova. Current observational constraints limit the mass of this companion to under 2.5 MQ (Crotts & Kunkel 1993). In this case there are at least three possible sources of wind asymmetry: tidal spin up of the progenitor star, tidally enhanced mass loss, and gravitational focussing of the stellar wind. Chevalier & Soker (1989) point out that for a. relatively narrow range of initial separations, the system will not enter a. CE phase, but will be close

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enough to spin up the progenitor star during its RSG phase. This would give the RSG a value of Q/Q.cr ~ 0.5. This may or may not produce an asymmetric RSG wind, but it would certainly lead to a rapidly rotating BSG, which we have seen above implies an asymmetric BSG wind. It has also been suggested that a companion in a wide orbit could produce enough tidal distortion to generate an enhanced wind in the equatorial plane (Tout & Eggelton 1988). This effect would likely occur only for the same small range in binary separation that tidal spin up would be important. As in the case for rotationally-enhanced mass loss in RSGs, a quantitative estimate of tidally-enhanced mass loss must await a better understanding of the origin of winds from RSGs. If a low-mass companion is in too wide an orbit to tidally affect the progenitor star, it can still affect the RSG wind through its gravitational field. Because the RSG wind is moving relatively slowly, even a. small companion can gravitationally focus the wind, concentrating the mass of the wind into the equatorial plane.

4 Ring Formation Given these different possible sources of wind asymmetry, let us consider how they might give rise to the observed ring. Most of the sources of asymmetry discussed in the previous section rely either directly or indirectly on stellar rotation. In fact, this is the only suggested source of asymmetry (other than magnetic fields) for a single star progenitor. For the binary models, rotation is almost inevitable. If rotation of the RSG progenitor is able to focus the wind into the equatorial plane, perhaps by the Bjorkman-Cassinelli mechanism as in early-type stars, then the hydrodynamical simulations of Blondin & Lundqvist (1993) suggest that a compact ring can be formed in the interacting winds model. However, conservation of angular momentum implies that if rotation is important in the RSG wind, it is very important in the BSG wind. The interacting winds model must then include both an asymmetric RSG wind and an asymmetric BSG wind. If the shell expansion is sufficiently slow that the isobaric approximation is valid, then the asymmetric BSG wind becomes an unimportant feature of the model. If the isobaric approximation does not hold, then it becomes possible, if not likely, that an asymmetric BSG wind (due to rapid rotation) will lead to an asymmetric shell. Whether the asymmetry in such a shell matches the morphology of the ring around SN 1987A must await the results of further hydrodynamic simulations. We discussed a few scenarios for asymmetric CSM in §3 that did not in-

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volve a wind from a rapidly rotating star. These included mass ejection in a CE phase, tidally enhanced winds, and gravitationally focussed winds. Although the mass ejection in the CE phase has been shown to be strongly confined to the equatorial plane (Livio & Soker 1988), the typical outflow velocity is of order the orbital velocity of the merging companion. This is much higher than the 10.3 km s"1 expansion of the ring, and thus the identification of the ring material with the outflow from a CE stage must be considered tentative. One might envision that the mass outflow at a later stage in the CE evolution might not be so fast, but yet still be asymmetric. In the limit that this mass outflow is really just a stellar wind from a rapidly rotating RSG, we are back to the rotation models discussed previously. More work must be done on the dynamics of mass loss from CE systems to verify the validity of this class of models. In any case, the end result of a CE evolution must be a rapidly rotating star, which again leads us to the conclusion that the immediate BSG progenitor to SN 1987A had an asymmetric stellar wind. A tidally-enhanced RSG wind would presumably produce a mass concentration in the equatorial plane as assumed in the hydrodynamical models of Blondin h Lundqvist (1993). But again, we do not have an accurate description of such a model. In particular, would such a wind have a faster or slower velocity in the equatorial plane (a very slow velocity being required for the ring around SN 1987A)? Would the binary be close enough to also spin up the RSG, leading eventually to a rapidly-rotating BSG progenitor? These questions could be answered with hydrodynamic simulations if we had a good model for the origin of winds in RSG stars. Gravitational focussing appears to be the only source of asymmetry that is not accompanied by the ultimate conclusion that the BSG progenitor was rapidly rotating. However, it is not yet clear that the strong density asymmetry required in the models of Blondin k. Lundqvist (1993) could be produced by a low-mass companion. 5 Conclusions Although the existence of a circumstellar shell around SN 1987A is in fact expected based on the assumed evolutionary history of the progenitor, it is much harder to explain the severe asymmetry of the observed ring. The most likely source of asymmetry in the progenitor wind is rapid rotation in either the RSG or BSG phase of evolution. While a rapidly rotating BSG may occur in a single star system, a rapidly rotating RSG necessitates a binary system. In all binary scenarios (except perhaps a CE evolution in

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which all the orbital angular momentum went into the ejected envelope), the final BSG star must have been rotating rapidly, and hence must have had an asymmetric wind in accordance with the Bjorkman-Cassinelli model. There are, however, other scenarios that could work, but have not been sufficiently explored. These include a tidally enhanced wind from the RSG, a gravitationally focussed RSG wind due to a solar mass companion, or perhaps a gentle form of envelope ejection in a CE evolution. Finally, a few words on future observations. Unfortunately, we will not be able to observe a possible companion star as small as a solar mass (Crotts, private communication), so we may never be able to rule a binary system out. We can, however, expect to get more information on the circumstellar structure when the young supernova remnant, begins to interact with the ring and lobes, and material exterior to the ring. In particular, any sort of asymmetry in the RSG wind will become readily apparent. Also, high resolution observations of the SNR may be able to determine the extent of any asymmetry in the progenitor BSG wind. I would like to thank many of the Herstmonceux conference participants, including Roger Chevalier, Vincent Icke, Mario Livio, and Noam Soker, and especially my collaborator in this field, Peter Lundqvist, for enlightening discussions on this subject. References Arnett, W. D. (1991), Astrophys. J., 383, 295. Bjorkman, J. E., & Cassinelli, J. P. (1993), Astrophys. J., 409, 429. Blondin, J. M., k Lundqvist, P. (1993). Astrophys. J., 405. 337. Chevalier, R. A., k Fransson, C. (1987), Nature, 328, 44. Chevalier, R. A., &; Soker, N. (1989), Astrophys. J., 341, 867. Crotts, A. P. S., k Heathcote, S. R. (1991). Nature, 350, 683. Crotts, A. P. S., k Kunkel, W. E. (1991), Astrophys. ./., 3G6. L73. Crotts, A. P. S., k Kunkel, W. E. (1993), IAU Circular, 5691. Eriguchi, Y., Yamaoka, H., Nomoto, K., k Hashimoto, M. (1992), Astrophys. J., 392, 243. Jakobsen, P., et al. (1991), Astrophys. J., 3C9, L63. Kahn, F. D. (1982), in IAU Symp. 103, Planetary Nebulae, ed. D. R. Flower (Dordrecht: Reidel), p. 305. Kwok, S. (1982), Astrophys. J., 258, 280. Livio, M., k Soker, N. (1988), Astrophys. J., 329, 764. Lundqvist, P., k Fransson, C. (1991), Astrophys. J., 380, 575. Luo, D., k McCray, R. (1991). Astrophys. J., 379, 659. Owocki, S. P., Cranmer, S. R., k Blondin, J. M. (1994). Astrophys. J., in press. Tout, C. A., k Eggelton, P. P. 1988, M.N.R.A.S., 231, 823. Wampler, E. J., Wang, L., Baade, D., Banse, K., D'Odorico, S., GouifFes, C , k Tarenghi, M. (1990). Astrophys. J., 362, L13. Wang, L., & Mazzali, P. A. (1992). Nature, 355, 58.

Supernovae with dense circumstellar winds N. N. Chugai Institute of Astronomy, Russian Academy of Sciences Pyatnitskaya 48, 109017 Moscow, Russia

Abstract The circumstellar (CS) wind around a type II supernova (SN II) can be revealed through the optical emission induced by the collision of SN ejecta with the wind. The optical manifestations of the ejecta-wind interaction provide an excellent tool for the study of the mass-loss history of pre-SN II at the final red supergiant stage. There is strong evidence that pre-SN II with an extraordinarily high mass-loss rate, M > 10~ 4 MQ yr" 1 , originate from the low-mass end of the massive star range (Mm, ~ 8 - IOMQ), while pre-SN II-P originating from Mmi > 12A/C, are characterized by a very low mass-loss rale, M < \()~bMQ yr~'.

SN 1979C (a type II-L), known for its powerful radio emission, was the first SN II where the late-time Ha luminosity was attributed to the ejecta-wind interaction (Chevalier & Fransson 1985). Yet the success of the radioactive model for the late-time luminosity of SN 1987A raised the problem of choosing between radioactive and shock-wave mechanisms in SN II. One possible solution was prompted by the observed excess in the Ha luminosity of SN 1980K (also type II-L and a strong radio emitter) at t — 670 days, relative to the predictions of the radioactive model (Chugai 1988). The interpretation of the excess in terms of the ejecta-wind interaction was supported by the strong radio luminosity and wide flat-top profile of Ha. In a sample of six SN II whose Ha fluxes were available (69L, 70G, 79C, 80K, 87A, 87F), three were found to possess significant excesses in their Ha luminosities (Chugai 1990). As mentioned already, two of these (79C and 80K) were radio emitters, but the third, SN 1987F, had not been detected at radio wavelengths. Nevertheless, the strong Ha luminosity of SN 1987F on day 150, L{Ra) « 1.5 1041 ergs s" 1 , indicates a dense wind with a parameter w = M/uw « 10 17 g cm" 1 , compared to 8 1015 g c m ' 1 for SN 1979C. Remarkably, the two SN II with the strongest excesses in their \la luminosities (79C and 87F) show a broad emission profile in this line with no absorption component. This property is consistent with the emission of the line from a narrow over-excited outer layer of the SN envelope. A similar pure-emission Ha profile is characteristic of the supernovae selected by Schlegel (1990) in the new subclass SN Tin. This property seems to be an indication that the luminosity of all SN Iln is strongly affected by the ejecta-wind interaction. This is particularly convincing in the case of the 148

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SN Iln SN 1988Z. Strong, variable, narrow, nebular lines in its spectrum are indicative of a dense CS wind (Filippenko 1991), while the broad Ha emission, lacking an absorption component, is consistent with the ejecta-wind interaction mechanism (Chugai 1990). This suggestion is confirmed both by the tremendously high Ha luminosity compared to that of SN 1987A and by the pronounced excess over the radioactive model (Turatto et al. 1993). In addition to attaining similar maximum Ha luminosities, SN 1987F and SN 1988Z exhibited similar light curves, characterized by an unusually slow decline in the B, V, R magnitudes. This resulted in extraordinarily high bolometric luminosities at the stage t > 100 days (Filippenko 1989; Stathakis and Sadler 1991; Turatto et al. 1993). This peculiar behaviour in their luminosities places these supernovae out. of the standard photometric subclasses II-P and II-L. As indicated above, the unusually strong excess in the bolometric luminosities of SN 1987F and SN 1988Z can be attributed to the energy released during the collision of the SN ejecta. with a dense wind. The required wind density parameter is w ss 101' g cm" 1 (Chugai 1992). This estimate may be reduced to w ss 210 1 6 g cm" 1 if the efficiency of the transformation of the kinetic energy into radiation is close to 100%. Nevertheless, even in this case, the CS winds around SN 1987F and SN 1988Z would be the densest ever detected around SN II. To emphasise the unusually strong effects of the ejecta-wind interaction on the Ha and broad-band emission, we tentatively designate such supernovae SN IIsw, with "sw" standing for "superwind" or "shock wave". Their key properties are (i) a strong Ha luminosity > 1041 erg s" 1 ; (ii) a, slow decay in broad-band optical luminosity relative to normal SN II; (iii) the lack of an Ha absorption component. Whether or not all SN Iln are identical to SN IIsw is unclear, since properties (i) and (ii) have been established only for SN 1987F and SN 1988Z. The details of the formation of the optical spectrum of SN IIsw are poorly understood. In the case of a sufficiently dense CS wind, w > 10 16 g cm" 1 , one expects that the reverse shock wave propagating into the ejecta. is radiative for approximately one year after the explosion, and that a cool dense shell forms (Chevalier & Fransson 1985). The cool dense shell, irradiated by X-rays, could be responsible for the bulk of the optical radiation in SN IIsw. However this shell is liable to fragmentation clue to Ray leigh-Taylor instability (Chevalier & Fransson 1985). If certain conditions are met (relatively large geometrical thickness compared to Sobolev length; large covering factor) the line profile, even from a clumpy shell, can be formally described in the Sobolev approximation. For SN 1987F on day 150 the Ha profile can

150

N. N. Chugai: Supernovae with dense circumstellar winds

be reproduced by an optically thick outer shell with zero velocity gradient attached to the partially opaque inner ejecta with rp ss 1 (Chugai 1991). However, this simple model fails in the case of SN 1988Z . In addition to the broad component of Ha (FWHM « 104 km s" 1 ) and the unresolved line from the photoionized wind, a third, intermediate component with FWHM « (1 - 2) x 10 3 km s" 1 is seen (Filippenko 1991; Turatto et al. 1993). A promising model for the origin of this spectrum invokes a clumpy wind, with the intermediate component originating from a radiative shock wave being driven into the dense clumps (Chugai & Danziger 1994). In a similar model, the interaction of the SN ejecta with a clumpy wind may account for the optical and X-ray emission from the strong radio/X-ray supernovae SN 1986J and SN 1978K (cf. Chugai 1993). The proposed dumpiness of the pre-SN II wind is supported by the increasing evidence for clumpy structure in RSG winds, revealed by molecular radio line observations (cf. Olofsson, this volume). SN II-P constitute roughly two thirds of all SN II but, so far, have never shown any of the signatures of a CS wind. An upper limit for the wind density around SN II-P can be obtained from the maximum velocity, vmax, determined from the blue wing of the Ha absorption. This velocity is obviously lower than the velocity of the ejecta at the reverse shock wave, vmeLX < Rc/t, where Rc is the radius of the contact discontinuity, and t is the expansion time. To estimate the maximum wind density, a constant density in the inner part of the SN II, p oc v~k in the outer part, and a wind density profile p = w(4nR2)~1 is adopted. The self-similar solution for Rc (cf. Chevalier 1982) is used, and M = 12M© for the ejecta mass, 1051 ergs for the kinetic energy and k = 8 are assumed. From observation, vmax = 11900 km s" 1 in SN 1969L on day 59 (Benetti 1991) and vmaK = 13000 km s" 1 in SN 1990E on day 32 (Schmidt et al. 1994). The resulting upper limit is to < 1015 g cm" 1 for both SN II. This limit decreases if k or M increases. The derived upper limit shows that the wind density around SN II-P is relatively low compared with that of SN 1980K (type SN II-L) where i » « 3 1015 g cm" 1 (Lundquist & Fransson 1988). The origin of the dramatic difference in the mass-loss rates among preSN II, ranging from M < 10~5 M© yr" 1 for SN II-P to ~ 5 10" 4 M Q yr" 1 for SN IIsw, is unclear. Basically, two factors may affect the mass-loss rates of pre-SN II viz. main-sequence mass, and the separation of the components if the SN II progenitor is in a binary system. With some suitable initial values for the main-sequence mass and binary separation, the common envelope regime at the RSG stage might stimulate a robust mass-loss at just the appropriate epoch viz. ~ (several)xlO 4 yrs prior to the explosion. This

N. N. Chugai: Supernovae with dense circumstellar winds

151

possibility, while plausible, has not yet been given detailed consideration (A.V.Tutukov, private communication). Alternatively, the range of main-sequence masses alone could be responsible for the observed variation in mass-loss rates among pre-SN II. Although standard mass-loss rates for massive stars (de Jager & Nieuwenhuijzen 1988) are insufficient to account for the rates inferred in pre-SN IIsw, nevertheless some facts favour this alternative. Firstly, the low ejecta mass of SN 1988Z indicates a progenitor main-sequence mass in the range 8 — 10 MQ (Chugai & Danziger 1994). Secondly, SN II-P always exhibit a strong saturated [01] 6300, 6364 A doublet, suggesting that supernovae of this subclass originate from stars which produce large amounts of oxygen. Such stars have masses M m s > 12 MQ (cf. Woosley & Weaver 1986). Thirdly, the masses of SN II-L seem to be generally lower than those of SN II-P (cf. Chugai 1990). All these facts may be reconciled with the data on the pre-SN II winds if the main-sequence mass increases along the sequence SN IIsw, SN II-L, SN II-P, while the mass-loss rate at the final R.SG stage decreases with main-sequence mass in the range 10-12 MQ. This scheme suggests that preSN II-P (Mms > 12 MQ) lose their mass according to the de Jager law, while the mass-loss rates of SN II-L (Mms « 1 0 - 1 2 MQ) and SN IIsw (Mms ss 8 — 10 MQ) are substantially higher. It is tempting to identify the tremendously high mass-loss rates of preSN IIsw with a superwind phenomenon in stars M m s « 8 — 10 MQ. This superwind might have an origin similar to that seen in WD producers (Mms < 8 MQ). However, given the relatively low rate of SN IIsw, only 10-20% of all the stars in the range 8-10 MQ should result in such supernovae. Other supernovae arising from this mass range must be unobserved. This could happen if e.g. (i) the hydrogen envelope is completely lost before the supernova explosion, so that the pre-SN becomes a blue helium star, and (ii) the amount of ejected 56 Ni is very low (< O.OIMQ). At explosion, such a supernova would be faint with a narrow (5-10 days) maximum of absolute magnitude only « —14. The hydrogen lost by the progenitor would lie not far from the supernova, at a distance R = vwtss ~ 10 17 cm, where ^BS * 10 3 ~ 4 yr is the duration of the final blue stage and vw w 10 km s" 1 is the RSG wind velocity. Such faint SN II may be recoverable in optical, radio and X-ray observations after several years, due to the collision of the supernova ejecta with the dense wind (e.g. SN 1978K and SN 1986J?).

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References Benetti, S. (1991). SN 1987A and other supernovae, eds. I.J. Danziger & K.Kjar (ESO, Garching), p. 339. Chevalier, R.A. (1982). Astrophys. J., 258, 790. Chevalier, R.A. & Fransson, C. (1985). Supernovae as Distance Indicators, ed. N. Bartel (Springer, Berlin), p. 123. Chugai, N.N. (1988). Astrophys. Space Sci., 146, 375. Chugai, N.N. (1990). Sov. Astr. Lett., 16, 457. Chugai, N.N. (1991). Mon. Not. R. astr. Soc, 250, 513. Chugai, N.N. (1992). Sov. Astr., 36, 63. Chugai, N.N. &; Danziger, I.J. (1994). Mon. Not. R. astr. Soc.,, in press. Chugai, N.N. (1993). Astrophys. J. Lett., 414, L101. Filippenko, V.A. (1989). Astron. J., 97, 726. Filippenko, V.A. (1991). Supernovae, ed. S.E. Woosley (Springer, New York), p. 467. de Jager, C. & Nieuwenhuijsen, R. (1988). Atmospheric Diagnostics of Stellar Evolution, ed. K.Nomoto (Springer, Berlin), p. 122. Lundqvist, P. &; Fransson, C. (1988). Astron. Astrophys., 192, 221. Schlegel, E.M. (1990). Mon. Not. R. astr. Soc, 244, 269. Schmidt, B.P., Kirshner, R.P., Schild, R. et al. (1994). Astrophys. J., (in press). Sramek, R.A., Weiler, K.W. & Panagia, N. (1990). IA U Circular No. 5112. Stathakis, R.A. & Sadler, E.M. (1991). Mon. Not. R. astr. Soc, 250, 786. Turatto, M., Cappellaro, E., Danziger, I.J., Benetti, S., Goiiiffes, C. & Delia Valle, M. (1993). Mon. Not. R. astr. Soc, 262, 128. Woosley, S.E. & Weaver, T.A. (1986). Ann. Rev. Astron. Astrophys., 24, 205.

Compact Supernova Remnants Roberto J. Terlevich Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, U.K.

Abstract Two new kinds of peculiar type II supernovae (SNe) have been observed recently: namely the very luminous type II radio supernovae (RSNe) and the so-called Seyfert 1 imposter. I will show that a simple model of interaction of supernova (SN) ejecta with a high-density homogeneous circumstellar medium (CSM), combining analytic and numerical hydrodynamic simulations together with static photoionization computations, can describe their observed emitted spectrum, optical light curve, X-ray luminosity and emission line widths. I suggest that these two new kinds of SNe are not peculiar type Us, but are, in fact, the optical or radio manifestation of the same phenomenon, i.e. the interaction of the SN ejecta with a high density CSM. During the interaction with a high density CSM a young remnant can radiate most of its kinetic energy and outshine the SN event itself; therefore to emphasize the unique aspects associated with this type of event, I suggest calling this group of small, luminous and rapidly evolving remnants, compact supernova remnants (cSNRs).

1 Introduction The defining characteristic of type II SNe is the presence of very broad Ha emission with a strong P-Cygni profile. A small number of peculiar type II SNe have been found in recent years which are either very bright in the optical continuum with very strong and broad Ha emission without a P-Cygni profile, or are strong radio sources then called radio supernovae. Probably significant is the fact that these SNe tend to be associated with regions of active star formation. The optical light curve of peculiar type II SN, such as SN 1987F (also know as Seyfert 1 imposters; see Filippenko 1989) and SN 1988Z, is characterized by a high luminosity maximum followed by an extremely slow decay. The decay is so slow that the broad Ha emission can be observed for several years after discovery. SN 1988Z is also very bright at radio frequencies. In the past 10 years, since the discovery of radio emission in SN 1979C in M100 in the Virgo cluster, several luminous RSNe have been discovered and studied in detail (Weiler et al. 1989, 1990). Some appear to be very luminous at optical wavelengths and with optical spectra dominated by relatively broad Balmer lines showing no P-Cygni profiles. The most luminous examples so far are SN 1986.1 and SN 1988Z. The radio spectrum of RSNe 153

154

R. J. Terlevich: Compact supernova remnants

is inverted at low frequencies at early times, with a turnover above a critical frequency that decreases with time. This behaviour is interpreted in terms of time-dependent free-free absorption by circumstellar material (Weiler et al. 1989). Both the high Ha luminosity with no P-Cygni profiles, and the strong radio emission can be interpreted as resulting from an interaction between the expanding SN ejecta and a homogeneous, dense circumstellar shell created by the interaction of the slow wind from the progenitor star with the high ambient pressure of the surrounding HIT region (Chevalier 1982; Terlevich et al. 1992; Chugai & Danziger 1994). Supernova remnants evolving in a dense and homogeneous CSM (n > 105 cm" 3 ) reach their maximum luminosity (L > 107 L@) at small radii (R < 0.1 pc ), soon after the SN explosion (t < 20 yr) and while still expanding at velocities of more than 1000 kms~1(Shull 1980; Wheeler et al. 1980; Draine and Woods 1991; Terlevich et al. 1992). In these compact SNR.S, radiative cooling becomes important well before the thermalization of the ejecta is complete, making the remnant miss the Sedov track. As a result, the shocked matter undergoes a rapid condensation behind both the leading and the reverse shocks. Two concentric, high-density, fast-moving thin shells are then formed. The cool, dense shells, the freely expanding ejecta, and a section of the still dynamically unperturbed interstellar gas, are all irradiated and ionized by the photon field produced by the radiative shocks. This contribution describes an attempt to explain the most luminous optical and radio SN events observed in the past few years. 2 cSNRs, or SNRs evolving in a high density medium The interaction of the SN ejecta with a dense CSM causes a shocked region of hot gas enclosed by two shock waves: on the outside the leading shock, and on the inside the inward facing "reverse" shock. The leading shock (V ~ 104 kms" 1 ) encounters dense circumstellar material and raises its temperature to ~ 109 K. The reverse shock, which is initially substantially slower (V ~ 103 kms" 1 ), begins to thermalize the supernova ejecta to temperatures of about 107K. Early analytical and numerical computations of the evolution of SNRs in a dense medium (Chevalier 1974; Shull 1980; Wheeler et al. 1980) showed a speeded-up evolution compared with the "standard" solution in a medium of no — 1 cm"3 . All evolutionary phases (free expansion, thermalization of the ejecta, the quasi-adiabatic Sedov phase, the radiative and the pressure modified snow-plough phases)

R. J. Terlevich: Compact supernova remnants

155

which have been thoroughly studied for the standard case, are substantially speeded up. For SNe evolving in circumstellar densities of the order of no ~ 107 cm" 3 , the cooling-time (tc) and cooling-length (r c ) scales for the post-shock temperatures are very small. f

f. ^ n 9

rv

LQ —

°

w»i^^^—^^—

1

14

vr

(A\

y i«

i j. i

U3

r c ~-* c t, s ~1.8xlO —f- cm, 4

(2)

TJ7A23

where v$ = u/10 8 km s" 1 and A23 = A/10" 2 3 erg cm 3 s""1. Thus, radiative losses become important at very early times when the shock velocities and temperatures are vs > 103 km s" 1 and Ts > 107 K, well before the ejecta is even thermalized. This means that a large flux of ionizing photons will emerge from the shocked gas at X-ray energies. For a supernova remnant which injects 1051 ergs into a medium of constant density n-j = n o / 1 0 7 c m - 3 the onset of the radiative phase behind the leading shock (assuming f-f cooling only) causing the formation of a dense outer shell, (Shull 1980, Wheeler et al. 1980, Draine and Woods 1991), begins at a time tsg given by tsg = 230 E\l%n~3lidays

(3)

where £51 is the energy deposited by the SN in units of 10 51 ergs. At this stage, the shock is at a radius of Rshock = 0.01 £ 5 ( iij '

(-—J

pc

(4)

with velocity,

Vshock = 4600 EH8 A1' (j-) ' ^ km s"1

(5)

temperature, tsg

and luminosity, Lshock = 2 x 10 43 E75l8 n 3 / 4 ( / - ) ~ n / < ergs" 1 .

(7)

156

R. J. Terlevich: Compact supernova remnants

These approximate formulae assume that the ejecta has already been fully thermalized. However, for the case of interest in this work (i.e. for no > 105 cm" 3 ) strong radiative losses occur before thermalization is completed. Because the cooling processes radiate the thermal energy at the same rate as thermalization proceeds, the Sedov phase is totally inhibited and thus there is no self-consistent analytic treatment for the evolution of such remnants. It is therefore necessary to follow the detailed time evolution of the gas flow. In particular, special care should be given to the post-shock structure which is sensitive to the details of the ambient density distribution and to the temperature dependence of the radiative cooling function. The cooling time scale, tc, for an optically thin plasma, is proportional to the inverse of the gas density and the evolution proceeds faster at higher ambient densities. On the other hand, the temperature dependence of the cooling function is different for different temperature ranges. Adiabatic shocks are stable but cooling instabilities can develop over a wide range of radiative shock conditions (Avedisova 1974; Falle 1975, 1981; McCray, Stein and Kafatos 1975; Chevalier and Imamura 1982; Imamura 1985; Bertshinger 1986). The details of the transition from a nearly adiabatic to a strongly radiative shock depend on the ability of the gas to readjust to the cooling rate. Pressure gradients tend to be smoothed out in a. sound-crossing time, td, and the ratio tc/td provides an estimate of the conditions prevailing in the cooling gas. For tc/td > 1, at moderate cooling rates, the gas elements are continuously compressed as their temperature falls and the cooling process operates quasi-isobarically at the pressure attained by the gas immediately behind the shock. For i c /i^ < 1, however, the cooling rate dominates over any pressure readjustment and the process becomes quasi-isochoric at the post-shocked density of the cooling gas elements. A large pressure imbalance then develops in the flow, and new additional shocks are generated which end up compressing the cooled gas. This process, termed "catastrophic cooling" (Falle 1975, 1981), appears during thin shell formation and the instabilities continue to operate during the rest of the radiative shock evolution (Chevalier and Imamura 1982; Bertshinger 1986; Cioffi et al 1988; Tenorio-Tagle et al 1990). The catastrophic cooling acquires a central role in the case of supernovae evolving in high density media due to the strength of the radiation produced upon cooling, and the rapid variations inherent in the shock propagation. These features imply that a large flux of ionizing photons will emerge from the shocked gas. The wide range of gas temperatures in the cooling region results in a power-law-like spectrum at UV and X-ray frequencies.

R. J. Terlevich: Compact supernova remnants

157

3 Numerical hydrodynamic models I present here a short summary of the calculations that follow the evolution of a SNR caused by the release of EQ = 1051 ergs into a constant density medium of n0 = 107 cm" 3 (Terlevich et al. 1992; Tenorio-Tagle et al. 1994). The results of numerical simulations in one and two dimensions indicate that the onset of radiative cooling starts before thermalization is completed for values of no > 105 cm" 3 . Given the large densities and thus the strong radiative cooling, the combination of values of thermal (Eth ~ 0.7 X Eo) and kinetic (E^n ~ 0.3 x EQ) energies which characterizes the Sedov solution is never achieved. Instead, the thermal energy content is radiated away within the first 5 years of evolution, before the thermalization of the ejecta is completed. The calculations also show how both energies continuously drop after strong cooling occurs, and values of Eth much smaller than 0.1 x £o are soon achieved. Such low values of Eth occur in the standard case some 106yr after the explosion. The present solution thus indicates a speeded up evolution and, at t = 10 years, most of the injected energy has been radiated away. The remaining kinetic energy, a. major fraction of which is stored in the remnant outer shell, also decays rapidly. Clearly, given the similar rate of change of both Eth a n d Ekin towards the end of the calculated evolution, the kinetic energy is radiated as soon as it is thermalized behind the shock waves. Strong radiative cooling leads to the development of a thin-shell at the edge of the remnant as well as the collapse of the shocked ejecta behind the reverse shock. The 2-D calculations, which also allow for the development of cooling instabilities, clearly show, as in 1-D, the steady approach of the two shells of cool matter, promoted by the weakening and withdrawal of the reverse shock while the unshocked ejecta fills the remnant interior. The various numerical solutions, despite different resolutions and initial conditions, all show a small (Rsnr ~ a few times 1016cm) rapidly radiating remnant, with two massive and geometrically thin concentric shells moving at large speeds and bounded by radiative shock waves. A small fraction of the swept-up mass lies between the two shocks and remains at very high temperatures (about 108K). On the other hand, in the central regions of the remnant, the ejected matter continues to expand homologously, and thus presents a density and velocity distribution which reflect the initial conditions, until it is thermalized at the reverse shock. This central region, together with a section of the still unperturbed background gas and the two thin, cool shells, are clearly subjected to the ionizing radiation provided by the radiative shocks.

158

R. J. Terlevich: Compact supernova remnants Table 1. Shock evolution for no — 10 7 cm~3 Time Ug 1 2

3 4

6 10 25 50

Time years

Log(fl)

0.63 1.26 1.89 2.52 3.78 6.30 15.7 25.6

16.5 16.6 16.65 16.7 16.75 16.8 16.9 17.0

cm

v,h kms *

Log(T) K

4600 3700 2760 2250 1680

8.3 7.8 7.6 7.4 7.2

1170 610 370

6.9 6.3 5.9

Log(L) ergs" 1 43.3 42.8 42.6 42.4 42.1 41.7 41.1 40.6

Log nshell

cm'3 12.0 11.6 11.3 11.1 10.9 10.6 10.0 9.6

The emitted spectrum is similar for both the main shock moving into the interstellar medium and the reverse shock thermalizing the ejecta. The details, viz. the specific temperature structure of the cooling region and the resulting photon spectrum, cannot be resolved with the numerical time dependent calculations and have been computed with analytical solutions (see Terlevich et al. 1992). The wide range of temperatures in the cooling region results in a power-law-like spectrum at UV and X-ray frequencies. In a strict sense the analytical formulae, described in the previous section, can only be applied when tsg > tth, and this occurs only for no < 5.0 X 10 5 cm~ 3 . Comparing our numerical and analytical results we found that at densities n 0 > 5.0 x 10 5 cm~ 3 , the analytical formulae still give a good description of the run of luminosity, shock velocity, size and shock temperature with time. One important difference is that for these densities, thin shell formation and peak luminosity occur at about 2.5tsg. Table 1 shows the results for a. shock with a total energy of 3 x 1051 ergs" 1 , interacting with a medium with initial density nQ = 107 cm" 3 for several units of tsg = 230 days. Table 1 does not include the luminosity of either the hot cavity or the reverse shock. In the homogeneous case the dominant flux is that of the leading shock (Terlevich et al. 1992).

4 Photoionization models Typical values for the density, size, column density, velocity of the cooled regions of gas, and ionizing flux of the radiative shock, were obtained from the analytical and numerical hydrodynamical models described in the previous sections. These parameters were used as input to the photoionization code CLOUDY (Ferland 1990). Table 2 lists the results of the photoionization

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159

Table 2. cSNR line luminosities Time t.g

Log L(Ha ) ergs" 1

1 41.55 2 3 4 6 10 25 50

41.47 41.38 41.29 41.08 40.73 39.81 39.12

Balmer

Log L(H0) ergs" 1 41.04 40.86 40.67 40.48 40.07 39.42 38.71 38.42

Decrement 1.9 2.8 4.2 6.3 13.5 44 89 63

3.3 4.1 5.1 6.5 10.2 20.4 12.6 5.1

models for the same set of parameters listed in Table 1 and solar abundance gas. Values of Ho luminosity in Table 2 for tsg < 2.5 should be treated with caution because before thin shell formation they represent an upper limit to the total Ha luminosity emitted by the cSNR. Values are reliable after thin shell formation, i.e. for t > 2.5 x tsg. General trends can be recognized, i) Both Ha and H/? luminosities decrease steadily with time ii) The Balmer decrement increases reaching a maximum of around 20; only in the early phases is the value of the Balmer decrement close to the Case B recombination value of about 2.85, and iii) Lya is very strong at later times becoming the main coolant in the photoionized region. These results apply only to the broad component with line widths similar to the shock velocities listed in Table 1.

5 Comparison with observations The nearest example of a RSN is probably SN~1955 in the core of the nearby starburst galaxy M82. This RSN, with an age between 20 and 50 yr, has a radio luminosity equivalent to 200 times that of the most luminous galactic SNR, Cass A. This remnant and SN 1986J are the only ones for which VLBI radio images are available. The interpretation of the radio data indicates high CSM densities (~ 107 cm" 3 ) and large progenitor masses (M> 20 M 0 ) for these two RSNe. The most dramatic example of bright RSN was discovered in the distant irregular galaxy Mk297. Mk297A reached a maximum luminosity in 20cm equivalent to 30,000 Cass A.

160

R. J. Terlevick: Compact supernova remnants Table 3. cSNR

candidates

SN name

Galaxy

a 1950. 0

Mrk 297A SN 1988Z SN 1986J SN 1980K SN 1987F SN 1979C SN 19881 41.9+58 SN 1978K

Mrk 297 Zw095-049 NGC 891 NGC 6946 NGC 4615 NGC 4321 (M100)

16 03 10 49 02 18 20 34 12 39 12 20 10 18 09 51 03 17

— NGC 3040 (M82) NGC 1313

01.72 10.62 22.6 26.7 07.7 26.71 17 41.96 38.62

6 1950.0 20 40 43.6 16 15 56.4 42 06 18.9 59 55 56.5 26 20 49 16 04 29.5 35 54 .7 69 54 57.4 -66 33 03.4

Notes: * Coordinates are for galaxy nucleus. SN was 5" E and 1" N.

Other notable examples are the "Seyfert 1 imposters" SN 1987F and SN 1988Z. The enormous Ha and bolometric luminosities of these two objects have been interpreted as resulting from the interaction of the ejecta with a high density (n ~ 107 cm" 3 ) CSM. The case of SN 1988Z is particularly important because strong radio emission has been reported and the CSM density was measured, using the [OUT] narrow lines ratio, to be about 107 cm" 3 . SN 1988Z is therefore particularly interesting in that it shows a link between luminous RSNe and cSNRs. Almost all the listed cSNR candidates are reported to be associated with HII regions or regions of active star formation. 5.1 SN 1988Z At a distance of 133 Mpc (cz - 6670 kms" 1 ; for HQ - 50 kms" 1 Mpc"1) this is one of the most distant examples of a cSNR. SN 1988Z is one of the most extreme cases of bright type II with very slow rate of decay and lack of P-Cygni profiles in the emission lines. The rate of decay is so slow that at an age of 568 days (this is a lower limit to the age) SN 1988Z was more than 3 magnitudes brighter than a typical plateau type II (type IIP) (Stathakis and Sadler 1991). Good spectroscopy has been reported up to day 1149 (Stathakis and Sadler 1991; Turatto et al. 1993). Narrow emission lines from the CSM were detected in spectra, taken on day 73. Electron densities in the range 4 X 106 to 2 X 10' cm" 3 are determined from the ratio of [OIII]5007A to [OIII]4363A for an assumed electron temperature in the region of 20,000-7,000 K. This is the best determination to date of the

161

R. J. Terlevich: Compact supernova remnants

42

i

i

i

1

••

1

1

!

1

1

!

1

•

O' •

•

•

•

•

•

•

«

.

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• •

o...

•

o

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c

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1

200

400

i

600

800 Time

1000

1200

1400

Fig. 1. Observed light curve of SN 1988Z in Ho (Dots) compared with the prediction of the simple homogeneous model with no = 10' cm" 3 (Open circles). Luminosity in ergs" 1 time in days.

pre-shock circumstellar density and coincides with the assumed density of the canonical model in Tables 1 and 2. Figure 1 shows the observed Ha luminosity of SN 1988Z as a. function of time. The data were taken from Stathakis and Sadler (1991) and Turatto et al. (1993). The origin of the time axis is the one adopted by the observers, viz. an explosion epoch of 1988 December 1. The predicted Ha luminosity curve from Table 2 is also shown; again the origin of the time axis corresponds to the SN explosion epoch. The predicted light curve gives a reasonable but not excellent prediction of the observed light curve. Both the value and epoch of the maximum luminosity in Ha seem well predicted, but the decay of the model is too slow compared with the observations. This suggests either a CSM density slightly higher than the value adopted here or a radial gradient in the CSM density. Integration of either the Ha light curve or the blue light curve indicates, after applying a moderate bolometric correction, that the total energy radiated by SN 1988Z in the first 4 years was in excess of 2 X 1O50 ergs. This

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R. J. Terlevich: Compact supernova remnants

suggests that a substantial fraction of the total kinetic energy was radiated in this phase, supporting the classification as a compact remnant rather than a supernova (see also Chugai 1992) Radio emission at the level of about 1 mJy at 6cm was detected in 1990 July (Sramek et al. 1990). This detection indicates that SN 1988Z is among the brightest RSN observed; in fact it is even brighter than SN 1986J. This suggests a link between RSNe and cSNRs, or between strong radio emission and strongly radiative shocks. This is confirmed by the analysis of the optical and radio data of SN 1986J and comparison with the model predictions. 5.2 SN 1986J Powerful radio emission was detected in the disk of NGC891 in 1986 August by Rupen et al. (1987). With a radio luminosity equivalent to 2,000 times Cass A, this was, at that time, the most luminous RSN discovered. The optical spectra was dominated by broad Ha emission with a FWHM= 1070 km s"1 a very high Balmer decrement Ho /H/3 ~ 60, and broad [OIII] lines. The location of this SN near the minor axis of an edge-on spiral galaxy and the large observed Balmer decrement seem to suggest a large amount of reddening. The estimated time of the explosion is 1982 September, giving an age of about 4 years at the time of the spectral observations. The fact that it was possible to see the SN 4 yr after explosion is apparently in contradiction with the expected large extinction. The observed average Ha luminosity in 1986 September/October is 5.1 x 1038 ergs" 1 and the observed average Balmer decrement is about 60. The age of the remnant was, at the time of the observations, more than 5 years or almost 8 tsg. For this age we obtain from Table 1 a shock velocity of 1400 kms" 1 , an Ha luminosity of 8 x 1040 ergs" 1 and a Balmer decrement of 14. Assuming 14 to be the intrinsic Balmer decrement, an extinction coefficient of c = 1.9 would be needed to explain the observed Balmer decrement. Correcting the observed Ha flux for this value of c gives an intrinsic luminosity for SN 1986J of 4.2 x 104Oergs~1 . This result is within a factor of two of the predicted value of the Ho flux (Table 1.) Bregman and Pildis (1992) detected X-ray emission from SN 1986J in 1991 August. The thermal emission model fits show two \ 2 minima: a low column density x 2 minimum giving a best fit with a luminosity 2.5 X 1O40 ergs" 1 between 0.1 and 10 keV, for T = 3.9 keV and a high column density X2 minimum indicating a softer spectrum (T = 0.35 keV) and an X-ray luminosity of 1042 ergs" 1 . When fitted by a power law over a bandpass of 0.1 - 2.5 keV, the spectrum of SN 1986J is described by a very soft

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spectral index of a = —2.1. This again suggests the presence of a very soft component in the X-ray spectrum. At the time of the X-ray detection the age of the cSNR was 10 years or I6tsg. The predicted X-ray luminosity at that time is about 2.5 x 10 41 ergs" 1 , ten times larger than the best fit solution but smaller than the value for the soft spectrum case. At this age the shock spectrum is very soft with a temperature of about 0.4 keV, similar to the best fit for the second X2 minimum in the thermal solution.

5.3 SN~1955 in M82 (41.9+58) This RSN is the most luminous radio source in M82. At a distance of 3.3 Mpc this is also the nearest example of a RSN. Its spectral evolution has been followed since 1973 and it shows increasing optically thick flux at low frequencies and decreasing flux at high frequencies (Kronberg et al. 1985). This RSN has a radio luminosity equivalent to more than 200 times that of the most luminous galactic SNR, Cass A, and together with SN 1986J, are the only two RSNe for which VLBI radio images are available. The measured diameter and the inferred column density (from the observed critical frequency of thermal absorption) requires that the shell of pre-supernova ejecta had a thickness of about 3 x 1016 cm and a density of no ~ 107 cm"3 at the time of the Kronberg et al. (1985) observations. These values are remarkably similar to the expectations of our canonical model in Tables 1 and 2. All the evidence points to a massive (M> 20 M@) progenitor. A compact and variable X-ray source with a luminosity of about 3 X 1039 ergs" 1 in the band 0.1-2.5 keVhas been identified at the radio position of this RSN (Kronberg et al. 1985, Collura et al. 1994). The X-ray flux seems consistent with the prediction in Table 1 and the detection of rapid X-ray variability on time scales of less than a day raises very important questions relating to the origin of such variability (Terlevich and Fabian in preparation). 5.4 SN 1981F SN 1987F was discovered in 1987 March. Early spectra revealed broad Ha superimposed on a nearly featureless continuum. The Ha profile did not have the characteristic P-Cygni profile of type II SNe (Filippenko 1989). The observed light curve was extremely flat (although steeper than that of SN 1988Z), fading by about 2.5 magnitudes in 400 days in the R. band; a typical type IIP will fade about 5 magnitudes in V in the same interval of

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time. A few months after discovery the spectrum was dominated by broad hydrogen, Fell and Call emissions. The estimated density of the ejected envelope at this late time was > 107 cm" 3 . The maximum in the observed Ha luminosity was 2.0 X 1041 ergs" 1 at an age of about 185 days. The FWHM of the Ha emission was about 7000 km s - 1 at that time. Comparing the results of the observations with the models of Table 1 and 2 we see that again the simple homogeneous models for n0 = 107 cm"3 give a reasonable description of the line widths and Ha luminosities. 6 Conclusions and future prospects We have seen that the optical and radio observations of cSNRs indicate a high circumstellar density of around 10" cm"3 and remnant sizes of a few times 1016cm with line widths (FWHM) between 1,000 and 10,000 kms" 1 . The size, Ha luminosity, X-ray luminosity and time scale of evolution predicted by the simple spherical homogeneous model of a radiative SN shock in a medium with density 107 cm"3 shows good agreement with the observations of cSNRs. This gives support to the idea that these systems are indeed strongly radiative shocks in the circumstellar environment. The study of strongly radiative astrophysical shocks is at an early stage. These shocks provide a unique laboratory capable of yielding valuable insights into physical processes in extreme conditions. They give unique information about the CSM and the ejecta of massive stars and can provide clues about some important aspects of galaxy evolution and of active galactic nuclei. Much work is needed, on the one hand to identify cSNR candidates early enough in their evolution, on the other to obtain good quality data over the widest possible spectral range. Radio and X-ray frequencies are, besides the optical, the most important spectral windows for observations of cSNRs. High quality and high resolution spectral information is particularly needed in the X-ray region where most of the energy is radiated. UV monitoring with high time-resolution will be very valuable if cSNRs prove to be variable on time scales of weeks. Due to their high luminosities, the study of cSNRs can be performed even at very large distances. The sample of cSNRs listed in Table 3 is perhaps the best data set for the study of shocks in a high density environment. Future developments in the theory should explore the effects of cooling out of equilibrium in strongly radiative shocks, the generation of cooling instabilities and luminosity variability and the effects of magnetic fields in the evolution of these cSNRs.

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Acknowledgements I would like to take this opportunity to thank Elena Terlevich for her support and help with this paper. Most of the work reported here is the result of a long-term collaboration with Guillermo Tenorio-Tagle, Pepe Franco, Michal Rozyczka and Jorge Melnick. I would like to thank Nino Panagiafor invaluable discussions, and Gary Ferland for the use of his code CLOUDY. References Avedisova, V. S., 1974, SvA, 18, 283 Bertshinger, E., 1986, ApJ , 304, 154 Bregman, J. N. and Pildis, R. A., 1992, ApJ , 398, L107 Chevalier, R. A., 1974, ApJ , 188, 501 Chevalier, R. A., 1982, ApJ ,259, 302 Chevalier, R. A., and Imamura, J. N., 1982, ApJ , 2C1, 543 Cioffi, D. F., McKee, C. F., and Bertshinger, E., 1988, ApJ , 334, 252 Collura, A., Reale, F., Schulman, E. and Bregman, J.N., 1994, ApJ , 420, 163 Chugai, N.N., 1992, SvA, 36, 63 Chugai, N.N. and Danziger, I.J., 1993, MNRAS, (in press) Draine, B.T. and Woods, D.T., 1991, ApJ , 383, 621 Falle, S. A. E. G., 1975, MNRAS, 172, 55 Falle, S. A. E. G., 1981, MNRAS, 195, 1011 Ferland, G., 1990, OSU Astronomy Dep. Internal Report (90-02) Filippenko, A., 1989, AJ, 97, 726 ' Imamura, J. N., 1985, ApJ , 296, 128 Kronberg, P.P., Biermann, P. and Schwab, F.R., 1985, ApJ , 291, 693 McCray, R, Stein, R. F., and Kafatos, M., 1975, ApJ , 196, 565 Rupen, M.P., van Gorkom, J.H., Knapp, G.R. and Gunn, J.E., 1987, AJ, 94, 61 Shull, J. M., 1980, ApJ , 237, 769 Sramek, R.A., Weiler, K.W. and Panagia, N., 1990, IAU Circ No. 5112 Stathakis, R.A. and Sadler, E.M., 1991, MNRAS, 250, 786 Tenorio-Tagle, G., Terlevich, R., Franco, J., and Rozyczka, M., 1994, (in preparation) Tenorio-Tagle, G., Bodenheimer, P., Franco, J., and Rozyczka, M., 1990, MNRAS, 244, 563 Terlevich, R., Tenorio-Tagle, G. Franco, J. and Melnick, J., 1992, MNRAS, 255, 713 Turatto, M., Cappellaro, E., Danziger, I.J., Benetti, S., Gouiffes, C. and Tarenghi, M., 1993, MNRAS, 262, 128 Weiler, K.W., Panagia, N., Sramek, R.A., van der Hulst, J.M., Roberts, M. and Nguyen, L., 1989, ApJ , 336, 421 Weiler, K.W., Panagia, N. and Sramek, R.A., 1990, ApJ , 3C4, 611 Wheeler, J. C , Mazurek, T. J., and Sivaramakrishnaii, A., 1980, ApJ , 237, 781

The evolution of compact supernova remnants Guillermo Tenorio-Tagle Instituto de Astrofisica de Canarias, 38200 La Laguna, Tenerife, Spain.

1 Introduction This is a short summary of several calculations of the evolution of supernova remnants in a constant high density medium no > 106~8 cm"3 and an abundance in the range O.O1Z0< Z < IOZQ. The main difference found when comparing them with the standard calculation of a supernova evolving into a constant density medium HQ — 1 cm"3 is that radiative cooling becomes important very early in the life of the remnants. The radiative phase starts well before the ejecta is fully thermalized and while the expansion velocities are still in the range of several thousands of kms" 1 . Consequently, the remnants miss their Sedov evolutionary phase and, unlike the standard case, in these calculations full thermalization of the ejecta is only completed long after the moment of thin shell formation (see Terlevich et al. 1992, 1994a; hereafter referred to as papers I and II). The cooling event leads to large luminosities (> 109 LQ) in spans of time of only a few years, causing a major rapid depletion of the supernova's stored thermal energy, in only a few weeks or months. Strong radiative cooling leads to an ionizing spectrum (see paper I) and thus to an HII region with multiple components, as it photoionizes the recombining, rapidly-moving swept-up gas and the outer unperturbed matter. The ionizing radiation is also absorbed by the still unshocked and dense expanding ejecta. Such remnants, hereafter termed "compact SNRs", are capable of producing the strong ionizing flux that makes them appear as Seyfert I impostors (Filippenko 1989) when occurring in dense regions far away from the nucleus of galaxies. Also, their rapid hydrodynamical evolution leads to the dimensions and luminosities as well as to the expansion velocities detected in radio supernovae (Weiler et al. 1989). Their enhanced radiation is primarily promoted by the high background densities, which may have been established during a slow and massive wind prior to the supernova explosion and contained as circumstellar matter by the high pressure of the surrounding gas. Here we describe the basic time evolution of compact SNRs looking also at particular moments that, when thoroughly analysed, lead to an alternative sound explanation to the phenomena observed in active galactic nuclei. 166

G. Tenorio-Tagle: The evolution of compact supernova remnants

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2 Numerical Calculations and Results 2.1 General Description Several two dimensional calculations of the evolution of supernova remnants in a constant high density medium (n > 106 cm"3) have been performed with the help of the hydro code described by Rozyczka (1985). Radiative losses, allowing in every run for a different Z metal abundance, were calculated from the basic interstellar cooling law of Raymond et al. (1976). The interstellar cooling law for the range of metal abundances considered here {\Q~2ZQ< Z < 10Z@) presents an increasingly larger cooling rate in the range 6 X 107 K > T > 3 x 105 K, as well as a shift of the minimum value of the function towards higher temperatures, for larger values of Z. The minimum value of the cooling function marks a definite change in the slope of the cooling curve implying that gas cooling from very high temperatures (say T > 4 x 107 K) enters, sooner or later, a cooling regime that causes a stronger cooling the further the gas cools down. This, together with the shift towards higher temperatures as a function of Z, marks the onset of strong radiative cooling in compact SNRs. 2.2 Boundary and Initial Conditions In all runs, the supernova ejected matter is at time t = 0 inserted in a. small portion of the computational grid. The procedure assumes an homologous growth and thus, the ejecta has a density proportional to R~fecta and velocity distribution proportional to Rejecta, as prescribed by the models of Arnett (1988). The insertion procedure ensures that 1051 erg are deposited in the form of kinetic energy with the largest speeds of about lC'kms" 1 . Given the homologous growth and consequent temperature drop of the ejecta, at t = 0 the amount of thermal energy in the grid is of the order of a few times 1048 erg. 2.3 The Time Evolution of Compact SNRs Figure 1 shows the run of luminosity versus time where a density no = 107cm~3, and an abundance Z = 10^0were assumed. The total luminosity of the remnant increases steadily at first, as t 0 8 , until strong radiative cooling becomes important leading to a total maximum luminosity of the order of a few times 1O9L0. The remnant then rapidly fades (L proportional to t~ n / 7 ) while exhausting its remaining energy. The main features, or main and secondary maxima, in the luminosity-time diagram (see Figure 1) correspond to major structural changes experienced by the compact

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G. Tenorio-Tagle: The evolution of compact supernova remnants

remnant as a function of time. The absolute maximum corresponds to rapid cooling behind the leading blast wave. This maximum occurs as soon as the temperature of the shocked gas decays to values T < 5 x 107 K. Before then, radiative losses from gas cooling from even higher temperatures, have only led to a steady growth of luminosity because of the larger size acquired by the remnant as a function of time. However, as soon as the temperature drops below 5 x 107 K the interstellar cooling function changes slope, leading to a larger cooling rate the further matter cools down, and this is sufficient in the high density regime to cause the run-away cooling described in detail in papers I and II. This rapid cooling causes a major loss of pressure behind the outer shock leading to its deceleration, while promoting the condensation of the swept up gas (see paper II). The second maxima occurs almost immediately afterwards. It is produced during the process of thin shell formation, when the cool swept up matter is collected together by the passage of several secondary shocks into a thin, cool and dense outer shell. Following this, conservation of momentum in the shell restores the leading shock velocity to values comparable to the ones it had before strong radiative cooling took place. Strong cooling behind the reverse shock produces the third apparent maximum in the total luminosity curve. This promotes the collapse of the swept up ejecta into a secondary thin and cool inner shell. The consequent loss of pressure behind the reverse shock leads to its withdrawal, and while favouring the approach and merging of the two shells, it delays the complete thermalization of the ejecta (see paper I). The last feature in the total luminosity curve produced mainly by an enhanced emission from gas cooling from T < 2 x 105K, is caused by the collision and merging of the two remnant shells.

2.4 Comparison with

Observations

A detailed comparison of the model results with the observations of AGNs can only be made with sources monitored for a. long period of time. The best example of these is perhaps NGC 1566 which has been followed for 15 years by Alloin et al. (1986). In that time four major periods of activity were detected. Each of these lasted about 1500 days while releasing 1051 erg. We associate these four events with four supernova explosions and their rapidly cooling compact remnants. The last one of these energy outbursts was observed with a higher temporal resolution to show a series of three (maybe four) rapid bursts (see Alloin et al. 1986, their Figure 4). Each of these bursts presents a steep rise time of about 20 days and a much longer

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169

decay of about 400 days. Similar features are present in the results displayed in Figure 1. Alloin et al. (1986) noted the spectral resemblance that the NGC 1566 energy bursts from January and November 1982 have with Type II supernovae, and in particular with SN 1970g. All rapid energy bursts within a major period of activity could however be better matched if interpreted as produced by a rapidly evolving compact SNR. In this way, the FWHM of the emitting lines will reflect the remnant expansion speed at the time that strong radiative cooling sets in, which is of the order of a few thousands of k m s " 1 , in agreement with the 2000km s" 1 observed in NGC 1566. Also as shown in the previous section, the maximum intensity of consecutive bursts decays steadily but rather slowly as a function of time. Factors of 4 - 10 between first and last luminosity maxima are indicated in the various runs. This is in good agreement with NGC 1566 and thus t.lie decay should not be interpreted in terms of the light curve of a. supernova, which as noted by Alloin et al. (1986), fades orders of magnitude earlier than the luminous events expected during the time evolution of compact SNRs.

3 T h e n a t u r e of t h e lag A further detailed analysis of the events promoted by strong radiative cooling had led to an important result concerning the applicability of the starburst model to the realm of AGNs. The lag, the observed delay between abrupt changes in the continuum ionizing radiation followed after some time by changes in the intensity of the emission lines from the broad line region (BLR) of AGNs have been accurately matched by the cooling events during maximum luminosity of compact SNRs (see paper II). The lag, usually interpreted as a result of the geometry of the accretion disk sources, implying a measure of the distance at which the BLR sits away from the ionizing source, and that has also been used to deduce the mass of the putative central black hole (see e.g. Netzer 1993), has a different explanation in the starburst model of AGNs. Here, the lag results from the time-dependent changes in the ionization parameter within the layer of gas swept by the supernova blast wave, as matter adjusts itself to the drastic drop in pressure suddenly promoted by strong radiative cooling, and consequently in this model the lag has nothing to do with the size of the emitting source! Detailed calculations of radiative shocks evolving in a high density medium performed with a greater time resolution have shown in detail the sequence of events that take place as remnants approach and reach maximum luminosity. The matter involved in the process suddenly has to readjust to the

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0.1

TIME ( y r )

10

Fig. 1. The luminosity output, of CSNRs. The integrated luminosity produced by CSNRs throughout their evolution, versus time.

large pressure imbalance promoted within the shocked gas by the onset of strong radiative cooling. The final outcome is the formation of a. thin shell at the edge of the remnant, several orders of magnitude denser than the original background medium. Gas condensation however, does not happen immediately, as it requires of the passage of secondary shocks through the cool region for this to acquire the appropriate density (and thus pressure). The secondary shocks emanate from the hottest section of the remnant interior, overtaken earlier during the evolution when the shock speed was larger, and that had taken longer to cool. The shocks follow the blast wave in an attempt to communicate the interior pressure, however, given the increasingly larger densities behind them and thus the correspondingly shorter cooling distances, inevitably become also rapidly radiative. Meanwhile cooling proceeds and moves as a wave, ahead of the secondary shocks, into gas more recently overtaken by the progressively slower blast wave to eventually catch up with it. The blast wave then slows clown for two reasons: because of ge-

G. Tenorio-Tagle: The evolution of compact supernova remnants ometrical dilution as the remnant grows and because it has now suddenly lost its piston pressure due to strong cooling behind it. The gas steadily overtaken by the cooling front is the source of ionizing continuum radiation, to be observed as a variation in the continuum of the AGN. This radiation is immediately absorbed by the reshocked matter. By the cool layer of gas continuously changing density after the passage of secondary shocks. The combination of a steadily denser layer constantly irradiated, as the cooling wave progresses through the layer of shocked gas, leads to a continuous decrease in the effective ionization parameter and results into a rapidly changing ionization structure of the fast-moving photoionized gas. The numerical calculations show that the width of the photo-ionized shocked region, traversed by the cooling front and continuously swept by secondary shocks to condense it into a cool thin shell, is only about 1013cm (with a light travel time of 103 sec) and yet, lags of up to several days, weeks, and even months, are generated for different lines. In general, the calculations predict shorter delays for high ionization lines than for low ionization ones. A detailed comparison with the results from the NGC5548 extensive monitoring campaign, agree both with the time delays for different lines, and with the intensity values reach by the various lines (Paper II). These results have also been independently corroborated with a different computational scheme (see Plewa et al. 1994). The calculations thus show that the compact supernova remnant model is capable of giving an accurate and detailed description of the temporal behaviour of the BLR, as well as accounting for all of its intrinsic properties with a, minimum of free parameters. 4 Rapid X-Ray Variability in Compact SNRs The role played by inhomogeneities in the supernova ejecta on the generation of rapid X-ray variability has also been recently explored (see Terlevich et al. 1994b). The interaction of supernova, fragments with the internal structure of a compact supernova, remnant has been followed with analytical approximations and 2-D hydrodynamical simulations. Here we focus on those interactions expected to lead to the largest energy bursts and the shortest durations. The model calculations show the evolution of fragments into thin, dense and cold "tortillas" as they encounter the reverse shock. These tortillas eventually cross the hot cavity and collide at large speeds with the remnant outer shell, causing luminous and short lived bursts that radiate most of their energy at X-ray frequencies. The range of predictions of the models has been compared with the observations, focusing particularly

171

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G. Tenorio-Tagle: The evolution of compact supernova remnants

on time scales, luminosities and light curve properties of rapid X-ray variable AGN. The models account for the complete observed range, including the highest luminosities and the shortest X-ray flare time scales, as well as for the power spectrum observed in AGN. In the past few years evidence for the existence of dense and fast moving fragments of SN ejecta has grown substantially. Metal rich high velocity knots moving at thousands of kms" 1 , have been found in Cas A (Braun et al. 1987 and references therein), Tycho (Seward et al. 1983), Puppis A (Winkler et al. 1988), Kepler (Bandiera k van den Bergh 1991) and in a number of extragalactic remnants (Lasker & Galinowski 1991, Hanuschik et al. 1993, Fesen & Matonick 1993). Theoretical work, in particular for the case of SN1987A, strongly suggests that Rayleigh-Taylor instabilities are the most likely origin of the condensations in the ejecta (Arnett et al. 1989). The effects of a fragmented ejecta in the evolution of "normal" remnants, i.e. those evolving in a low density interstellar medium, have been investigated for a wide variety of conditions (see review by Franco et al. 1991). In particular, Tenorio-Tagle et al. (1991) and Franco et al. (1993) have studied the interaction of SN fragments with the internal structure of single and multiple SNRs. In small remnants, high density fragments move almost unimpeded through the hot cavity to be rapidly thermalized as they impact the outer remnant shell. The multitude of fragment-shell interactions result into well localized strong shocks that partially disrupt the shell while causing a larger X-ray remnant radius. Fragments form soon after the explosion, via Rayleigh-Taylor instabilities, deep in the inner parts of the ejecta, at the interfaces between layers of different chemical compositions (Arnett et al. 1989). The resultant fragment properties, thus depend on details of the structure of the ejecta (see Franco et al. 1993). For simplicity we have assumed that the fragments (f) _

n/(r)

have a density contrast e = . L\ > 1 with respect to the interfragment medium (IFM) which is constant with r, and a velocity characteristic of their location. Fragments of different sizes, and at different locations have been inserted within the ejecta for a number of calculations. In all cases the flow has been assumed to undergo an homologous diverging expansion and thus the value of e remains constant for each fragment and its neighbouring gas, until the fragments meet the reverse shock. The resolution of the numerical grid limits the minimum initial dimensions of a properly resolved fragment, although the homologous expansion soon leads to large fragment sizes. Note that the evolutionary features and the physics of the interaction are independent of the fragment size, and thus the calculations accurately

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illustrate the time evolution. Figure 2 shows isodensity contours and the velocity field at selected evolutionary times for one of the calculated cases. As described in Section 2, the remnant caused by the violent release of the ejecta presents two well-defined shocks: the leading blast wave and the reverse, inward-facing, shock. At first, before t = 4yr, strong cooling is in full operation, but it is slow and the flow remains isobaric. During this time the density of the cooling gas changes only slightly. Soon afterwords, at t = 4 yr, a well defined thin, very dense cool shell is present right behind the blast wave. This also happens at t = 6.3yr behind the reverse shock where matter condenses also into a secondary inner shell. Note that at that time the ejecta is not yet fully thermalized and it is still impacting the disrupted shell with large velocities. At that time however, the energy left over amounts to about 1/10 of the initial value, and the luminosity of the remnant has fallen by an order of magnitude from the maximum emission (see Figure 1). The bulk of the thermal energy stored in the remnant was in fact rapidly radiated away before t = 3-5yr. From then onwards, a continuous and immediate release of energy occurs as soon as the remaining kinetic energy is thermalized. Figure 2 also shows the initial conditions, the homologous expansion, as well as the initial fragment interaction with the reverse shock. These are followed by frames showing the motion of the fragment through the hot cavity, as it is overtaken by a weaker reverse shock to evolve into a condensed "tortilla" with its smallest dimension along the flow direction, while catastrophic cooling also causes the formation of a thin dense shell at the edge of the remnant. Note that the lateral walls of the fragment are also condensed by a strong shock that develops as soon as the fragment sides enter into contact with the hot cavity. The final frames display the tortilla-outer shell interaction.

4-1 Analytical Description

of the X-Ray

Bursts.

All ejected fragments eventually meet the reverse shock which tends to decelerate them while thermalizing their kinetic energy. The shock strength is inversely proportional to the square root of the incoming gas density and thus, high density contrast fragments (e > 10) are compressed by weaker sections of the reverse shock. Given the high densities, and the correspondingly smaller cooling times these weaker shocks are regarded as isothermal, leading to large compression factors. Meanwhile, the shocked fragments acquire a tortilla shape with its smallest dimension in the radial direction. These cold tortillas move with high velocities through the hot cavity generated by the reverse and outer shocks. The tortillas eventually reach the edge

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G. Tenorio-Tagle: The evolution of compact supernova remnants

'•».. = I.IS x 10* cir

fragment

' =• 1.44 X 10' 1 " - . . = 8.C5 x 10* cm I"'

I = l.2< X 10* 1 "n«» = 3.94 x 10' cm •"

I = J.12 x 10' >

1 =r 1.02 X 10* •

« - . , = 6.00 x 10* cm I"1

t ' ^ , o 3.80 x 10* cm s"

' « 6 67 x 10' • ••-.. - 4.65 x 10« cm i - 1

Fig. 2. The evolution of cSNRs. Two dimensional representation of the density distribution and velocityfieldwith scales indicated at various times. The contours are logarithmically spaced with 6\ogn = 0.33. The crowding of contours indicates shocks and/or density discontinuities while sudden velocity changes also help to identify the location of the various shocks developed throughout the evolution.

G. Tenorio-Tagle: The evolution of compact supernova remnants

175

of the remnant and either catch up with the leading shock or the thin shell of compressed swept up matter depending on whether or not catastrophic cooling has already taken place behind the leading shock. The interaction of the fast moving tortillas with the outer structure generate bursts of X-ray emission promoting the rapid and strong X- ray variability of the source. The thin shell bombardment phase can last several years during which the remnant is highly variable in X-rays on time scales down to one hundred seconds. The parameters of the tortillas (denoted by superscript 't1) can be estimated from Papers I and II. In the cSNR model of Paper I, at t = tsg = 230 days, the leading shock velocity is V3hock = 5000 km s" 1 and its radius is Rshock — 3.0 X 1016cm, while the outer thin shell has a density nshell ~ 1012 cm~ 3 and thickness ARsheii ~ 1011 cm. At the same time the density of the IFM just in front of the reverse shock is very similar to that of the circumstellar medium, i.e. ~ 10 7 cm~ 3 . A fragment with e = 100 has a density of nj = 10 9 cm~ 3 just before it meets the reverse shock. The reverse shock effective velocity in the IFM is about 5000kms - 1 , and thus the shock velocity into the fragment is 909km s ^ a n d the cooling time is about 2.2 x 10'1 sec. A fragment with initial size / = 10 14 cm, has a total shock crossing time of 1.1 x 106 sec and therefore is completely shocked before reaching the thin outer shell. As the cooling time is much shorter than the crossing time, the shock is isothermal, and the compression factor equals the square of the Mach number M. For the parameters considered here, M2 — (909/10) 2 = 8100 (assuming that the sound speed is 10 km s~afor the fragment gas in front of the shock). The final density of the fragment is about nj = 8 x 1012 cm" 3 and as the compression is only in the radial direction the final result is a "tortilla" 10 14 cm in diameter and with thickness A/' = 1.2 X 10 10 cm. The tortilla leaves the reverse shock with a space velocity of F* = 9900 km s" 1 , corresponding to a relative speed (with respect to the shocked IFM) of vt~IFM = 4900 k m s " 1 . After about t ~ 1.0 x 10 1 6 /4.9 x 108 = 2 x 107sec, the tortilla reaches the edge of the remnant. There, due to the deceleration of the outer shell, the relative velocity between the tortilla and the shell is about 7000 km s - 1 a t the time of the collision. The densities of the thin shell and the tortilla are very similar and around n* ~ nsheii — 3 x 1012 cm" 3 . The collision leads to two similar shocks, one into the tortilla the other into the shell. Given the similar densities the velocity of both shocks is about 3500 km s" 1 . For the velocities and densities involved in this interaction the cooling time is, rcoo\ ~ vg / n.\3 ~ 1.8 minutes while the shock crossing time is Tlxi = .,?'' ~ V2*-1,1?^ ~ 0.5 minutes. shock

This implies that the kinetic energy of the tortilla is thermalized and emit-

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G. Tenorio-Tagle: The evolution of compact supernova remnants

ted in about a cooling time, i.e. ~ 2 minutes. The mass of the tortilla is Ml = (/') 2 A/*n* m; ~ 2 x 10 27 gm ~ 1O~ 6 M 0 , and its kinetic energy, E

k = \Mt

(y1-3^")

~ 5 x 1044 erg. The maximum luminosity reached

by a flare of duration Atfiare

is Lx — h-AT^— — 2 x l O ^ e r g s " 1 with

&tflare = \ Tcool — 1 min, and typical energy about 15keV. In general, larger tortillas produce more energetic flares and last longer. The above estimates of duration and peak luminosity corresponds to the "face-on" case when the observer sees the whole event simultaneously, case that corresponds to an interaction at the central part of the remnant shell as viewed by the observer. Interactions in other positions produce a light-travel time delay between different points in the tortilla. In the case of the /' = 1014 cm tortilla the maximum delay corresponding to an edge-on view is 3330 sec.

5 Concluding remarks In the starburst model of AGNs, sometimes viewed as exotic and/or unconventional, the applied physics are in fact most, conventional, as it uses the little, or the lot, that we know about real events: the physics of stars and stellar evolution and their interaction with the surrounding gas, and with these sound predictions are made. Under the assumption of a normal IMF, the supernova rate detected in NGC 1566 (about 1 every 4.5 years) should result from a coeval major burst of stellar formation, leading to a nuclear cluster with a mass of 5 x 108 MQ. Every supernova explosion from the evolving cluster that occurs in a high density medium (highly likely produced by slow winds from the progenitors) will lead to a compact SNR. The photoionizing radiation arising from the cooling shocks of each remnant will immediately be absorbed by the rapidly expanding recombined gas, producing as shown in papers I and II the broad emission lines with the intensities, relative ratios and lags with respect to the ionizing contunuum, in excellent agreement with the BLR of active galactic nuclei. The supernova rate will lead to detectable bursts of energy each releasing, a total of about 1051 erg in only a few years (< lOyr). By then the size of the remnants will be of the order of a few times 1016 cm. Each remnant will lead to several distinct energy bursts which have here been identified with: 1) the onset of strong radiative cooling behind the blast wave; 2) the completion of thin shell formation; 3) The onset of strong cooling behind the reverse shock; and 4) the collision and merging of both shells. The energy burst caused by these events agrees well with detailed observations of the

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long-term variability in AGNs. The remnants will effectively end up their evolution as rather fragmented shells in a time longer than the SN rate of the nuclear star cluster, thus providing at all times a minimum intensity value of the emitting region. During the evolution of each compact SNR (< 10 years), more than 90% of the initial energy of the explosion (1051erg) is radiated away, while the remnants acquire dimensions of only a few times 1016cm. The high circumstellar densities thus provoke the rapid onset of strong radiative cooling and with it the various hydrodynamical events that lead to the broad line region, to the continuum and line variability and also to the lag typical of these sources. The original problem of processing 1051 erg into the medium surrounding an exploding star acquires a. number of new facets when the early fragmentation of the ejected matter is taken into consideration. This has been shown here to be directly related to the rapid X-ray variability of AGN if the remnants evolve in a high-density circumstellar medium, as postulated in the Starburst model of AGN. Under such conditions, strong radiative cooling rapidly begins to drain the thermal energy of the remnant even before the ejecta is fully thermalized. Therefore, most of the ejected fragments, particularly denser ones like the oxygen-rich, fast-moving (5000 km s"1) knots of Cas A, experience a complex evolution to end up colliding at large relative speeds with the remnant outer shell. This is one of several possibilities capable of generating rapid X-ray variability in r.SNRs. Perhaps the one that leads to the most energetic and shortest duration energy bursts, reason that motivated this first exploration.

Acknowledgements. GT-T acknowledges support from the EEC grant for international collaboration No CI1*-CT91-O935. References Alloin, D., Pelat, D., Phillips, M.M., Fosbury, R.A.E. k Freeman, K., 1986, Ap.I, 308, 23. Arnett, W. D., 1988, ApJ, 331, 377. Arnett, W. D., Fryxell, B. & Miiller, E., 1989, ApJL, 341. L63. Braun, R., Gull, S. F. & Perley, R., 1987, Nature, 327, 395. Bandiera, R. fc van den Bergh, S., 1991, ApJ, 374, 186. Fesen, R. A. fc Matonick, D. A., 1993, ApJ, 407, 110. Filippenko, A., 1989, AJ, 97, 726. Franco, J., Tenorio-Tagle, G. Bodenheimer, P. & Rozyczka, M., 1991, PASP, 103, 803. Franco, J., Ferrara, A., Rozyczka, M., Tenorio-Tagle, G. fc Cox, D., 1993, ApJ, 407, 100. Hamilton, A. J. S., 1985, ApJ, 291, 523.

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Hanuschik, R., Sypromilio, J., Stathakis, R., Kimeswenger, S., Gochermann, J., Seidensticker, K.J. k Meurer, G. 1993, MNRAS, 261, 909. Lasker, B. M. k Golinowski, D. A., 1991, ApJ, 371,563. Netzer, H. 1993 in The Nearest Active Galaxies eds J. Beckman, L. Colina k H. Netzer. Madrid. Coleccion Nuevas Tendencias Vol 22, p. 219. Plewa, T., 1994 submitted to Violent Star Formation, from SODoradus to QSOs, ed G. Tenorio-Tagle, Cambridge University Press. Raymond, J., Cox, D. P. k Smith, B. W., 1976, ApJ, 204, 290. Rozyczka, M., Franco, J., Miller, W., Tenorio-Tagle, G. &: Terlevich, R., 1994, (in preparation). Rozyczka, M., 1985, A&A, 163, 59. Seward, F., Gorenstein, P. k Tucker, W., 1983, ApJ, 266, 287. Stathakis, R. A. k Sadler E. M., 1991, MNRAS, 250, 786. Tenorio-Tagle, G., Rozyczka, M., Franco, J. k Bodenheimer, P., 1991, MNRAS, 251, 318. Terlevich, R., Tenorio-Tagle, G., Franco, J. k Melnick, J., 1992, MNRAS, 255, 713 (Paper I) Terlevich, R., Tenorio-Tagle, G., Rozyczka, M., Franco, J. k Melnick, J., 1994a, MNRAS, (in press) (Paper II). Terlevich, R., Tenorio-Tagle, G., Cid-Fernandes, R. Franco, J. k Rozyczka, M., 1994b, MNRAS, submitted (Paper III). Winkler, P. F., Tuttle, J. H., Kirshner, R. P. k Irwin, M. J., 1988, Supernova Remnants and the Interstellar Medium (IAU Coll. 101), ed. R. S. Roger and T. L. Landecker (Cambridge University Press: Cambridge), p. 65. Weiler, K. W., Panagia, N., Sramek, R. A., Van der Hulst, J. M. Roberts, M. S. k Nguyen, L., 1989, ApJ, 336, 421.

Massive Supernovae in Binary Systems P. C. Joss 1 , J. J. L. Hsu 2 , Ph. Podsiadlowski3, and R. R. Ross4 1

2 3 4

Department of Physics, Center for Space Research, and Center for Theoretical Physics, Massachusetts Institute of Technology, Cambridge, MA 02139, U.S.A. Department of Astronomy, University of California, Berkeley, CA 94720, U.S.A. Institute of Astronomy, Madingley Road, Cambridge CB3 OHA, U.K. Department of Physics, College of the Holy Cross, Worcester, MA 01610, U.S.A.

Abstract The presence of a close binary companion can affect the evolution of a massive star through one or more episodes of mass transfer, or by merger in a common-envelope phase. Monte Carlo calculations indicate that ~ 20 - 35% of all massive supernovae are affected by processes of this type. The duplicity of the progenitor may be revealed by the illumination, in the supernova event, of axially symmetric material that had previously been ejected during the mass-transfer phase or by the expulsion of a common envelope. Moreover, the properties of the progenitor star, the peak supernova luminosity, and other observable features of the supernova event can be affected by prior binary membership. Binary interactions may be the cause of much of the variability among Type II supernova light curves, and may result, in Type Ib or Ic events in cases where the entire hydrogen-rich envelope has been stripped from the progenitor. Many of the peculiarities of SN 1987A and SN 1993J may well have resulted from the prior duplicity of the progenitor.

1 Introduction A large fraction of all stars are members of binary systems. It is therefore reasonable to consider the possibility that the properties of many massive supernovae (i.e., supernovae whose progenitors had initial main-sequence masses, Mms, greater than ~ 8 MQ) are influenced by prior interactions of the progenitor with a binary companion star. This possibility was brought into focus in recent years by the nearby Type II supernovae SN 1987A and SN 1993J, many of whose properties differed markedly from theoretical expectations. As a result, several studies have been undertaken to estimate the frequency of massive supernovae in binaries and the unique properties of the progenitors and the resultant supernova events that result from the evolution of a massive progenitor in a binary system (Podsiadlowski, Joss & Hsu 1992; Tutukov, Yungelson & Iben 1992; Hsu et al. 1993). Of particular interest, in the context of these Proceedings, is the possibility of mass ejection from the presupernova binary in an axially symmetric pattern, with 179

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the axis of symmetry coinciding with the orbital axis, and the subsequent illumination of the ejected matter by the supernova event. We here describe the main results of recent theoretical work on massive supernovae in binaries and briefly discuss the application of this work to SN 1987A and SN 1993J.

2 Evolution of Massive Supernova Progenitors in Binaries The principal effects of evolution in a close binary system on the progenitor of a massive supernova can be broadly divided into three categories (Podsiadlowski, Joss & Hsu 1992; hereafter PJH): (1) loss of part or all of the stellar envelope to the companion star, (2) accretion of matter from the companion star, or (3) merger of the two stars in a common-envelope phase. (In addition, a star in a close but detached binary may lose a large fraction of its envelope in an enhanced stellar wind [Vanbeveren 1987; Tout & Eggleton 1988] whose time-averaged morphology will display axial symmetry.) In cases (1) and (2), it is highly likely that the mass-transfer process will entail the loss of a. significant amount of matter from the system, while in case (3) the common envelope itself may well be ejected prior to the supernova event. On the basis of Monte Carlo calculations, PJH concluded that ~ 20-35% of all massive stars experience binary interactions of one of the above types before undergoing a supernova explosion. This is consistent with the findings of Tutukov et al. (1992), who concluded, by somewhat different means, that ~ 25-45% of all supernovae (including those involving low-mass progenitors) originate in initially close binaries. In the following paragraphs, we describe the salient features of each of the three modes of presupernova binary evolution described above.

2.1 Mass-Loss

Models

If the supernova, progenitor was originally the more massive of the binary components, it can lose mass to its companion via Roche-lobe overflow. This scenario has been considered in detail by Joss et al. (1988) and PJH. If the star first fills its Roche lobe while it is still on the main sequence, a contact system and eventual merger of the binary components is likely to result; the merged star should then have the properties of a rejuvenated main-sequence star. (Any mass that is lost from the system as a byproduct of this process is likely to dissipate before the merged star is able to reach the supernova stage.) Of greater interest, in the present context, is the possibility that the primary first fills its Roche lobe during the course of its post-main-sequence

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evolution. If the masses of the binary components are not too different, the resultant mass transfer can take place on a sufficiently long time scale that a common envelope does not form. In cases where the entire hydrogen-rich envelope is lost, the progenitor will become a helium star and will likely end its life as a Type Ib or Ic supernova. If, however, the mass-transfer process terminates when the progenitor still retains at least a few tenths of a solar mass of its envelope, its final presupernova radius and effective temperature will be nearly the same as those it would have had in the absence of mass loss. Monte Carlo calculations (PJH) indicate that a few percent of all massive stars (perhaps up to ~ 5% if systems with binary-enhanced winds are included) become supernovae of this latter type.

2.2 Accretion

Models

The original secondary in a close binary system with a Roche-lobe filling primary should accrete a substantial fraction of the mass lost by the primary. If the mass transfer commences before the original secondary has completed core hydrogen burning, the subsequent evolution of the secondary should mimic that of a more massive main-sequence star (Hellings 1983; PJH). Analogously to the case of mass-loss models, of greater interest here is the situation where mass transfer commences only after the original secondary has left the main sequence (Podsiadlowski & Joss 1989; de Loore & Vanbeveren 1992; PJH). Due to the accreted mass, the original secondary will generally become the more massive of the two stars, and its concomitantly accelerated evolution may cause it to reach the supernova stage prior to the original primary. If the orignal secondary is the first star to become a supernova, it will have a normal post-main-sequence companion at the time of the explosion; if, instead, the original primary readies the supernova, stage first, it should leave a neutron-star or black-hole remnant that will remain gravitationally bound to the original secondary until it, too, becomes a supernova. In either case, however, the explosion of the original secondary will eject more than half of the residual mass of the system, generally causing it to become unbound. Nevertheless, the companion object may become detectable after the supernova photosphere has receded sufficiently. Another diagnostic of supernova events of this type is the color of the immediate supernova progenitor; if the original secondary accretes a sufficient amount of mass, it will end its life as a blue supergiant rather than a red supergiant, which is the generally expected precursor for a Type II supernova that has evolved in isolation (Falk & Arnett 1977; Woosley k. Weaver 1985).

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2.3 Merger Models If the initial masses of the two stars are sufficiently different, the time scale for mass transfer, once it commences, will be much less than the Kelvin time of the secondary. As a result, the mass transfer will be unstable, and the system will develop a common envelope (Paczyriski 1976; Kippenhahn & Meyer-Hofmeister 1977; Podsiadlowski, Joss & Rappaport 1990; PJH); the primary should lose its entire hydrogen-rich envelope to the common envelope. Thereafter, dynamical friction between the secondary and the common envelope will cause it to spiral in toward the system center-of-mass. It is uncertain whether or when the common envelope will subsequently be ejected (see Hsu et al. 1993 for a discussion). If the envelope is ejected before either the secondary is dissolved or the binary components merge to form a single star, and if the core mass of the primary is greater than ~ 1.4 MQ at the time of the ejection, a Type Ib or Ic supernova may result; however, if the supernova explosion strips off a significant amount of the hydrogen-rich envelope of the secondary, the supernova event may be misclassified as Type II. In cases where mass transfer commences when the primary is still on the first red giant branch (case B transfer) and the common envelope is not subsequently ejected, the spiral-in time scale should be much shorter than the remaining evolutionary time for the primary; the binary components should therefore merge before a supernova, event occurs. If, instead, mass transfer does not commence until the primary has reached the asymptotic giant branch (case C transfer) and the common envelope is not ejected, it is uncertain whether merger will occur before the the primary becomes a supernova. When the binary components merge before the occurence of the supernova event, the net effect is very similar to that of the accretion scenario described in §2.2, and the merged star may well end its life as a blue supergiant. If the merger is not yet complete by the time of the supernova event, the immediate progenitor (i.e., the common envelope itself) may have the appearance of either a. red or a blue supergiant, depending on the values of various parameters for the initial binary system and the details of the common-envelope evolution. 3 Hydrodynamics of Massive Supernovae in Binaries We have recently completed a series of hydrodynamic calculations to explore the consequences of mass-loss and accretion/merger scenarios for the observational properties of the resultant supernova events (Hsu et al. 1993). We restricted our attention to cases where the progenitor retained at least a portion of its hydrogen-rich envelope, so the the event would be of Type

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II; the hydrodynamics of events in which the entire hydrogen-rich envelope had been stripped from the progenitor, leading to Type Ib or Ic events, has previously been explored by Ensman & Woosley (1988) and Shigeyama et al. (1990). We here briefly summarize the results of some of our calculations. To investigate the effects of mass loss from a massive supernova progenitor, either through mass transfer to a close-binary stellar companion (see §2.1) or via a strong intrinsic stellar wind, we followed the explosion of a star with an initial main-sequence mass of 12 MQ. In the absence of mass loss, such a star would have a hydrogen-rich envelope of mass M env ~ 8.7 MQ; we considered cases with residual envelope masses of 0/1, 1.9, and 4.9 MQ, as well as a case with no mass loss. The visual light curves of our four mass-loss models are shown in the lefthand panel of Figure 1. The most dramatic effects of a reduced envelope mass are (1) a much more rapid rate of decline of the light curve, (2) a higher peak luminosity (by as much as two magnitudes in the V band), and (3) peak photospheric velocities that are higher by as much as a. factor of ~ 2 (~ 2 X 104 km s" 1 for the model with the smallest residual envelope, compared to ~ 1 X 104 km s" 1 for the model with no mass loss). In order to explore the effects of an increase in the envelope mass of the progenitor via accretion from or merger with a binary companion during the course of its post-main-sequence evolution, we calculated the explosion of three stars, each of which had a final presupernova mass of 20 MQ. The first star had an initial main-sequence mass of 20 MQ and underwent no mass loss or gain during the course of its evolution; the other two stars had M ms = 17 and 15 M 0 and gained 3 and 5 MQ, respectively, during their post-main-sequence evolution. Both of the models that gained mass were blue supergiants, rather than red ones, at the time of the supernova event (see §§2.2 and 2.3, Podsiadlowski & Joss 1989, and PJII). The visual light curves for these three models are shown in the right-hand panel of Figure 1. The principal effects of the addition of mass are (1) a peak luminosity that is fainter by as much as 3.5 magnitudes in the V band (if we exclude the initial flash near / = 0, which is not modeled very accurately in our calculations), (2) a more rapid decline of the light curve, and (3) a reduction of the peak photospheric velocity by as much as a factor of ~ 2.5, from ~ 1.6 X 104 km s" 1 for the constant-mass model to only ~ 6 X 103 km s" 1 for the model that has gained 5 MQ during the course of its evolution.

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t (days) Fig. 1. (a) Light curves (absolute visual magnitude, My, as a function of elapsed time, t, since core collapse) for four supernova models whose progenitors had initial main-sequence masses, A/ ms , of 12 M© but had lost their hydrogen-rich envelopes to varying degrees during the course of their post-main-sequence evolution. All four hydrodynamic calculations assumed an explosion energy, E, of 1 foe and no energy input from the decay of radioactive material. Solid curve, M env — 0.4 MQ. Long-dashed curve, M env = 1-9 MQ. Short-dashed curve, A/env = 4.9 MQ. Dotted curve, M e n v = 8.7MQ (corresponding to no mass loss), (b) Same as (a), but for three models whose progenitors underwent accretion from, or merger with, a binary stellar companion during the course of their post-main-sequence evolution; the final presupernova mass was 20 M© in all cases. All calculations again assumed E = 1 foe and no energy input from radioactive decay, except where otherwise noted. Solid curve, M m s = 15 MQ. Long-dashed curve, M m s = 17 MQ. Dotted curve, Mms — 17 MQ, with additional energy from the radioactive decay of 0.071 MQ of Ni 56 and its decay product, Co 56 , deposited in the innermost layers of the ejecta; this light curve comes closest to matching the general properties of the light curve of SN 1987A, although it does not fit the observed light curve in detail (see Hsu el al. 1993 for a discussion). Short-dashed curve, M m s = 20 MQ, with no mass gained or lost by the progenitor during the course of its presupernova evolution, shown for comparison.

4 Application to Recent Supernovae 4.1 SN 1987A A number of authors (Fabian & Rees 1988; Joss el. al. 1988; Barkat & Wheeler 1989; Hillebrandt & Meyer 1989; Podsiadlowski & Joss 1989; Podsiadlowski, Joss & Rappaport 1990; de Loore & Vanbeveren 1992; PJH; Rathnasree 1993) have explored the possibility that Sk -69°202 , the progenitor of SN 1987A, had been a member of a binary system prior to the

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supernova event. Among the various binary scenarios that have been proposed, the most promising appears to be one in which Sk —69°202 underwent merger with a binary stellar companion in a common-envelope phase (Hillebrandt & Meyer 1989; Podsiadlowski, Joss & Rappaport 1990; PJH). (The plausibility of accretion scenarios has been somewhat diminished by the lack of evidence for either a prior supernova event or a normal or neutronstar companion following the recession of the photosphere of SN 1987A.) The major lines of evidence in support of a merger scenario include (1) the blue color of Sk — 69°202, which was in contrast to most prior theoretical expectations (see §§2.2 and 2.3, Podsiadlowski & Joss 1989, and PJH), (2) chemical peculiarities in the progenitor and in the supernova ejecta, which may result from the dredge-up of nuclear-processed material during the merger process (see Hillebrandt & Meyer 1989, PJH, and references therein), (3) the low peak luminosity of SN 1987A and the exceptionally strong effect of energy input from radioactive decay upon its light curve, which is in accord with the results of hydrodynamic calculations by Hsu et al. (1994) for accretion/merger models of Type II supernovae (see §3 and Fig. lb), and (4) the approximate axial symmetry of the circumstellar material (Wampler et al. 1990). In regard to this last point, it is intriguing to speculate that the ring structure around the supernova represents relatively dense, low-velocity material from the common envelope that was ejected in the orbital plane of the original binary, while the "Napoleon's Hat" nebulosity is material from a high-velocity stellar wind that was emitted by the hot progenitor star after the completion of the common-envelope phase and gained its axial symmetry by interaction with the pre-existing ring structure.

4.2 SN

1993J

The likelihood that the progenitor of SN 1993J lost, most of its hydrogen-rich envelope by transfer to a close binary companion, in the manner discussed in §2.1, has already been noted by a number of authors (Nomoto et. al. 1993; Podsiadlowski et al. 1993a; Ray, Singh k Sutaria 1993; Woosley et. al. 1994) and is discussed in some detail elsewhere in this volume (Podsiadlowski et al. 1993b). Here, we only observe that there may be an interesting evolutionary link between SN 1993 J and SN 1987A. If the companion of SN 1993 J accreted several solar masses of material during the mass-transfer process, at a, time when it had already evolved off the main sequence, it should end its life (~ 10 5 -10 6 years hence) as a blue supergiant, and this second supernova event should resemble SN 1987A. It is remarkable that both SN 1987A and SN 1993J, the two nearest known

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supernovae of the past century, have both displayed substantial evidence for origin in massive close-binary systems. Of course, evidence of prior duplicity becomes easier to obtain with increasing proximity of the supernova event. It is distinctly possible that, for supernovae of all types, prior duplicity will turn out to be the rule rather than the exception. Acknowledgements This work was supported in part by the U.S. National Aeronautics and Space Administration under grant NAGW-1545. References Barkat, Z. k Wheeler, J. C. (1989). Astrophys. J., 342, 940. de Loore, C. k Vanbeveren, D. (1992). Astr. Astrophys., 2G0, 27.$. Ensman, L. M. k Woosley, S. E. (1988). Astrophys. J., 333, 754. Fabian, A. C. k Rees, M. J. (1988). Nature, 335, 50. Falk, S. W. k Arnett, W. D. (1977). Astrophys. .1. Suppl., 33. 515. Hellings, P. (1983). Astrophys. Space Sci., 9G, 37. HillebrancU, W. k Meyer, F. (1989). Astr. Astrophys.. 219, 1,3. Hsu, J. J. L., Joss, P. C , Ross, R. R. k Podsiadlowski, Ph. (1994). Astrophys. J., (submitted). Joss, P. C , Podsiadlowski, Ph., Hsu, J. J. L. k Rappaport, S. (1988). Nature, 331, 237. Kippenhahn, R. &; Meyer-Hofmeister, E. (1977). Astr. Astrophys., 54, 539. Nomoto, K., Suzuki, T., Shigeyama, T., Kumagai, S., Yamaoka, H. k Saio, H. (1993). Nature, 364, 507. Paczyriski, B. (1976). In 7.4 U Symposium 13, Structure and Evolution of Close Binary Systems, ed. P. P. Eggleton, S. Mitton, k J. Whelan, pp. 75 (Dordrecht: Reidel). Podsiadlowski, Ph., Hsu, J. J. L., Joss, P. C. k Ross, R. R. (1993a). Nature, 364, 509. Podsiadlowski, Ph., Hsu, J. J. L., Joss, P. C. k Ross, R. R. (1993b). This volume. Podsiadlowski, Ph. k Joss, P. C. (1989). Nature, 338, 401. Podsiadlowski, Ph., Joss, P. C. k Hsu, J. J. L. (1992). Astrophys. J., 391, 246 (PJH). Podsiadlowski, Ph., Joss, P. C. k Rappaport, S. (1990). Astr. Astrophys., 227, L9. Rathnasree, N. (1993). Astrophys. J., 411, 848. Ray, A., Singh, K. P., k Sutaria, F. K. (1993). J. Astrophys. Astr., 14, 53. Shigeyama, T., Nomoto, K., Tsujimoto, T. k Hashimoto, M. (1990). Astrophys. J. Lett., 361, L23. Tout, C. A. k Eggleton, P. P. (1988). Astrophys. J., 334, 357. Tutukov, A. V., Yungelson, L. R. k Iben, I. (1992). Astrophys. J., 38G, 197. Vanbeveren, D. (1987). Astr. Astrophys., 182, 207. Wampler, E, J., Wang, L., Baade, D., Banse, K., D'Odorico, S., Gouifles, C. k Tarenghi, M. (1990). Astrophys. J. Lett., 362, L13. Woosley, S. E., Eastman, R.. G., Weaver, T. A. k Pinto, P. A. (1994). Astrophys. J., (in press). Woosley, S. E. k Weaver, T. A. (1985). In Nucleosynthesis and Its Implications On Nuclear and Particle Physics, Proc. 5th Moriaud Astrophys. Cotif., ed. J. Audouze k T. van Thuan, pp. 145 (Dordrecht: Ilciclel).

The Progenitor of SN 1993J and its Mass-Loss History Ph. Podsiadlowski1, J. J. L. Hsu 2 , P. C. Joss 3 and R. R. Ross4 1 2 3 4

Institute of Astronomy, Cambridge CB3 OHA, UK University of California at Berkeley, CA 94720, USA Massachusetts Institute of Technology, Cambridge, MA 02139, USA College of the Holy Cross, Worcester, MA 01610, USA

Supernova 1993J in the spiral galaxy M81 is the brightest supernova since SN 1987A and, like the latter, appears to be another peculiar type II supernova. Its early light curve is characterized by a very sharp initial peak (lasting for less than ten days) followed by a less rapid secondary brightening, which was qualitatively similar to the secondary brightening observed in SN 1987A. Humphreys et al. (1993) have identified a candidate progenitor consistent with the position of the supernova. Combining their UBVR. photometry with the I magnitude obtained by Blakeslee & Tonry (1993), they concluded that the colors of the apparent progenitor require the presence of at least two bright stars. One star is an early-type supergiant (most likely a lateB to early-A supergiant), the other a late-type supergiant (most likely a G to early-K supergiant). The bolometric magnitudes of both stars are in the range of - 6 to - 8 , with best-fit values of - 7 to - 7 . 5 (for an assumed distance of 3.3 Mpc). We have performed our own fits to the photometric data and obtained similar results. These best-fit magnitudes imply mainsequence masses of ~ 15 MQ, but the masses could be as low as 8 MQ or as large as 20 MQ. The image of the candidate progenitor appears extended on some plates (Blakeslee & Tonry 1993). This suggests that, at the distance of M81, the two stars do not form a close binary (although either star could have an undetected binary companion). Neither of the two inferred stars is a theoretically expected progenitor for a. typical type II supernova. However, a G or early-K supergiant can account for the early light curve, provided that it had lost almost all of its hydrogen-rich envelope before the supernova explosion (i.e., provided that it was a "stripped" supergiant). To constrain the properties of the progenitor, we performed a series of hydrodynamical calculations (see Podsiadlowski et al. 1993; Hsu et al. 1994). We found that the visual light curve can be well fitted with the explosion of a progenitor star with a main-sequence mass of 15 MQ, which suffered severe mass loss and had a residual hydrogen-rich envelope mass of ~ 0.2 MQ, and 187

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with an explosion energy of ~ 1051 erg. The secondary peak of the light curve and the subsequent exponential decay can be understood as the result of ~ O.IJV/Q of 56Ni that was produced in the supernova, and confined to the inner region of the ejecta (the exact amount of 5<5Ni depends sensitively on the assumed reddening and the distance to M81). Our explosion model predicts that, because of the low mass of the hydrogen-rich ejecta, SN 1993J should resemble a type Ib supernova at late times (i.e., the supernova spectrum should transform from type II to type Ib). Recent observations of the supernova (e.g., Filippenko and Matheson 1993) seem to have confirmed this prediction. It is unlikely that the progenitor lost most of its envelope in an ordinary stellar wind, since this would require a. much larger initial main-sequence mass than is consistent with the photometric constraints on the progenitor, as well as significant fine-tuning of the wind parameters (otherwise the progenitor would have been either a normal red supergiant or a Wolf-Rayet star). A more probable scenario is that the progenitor was a member of a close binary system and underwent stable case C mass transfer (i.e., as a red supergiant after helium core-burning; see Podsiadlowski, Joss & Hsu 1992). As a result of Roche-lobe overflow, the progenitor will lose most of its hydrogen-rich envelope, of which a significant fraction will be accreted by the companion. However, when the mass in the hydrogen-rich envelope has decreased below a certain critical mass (~ 0.3 MQ for a progenitor of ~ 15 M Q ) , the envelope of the progenitor will shrink significantly; as a result the star will no longer fill its Roche lobe, and mass transfer will cease. Inititially, the primary will still have the appearance of a red supergiant (although of somewhat earlier type), but will have only a low-mass residual hydrogen-rich envelope. The time remaining to the supernova event at this phase is such that a stellar wind might well reduce the envelope mass to ~ 0.1 MQ by the time of the explosion. The resulting progenitor is precisely the type required to explain the early light curve. In Figure 1, we present a typical binary calculation to illustrate stable case C mass transfer leading to the scenario described above. The initial binary consists of a 15 MQ primary and a 10 MQ secondary. We assume that the primary loses 3 MQ before it fills its Roche lobe, possibly in an enhanced stellar wind, in order to avoid dynamical mass transfer (see Podsiadlowski et al. 1992 for a further discussion of this issue). We also assume that mass transfer is non-conservative and that 50% of the mass lost by the primary is lost from the system (the rest is accreted by the secondary). With these assumptions we can follow the evolution of all binary parameters. The top panel of Figure 1 shows the radius, Ri, of the primary and the radius of

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- Roche-lobe radius

o

-4—I—I—I—\—HH—I

I

I

I

I | I

I —

I

I I I I I

I I I — I — — I

I

I—I

log A* (yr) Fig. 1. Illustration of stable case C mass transfer as a function of time since the beginning of mass transfer. Top panel: radius of the primary (solid curve) and its Roche lobe (dashed curve); middle panel: masses of the primary and secondary (solid and dashed curves, respectively) and the mass interior to the hydrogen-rich envelope of the primary (dotted curve); bottom panel: mass-loss rate from the primary.

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its Roche lobe as functions of time since the beginning of the mass-transfer phase. Since the radius of the primary is always less than or equal to the Roche-lobe radius, mass transfer is dynamically stable. The bottom panel of Figure 1 displays the mass-loss rate from the primary (including an assumed stellar-wind mass-loss rate of lO~ 5 M0yr~ 1 ) and shows that there are two distinct mass-transfer phases, separated by a short detached phase. In the first, rapid ("thermal") phase, which has a duration of a few hundred years, the primary's mass-loss rate is ~ 10~2 .A^yr" 1 , and the structure of the primary is far from thermal equilibrium. In the second, slow (evolutiondriven) mass-transfer phase, which lasts for a few times 104yr, the mass-loss rate is of order lO~4M0,yr~1, and the primary is in approximate thermal equilibrium. The mass-transfer rate in the thermal phase is of the same order as the Eddington accretion rate of the secondary (~ 10~2 ;\/(T, yr" 1 ). Hence we expect mass loss from the system to occur during this phase via. two distinct modes: (1) a high-velocity, bipolar mode emanating from the inner accretion disk around the secondary with a characteristic escape velocity of ~ 1000 km s"1, and (2) a low-velocity equatorial mode, possibly through the outer Lagrangian point, L2, with a characteristic velocity of ~ lOkms" 1 . By the time of the supernova, explosion, the mass loss should have created a complex nebula around the progenitor that will resemble, in many respects, the presupernova nebula around SN 1987A (Wampler et al. 1990), with an hour-glass shape and a disk-like structure at its waist. The disk-like structure will be terminated by a massive ring-like torus (ejected during the thermal mass-transfer phase) at a characteristic distance of about a light year. While it will not be possible to resolve this nebula, directly in the optical, the nebula, should be glowing in recombination radiation (resulting from the supernova, ultraviolet burst). This radiation may be detectable with IUE and may allow inferences to be drawn about the geometry of the nebula. Remnants of the hourglass/ring structure may also persist into the epoch when the supernova remnant becomes detectable in the radio. In addition, our model predicts that the circuinstellar material is nitrogen-rich (and perhaps even helium-rich), since part of it, has undergone CNO burning inside the progenitor. Finally, we note that the companion star, which is presently hidden below the photosphere of the supernova, should eventually reappear. If the supernova led to the formation of a pulsar, the detection of periodic modulation of its pulse period would directly confirm our proposed binary scenario.

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References Blakeslee, J. k. Tonry, J. (1993). IAU Circular No. 5758. Filippenko, A. V. & Matheson, T. (1993). IAU Circular No. 5787. Hsu, J. J. L., Podsiadlowski, Ph., Joss, P. C. & Ross, R. R. (1994). (in preparation). Humphreys, R. M., Aldering, G. S., Bryia, C. & Tliunnes, P. (1993) IAU Circular No. 5739. Podsiadlowski, Ph., Hsu, J. J. L., Joss, P. C. & Ross, R. R. (1993). Nature, 364, 509. Podsiadlowski, Ph., Joss, P. C. & Hsu, J. J. L. (1992). Astrophys. J., 391, 246. Wampler, E. J. et a/. (1990). Astrophys. J., 362, L13.

Narrow optical emission lines from supernova 1993J Robert J. Cumming 1 ' 2 , Peter Meikle2'1, Nic Walton 3 and Peter Lundqvist4 Royal Greenwich Observatory, Madingley Road, Cambridge CB3 OEZ, U.K. Blackett Laboratory, Imperial College of Science, Technology and Medicine, Prince Consort Road, London SW7 2BZ, U.K. 3 Isaac Newton Group, Royal Greenwich Observatory, Apartado 321, 38780 Santa Cruz de La Palma, The Canary Islands, Spain 4 Stockholm Observatory, S-133 36 Saltsjobaden. Sweden

Abstract We report narrow emission lines observed during the first 10 days of supernova 1993J. The earliest spectra showed resolved, P-Cygni-like Ha emission which declined on a timescale of about 2 days. Fast-declining, unresolved He II and coronal iron lines were also detected. A higherresolution spectrum taken on day 8 after the explosion showed that the Ha line had narrowed to about 90 km s"1 FVVHM and had lost its P-Cygni profile. The narrow line emission followed ionisation, by the EUV flash, of dense circumstellar material close to the supernova. A likely explanation for the rapid disappearance of the these lines was that the circumstellar gas was overrun by the expanding supernova shock. However, it may be necessary to also invoke the decline of trapped UV radiation to account for the rapid fading of the H« line.

1 Introduction The detection of narrow emission lines from supernova 1993.1 gives us a rare opportunity to study the circumstellar medium produced in the preexplosion phase. This can give insights into the evolution and mass loss history of the progenitor star, and test theories of the interaction of a supernova with its surroundings. 2 Observations Optical spectra of SN 1993J were obtained with the IDS and ISIS spectrographs on, respectively, the INT and WHT on La Palma. For the first two weeks, spectroscopy was carried out nightly, and less frequently thereafter. The resolution and wavelength coverage varied according to the scheduled observing programme being carried out at the time. The first spectrum was taken on 1993 March 29.88 by Enrique Perez and Derek Jones (Gomez, Lopez & Perez 1993), just over two days after the supernova explosion date, which we take to be March 27.5. The spectra were reduced and 192

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Table 1. Circumstellar line intensities in 10~lA erg s~x cm~2; Upper limits are 3a. days after Mar 27.5 2.38 3.38 4.56 5.40 8.40 13.43 66.39 79.39 89.38

A0(A) Mar Mar Apr Apr Apr Apr Jun Jim Jun

29.88 30.88 01.06 01.90 04.90 09.93 01.89 14.89 24.88

Ha 6563

He II 4686

[FeX] 6374

[Fe XI] 7892

6.5±0.8 3.8±0.5 1.3±0.3

—

5.9±2.0

3.7±0.4 3.0±0.4

<0.9 <1.4 — — — — —

<1.5 <0.7 — — — — —

1.0±0.3 0.6±0.2 <1.2 <2.3

0.9±0.3 0.38±0.04 <1 <0.05 <0.07 <0.08

— — — — —

flux-calibrated by Jim Lewis (RGO) using IRAF. These observations are described in detail in Lewis et al. (1994), and are available from the RGO on-line archive (Martin & Lewis, 1993). 3 Results Our spectra revealed narrow line emission which persisted to at least day 8.40. Narrow optical lines were also detected by other observers during this time (Andrillat 1993a,b; Filippenko et al. 1993; Porter et al 1993; Benetti et al. 1993). All the narrow lines declined rapidly during the first week and by day 17 the lines were reported to have vanished (Filippenko & Matheson 1993). After day 8.40, our next high-resolution spectrum was not taken until day 66.39, when we obtained for Ha a 3
3.1 The Ha profile In Figure 1 we show the evolution of the Ha line profile. On days 2.38 and 3.38 the line was resolved, and its shape suggests a P-Cygni-like profile. The blueshifted absorption extended to -800 km s" 1 , while the red wing reached about +1000 km s" 1 . Porter et al. (1993) report a similar profile. With respect to the local standard of rest (LSR) in M81 of-135 km s" 1 (Vladilo

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-2000

velocity (km s"1)

2000

Fig. 1. Evolution of the Ha profile. The spectra are shown at the same flux scale, but have been shifted vertically. et al. 1993), the Ho line peak showed a redshift of about +100 km s on day 2.38, falling to about +50 km s" 1 on day 3.38. The day 4.56 profile is of poor signal-to-noise ratio, but the high resolution spectrum obtained on day 8.40 shows the line still resolved, with a FWHM of only about 90 km s" 1 . The blueshifted absorption had disappeared, and the line profile had become symmetrical, centred on a velocity equal to that of the M81 LSR. Widths in the range 200-500 km s" 1 have been reported for day 3 (Andrillat 1993a; Benetti et al. 1993; Porter et al. 1993). The wide range of values was probably a result of the P-Cygni profile being observed at different sensitivities. On day 4.4 Andrillat (1993b) reported a width of only 145 km s" 1 which, taken with our day 8.40 width measurement, suggests that the broad component of the line disappeared rapidly during the first 10 days.

3.2 Profiles of other lines The spectral region containing He II A4686 was not covered until day 3.38, when we obtained a 3
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penko et al. (1993) on day 2.8 and by Benetti et al. (1993) on day 3.5. At a resolution of 380 km s" 1 the line was unresolved in our data. Due to the rapid fading of the line and the lower resolution employed on the following night, this line was not detected again. [Fe XJA6374 and [Fe XI]A7892 were detected on days 2.38 and 3.38, but not thereafter. At a resolution of about 250 km s" 1 neither line was convincingly resolved. When the helium and iron lines were of reasonable signal-to-noise ratio, we were able to examine their wavelength shifts. The peaks of the iron lines on day 2.38 and the helium line on day 3.38 were all redshifted by about 180 km s" 1 with respect to the M81 LSR.

4 Origin of the Ho emission Richmond (1993) reported a lack of any II II region emission within 150 pc of the progenitor. Moreover, the observed narrow lines are not typical of H II regions, which do not show strong coronal lines, and generally have strong low-excitation forbidden lines which are not observed here. Together with the rapid decline of the lines after the explosion, this points to the supernova as the ultimate cause of the emission. The presence of the lines less than 3 days after the explosion, and the speed with which they declined indicate that the gas probably lay close to the supernova, <1 light day, rather than being due to recombination in a cloud lying at a. large distance from the explosion. The most likely triggering mechanism was photoionisation of the circumstellar medium (CSM) by the EUV burst from the shock breakout; ionisation by X-rays from ejectP./CSM interaction may also have been important. The P-Cygni profile of the earliest Ha observations suggests that the emitting material was dense, and lay just beyond the expanding fireball. For the typically fastest ejecta velocities seen in supernovae this again places the the emitting circumstellar material at a distance of less than a light day from the supernova. Some of this CSM material must have been moving as fast as 1000 km s" 1 with respect to (and presumably outwards from) the supernova centre of mass. The movement to the blue of the line peak between days 2.38 and 8.40, the disappearance of the absorption trough, and the development of a symmetrical profile all point to the emission arising from decreasingly dense material. If the supernova centre of mass was roughly stationary with respect to the LSR, then the coincidence by day 8.-10 of the emission line peak blueshift with that of the LSR tends to confirm that the material lay close to the supernova. Light travel-time effects in a larger-scale expanding

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CSM would have produced a net blueshift with respect to the supernova centre of mass at this time. Why did the lines decline so quickly? After the gas had been heated and ionised by the EUV burst, it would have recombined with a. timescale which depends on the density. If the line emission was due to recombination, then to produce the observed fast decline, the density must have been 2 X 1012 cm" 3 . This is improbably high for circumstellar material. A possible explanation for the rapid decline is that the emitting CSM was not unusually dense, and did indeed produce the Ha emission by recombination, but in addition, the recombining gas was overrun by the expanding supernova shock (Filippenko et al. 1993). After being shocked, the gas would have a temperature of about JO9 K, too high to produce optical lines, and with velocities too fast to produce narrow lines. As already pointed out, the evolution of the Ha line indicates emission from decreasingly dense material. This is as one would expect as the shock moved outwards through a CSM originally created by, for example, a steady mass loss rate. To account for the narrowing of the line, we note that there would be prea.ccelera.tion of the circumstellar gas close to the supernova, by the EUV outburst (and to some extent by the early emission from the ejecta/CSM shock). After about a day the unshocked CSM would have a steep velocity gradient, so that the highest velocity CSM gas, close to the shock, would be quickly overtaken by the blast wave thus producing the observed narrowing of the line. To test this scenario, we have used a model in which the CSM is spherically symmetric and ionised out. to a. large radius by the EUV burst. We assume that in the subsequent recombination, the Ho emissivity is proportional to density squared, and that the shock radius increases solf-similarly according to the minishell model (Chevalier 1982). Given the VLBI-determined shock radius of SN 1993J (Marcaide et al. 1993), we expect the Ha luminosity to decline as t~08. This is much slower than the observed luminosity, which declined as r ( 2 3 ± 0 - 2 > . However, there may be an additional driving mechanism for the Ha emission besides recombination. The presence of a P-Cygni profile in Ha suggests that the CSM was fairly dense, and so was probably optically thick to SN continuum radiation of wavelength around Ly/3. Thus an enhanced ?i=3 population would be established in the CSM, amplifying the Ha emission. Since the intensity of this photo-excited emission is tied to the energy density of the continuum around Ly/J, and since the flux in this part of the SN spectrum declined very rapidly, this may, at least in part, account for the fast fading of the Ha line. It may also account, for the narrowing of the line, since the enhanced Ha emission would be strongest where the Ly/3 trapping

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of continuum radiation was most effective. This would be predominantly in the CSM region of steep velocity gradient near the SN photosphere. An interesting alternative to the above is pure recombination in a disk geometry. Keeping an r~2 density distribution in the disk, we find that as the shock overruns the recombining material the line flux fades as i~ 1 6 . This is closer to the observed time dependence, and did not require us to invoke the photoexcitation mechanism. What is the mass of hydrogen responsible for the narrow line emission? From the integrated Ha luminosity we obtain 0.02 M 0 , assuming one Ho photon per recombination, under case B, taking M///M to t a (=0.9 and a distance of 3.6 Mpc. However, the mass required may be less than this if the photoexcitation mechanism is important. Our non-detection of narrow Ha 2-3 months after the explosion implies an upper limit of about 5 x 1035 erg s" 1 in this line. At this level, the most luminous CSM line seen in SN 1987A, at its maximum intensity would have been just detectable at the distance of M81. The SN 1987A CSM lines originated in a ring of about 1.4 light yr across. If a similar ring surrounds SN 1993J, it is possible that when the EUV flash intercepts it, we will detect flaring of the narrow Ha emission during the next 1-2 yr.

Acknowledgements We thank Claes Fransson, Patrick Petitjean and George Sonneborn for helpful discussions. The William Herschel and Isaac Newton Telescopes are operated on the island of La Palma by the Royal Greenwich Observatory in the Spanish Observatorio del Roque de los Muchaclios of the Instituto de Astrofisica de Can arias.

References Andrillat, Y. (1993a). IAU Circular No. 5736. Andrillat, Y. (1993b). IAU Circular No. 5743. Benetti, S., Contarini, G., Gratton, R. & TuraUo, M. (1993). IAU Circular No. 5751. Chevalier, R. A. (1982). Astrophys. J., 258, 790. Filippenko, A. V. et. al. (1993). IAU Circular No. 5740. Filippenko, A. V. & Matheson, T. (1993). IAU Circular No. 5760. Lewis, J. R. et al. (1994). Mon. Not. R. astr. Soc., 2CC. L27. Marcaide, J. M. et al. (1993). IAU Circular No. 5785. Martin, R. & Lewis, J. (1993). Gemini, No. 40, 14. Perez, E., Gomez, G. & Lopez, R. (1993). IAU Circular No. 5733. Porter, A. C , Wells, L. A. & Rubin, V. C. (1993). IAU Circular No. 5748. Richmond, M. (1993). IAU Circular No. 5739. Vladilo, G., et al. (1993). Astron. Astrophys., 280, L l l .

UV Spectroscopy of SN 1993J and Detection of Highly-Ionized Gas Close to the Progenitor George Sonneborn 1 , Pedro Rodriguez Pascual 2 , Willem Wamsteker 2 and Claes Fransson 3 1

Laboratory for Astronomy and Solar Physics, Code 681, NASA/Goddard Space Flight Center, Greenbelt, MD 20771, U.S.A. 2 IUE Observatory, ESA-VILSPA, Casilla 50727, E-28080 Madrid, Spain 3 Stockholm Observatory, S-133 36 Saltsjobaden, Sweden

Supernova 1993J in M81 (NGC 3031) was discovered by Spanish amateur astronomers on 28.86 March 1993 (Ripero & Garcia 1993). The first IUE spectra were taken on 30.2 March at VILSPA a few hours after notification of the discovery (Wamsteker et al. 1993) and the supernova was regularly observed by IUE over the next several weeks. This paper summarizes the principal results of the IUE observations (see Fransson & Sonneborn 1994 and Sonneborn et. al. 1994 for observational details and more extensive discussion). The first photographic detection of the supernova was on 28.30 March at magnitude 13.6 (Merlin & Neely 1993). Modelling of the supernova V light curve indicates that the explosion occurred on 27.8 March and that shock breakout should have occurred at ~ 28.0 March (Shigeyama et al. 1994). Careful analysis of pre-outburst plates and images has identified the progenitor and shown that its colors are consistent with a late-type supergiant (cf. IAU Circular No. 5739) and the supernova's Type II classification. In Fig. 1 we show the first UV spectra of SN 1993.1 from 30.2 March to 3.5 April, where the very rapid cooling of the exploding photosphere is readily apparent. On 30.2 and 31.2 March the temperature was ~22,500K and ~14,500K, respectively. Here we have assumed a Galactic extinction law and EB-V — 0.18, as determined from a fit of the 2200 A dust feature on 30.2 March. Such a rapid cooling is expected just after shock breakout, and agrees quite well with the light curve calculations by Shigeyama et al. (1994). The strong UV continuum faded by a factor of ~ 200 by April 4 (A <1600A); smaller decreases occurred at longer wavelengths. The very rapid evolution was also apparent in the optical spectrum (cf. Meikle et al. 1993; Filippenko, Matheson, & Ho 1993). The only UV emission feature detected is the N V doublet at 1240A. It is very strong in the first spectrum, but, weakens quickly with time. The N V flux (Fig. 2) decreased by a factor ~ 43 between 30.2 March and 4.5 198

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10,-11

SN 1993J March 30.2 UT

1500

2000 2500 Wavelength(A)

3000

Fig. 1. Spectral evolution in the UV of SN 1993.1 during the first week after the explosion. The N V A1240 emission line is clearly visible. April (Sonneborn et al. 1993). By 22 April the N V emission was virtually absent. The N V luminosity on 30.2 March was ~ 2 X 1O40 erg s" 1 , assuming a distance of 3.5 Mpc and EB-V = 0.18. In high resolution SVVP spectra obtained on 30 and 31 March the N V lines show the presence of a dense circumstellar shell (see Fig. 3). The blue component (\\^ - 1242.80A) has a red wing extending to at least ~ 200 km s" 1 , while the red component (Aiab = 1238.82A) has a FWHM of only ~ 35 km s" 1 . From the narrow line widths it is clear that they can not arise in the supernova ejecta, but must originate in the circumstellar medium of the progenitor system. The relative fluxes and line profiles are, however, different from those expected. Under optically thin conditions the A1238.8 component is expected to be a factor two brighter than the A1242.8 component. FA'en if optically thick, the ratio is never less than one. The most natural origin for the circumstellar medium is the wind from the progenitor (or possibly a companion star). Tf the expansion velocity of the circumstellar medium, VCSM, where the N V line arises is £ 1000 km s" 1 then one expects a P-Cygni profile with an absorption component extending to -VCSM a " d an emission component with a lower extension to the red due

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10-10. (Dm

1 1

SN 1993J

• a

S> 10-12 X

I 10"13 CD

|

10"

• NV 1240 net flux a Continuum near 1240A

14

-15 10'

0

5 10 15 Days After 1993 March 28.0 UT

20

Fig. 2. Light curve of the N V A 1240 line. Solid circles: N V flux above the continuum; open squares: mean continuum level near 1240A.

to the occultation by the ejecta. If VCSM ~ 963 km s" 1 , the velocity separation of the N V components, the A1242.8 blueward absorption will interfere with the A1238.8 emission component. The scattered A1238.8 photons will emerge in the A1242.8 component, explaining the higher luminosity of this line. Although the 31.5 March high resolution spectrum is noisier (decreasing N V flux) there is some evidence for a decrease in the width of the A1242.8 line, as well as a more abnormal A1238.8/A1242.8 ratio. Interstellar absorption by gas in M 81 and the Galaxy is not a plausible explanation for the line asymmetries because interstellar N V lines seldom have optical depths larger than ~ 0.4 along sightlines out of the halo (e.g., Sembach & Savage 1992) and the moderate strength of the Si IV and C IV interstellar lines in the IUE high-resolution spectra of SN 1993J (de Boer et al. 1993) are consistent with negligible N V interstellar absorption. Velocities ~ 1000 km s" 1 in the circumstellar medium of a late-type luminous star are much higher than expected, since wind velocities are typically 10 - 30 km s" 1 . An explanation, proposed by Fransson, Lundqvist & Chevalier (1994, hereafter FLC94), is that radiative acceleration of the

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SN 1993J

0 March 30.4 UT

-600

-400

-200 0 200 Heliocentric Velocity (km/sec)

400

600

Fig. 3. The N V AA 1238.82- 1242.80 line profiles on 30.4 March.

progenitor's circumstellar gas by the very strong burst of UV and soft X-ray photons during the first days is responsible for the high velocity. The interaction of the supernova's ejecta (v ^ 19,000 km s" 1 ) with the circumstellar medium will give rise to a shocked region, bounded by one ingoing (in a Lagrangian sense) reverse shock wave and one outgoing circumstellar shock. These shocks are expected to have temperatures of order T = 107 — 109K. The high N V A 1240 luminosity shows that there was a large flux of photons above 77 eV during the first few days. The models by FLC94 show that the burst in connection with the shock breakout indeed ionizes the circumstellar gas nearly completely. Because of its high density, the stellar wind material inside ~ 1015 cm has time to recoinbine before the shock hits the gas, explaining the presence of the N V emission line and its rapid decay. The UV (e.g. this paper) and optical (e.g. dimming et al. , this volume) were the first wavelength ranges to reveal a significant circumstellar medium close to SN 1993J. The supernova was subsequently also detected in the radio (Weiler et al. 1993; Pooley & Green 1993) and X-ray (Zimmermann et al. 1993; Tanaka 1993) regions. The X-ray spectrum is consistent with thermal emission from the high-temperature shocks. The radio emission showed the same characteristic pattern as has been observed in other Type II and

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Type Ib supernovae. Emission was initially seen at short wavelengths and at successively longer wavelengths at later times. This behavior has been explained for other supernovae as a result of decreasing free-free absorption by the ionized gas in a circumstellar medium around the supernova (Chevalier 1982). As the supernova expands, the emission measure from the radio emitting region close to the shock through the circumstellar medium decreases, explaining the radio turn-on.

References Chevalier, R. A., 1982, ApJ 259, 302 de Boer, K., Rodriguez Pascual, P., Wamsteker, W., Sonneborn, G., Fransson, C , Bomans, D. J., and Kirshner, R. P. 1993, A&A 280, L15 Filippenko, A. V., Matheson, T., & Ho, L. C , 1993, ApJ 415, L103 Fransson, C , Lundqvist, P., & Chevalier, R. A. 1994, in preparation (FLC94) Fransson, C , & Sonneborn, G. 1994, in Frontiers of Space and Ground-based Astronomy, eds. W. Wamsteker, M. Longair, and Y. Kondo, Kluwer Acad. Publ., in press Meikle, P., et al. , 1993, Gemini, 40, 8 Merlin, J.-C. & Neely, A., 1993, IAU Circ. No. 5740 Pooley, G. G., & Green, D. A., 1993, MN 264. L17 Ripero, J., fc Garcia, F., 1993, IAU Circ. No. 5731 Sembach, K. P. & Savage, B. D. 1992, Ap.ISupp 83, 147 Shigeyama, T., Suzuki, T., Kumagai, S., Nomoto K., Saio, H., & Yamaoka, H., 1994, ApJ 420, 341 Sonneborn, G., Rodriguez, P. M., Wamsteker, W., Fransson, C , & Kirshner, R., 1993, IAU Circ. No. 5754 Sonneborn, G., et al. 1994, in preparation. Tanaka, Y., & the ASCA team 1993, IAU Circ. No. 5753 Wamsteker, W., Rodriguez, P. M., Gonzalez, R., Sonneborn, G., fc Kirshner, R., 1993, IAU Circ. No. 5738 ; Weiler, K. W., Sramek, R. A., Van Dyk, S. D., & Panagia, N., 1993, IAU Circ. No. 5752 Zimmermann, H. U., Lewin, W., Magnier, E., Predehl, P., Hasinger, G., Pietsch, W., Aschenbach, B., Trumper, J., Fabbiano, G., van Paradijs, J., Lubin, L., k. Petre, R., 1993, IAU Circ. No. 5748

Ryle Telescope observations of the radio emission from SN 1993J D. A. Green and G. G. Pooley Milliard Radio Astronomy Observatory, Cavendish Laboratory, Madingley Road, Cambridge CB3 OHE, U.K.

Abstract We present observations of SN 1993.1 made with the Ryle Telescope, Cambridge, UK, at 15.25 GHz. These show a sharp switch-on of the radio emission about eight days after the supernova explosion, with the emission then brightening at an approximately constant rate for the next month. The emission peaked about 75 days after the explosion, and then showed a gradual decline, with variations on the timescale of weeks. The long, steady rise in emission and the sharp switch-on do not fit with expected radio emission from 'mini-shell' models of radio supernovae.

1 Introduction The detection of SN 1993J in NGC3031 (=M81) by F. Garcia on March 28.86 was reported in an IAU Circular by R.ipero (1993) on March 30. This was the nearest supernova (SN) detected in the northern hemisphere since SN 1937C, and it has been, and will be, studied at many wavelengths in more detail than any other SN except SN 1987A in the LMC. SN 1993J showed hydrogen lines in its early spectra, and was initially thought to be a type II SN, but subsequently its spectra developed to resemble those of type Ib SN. Thus it has been classed as a peculiar type lib SN (Filippenko & Matheson 1993), and it is thought to be the result of the explosion of a massive star which has almost lost its outer hydrogen-rich envelope. SN 1993J was monitored from March 30 — only three days after the explosion, taken to be March 27.5 (Wheeler et ul. 1993) — at 15.25 GHz with the Ryle Telescope. This led to the first reported detection of radio emission from SN 1993J, on April 5 (Pooley & Green 1993a), although subsequent analysis revealed earlier radio detections both with the Ryle Telescope on April 4 (see below) and with the VLA at 22.5 GHz (VVeiler et al 1993) on April 2. Here we present observations of this radio supernova. (RSN) made up to July 1993 (see also Pooley & Green 1993b, who give more details of the observations and their calibration). This is the first RSN for which the rising emission has been studied in detail. 203

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D. Green & G. Pooley: Ryle Telescope observations of SN 199SJ

40

SO

80

d«T (from I9S3 March 27.5)

Fig. 1. The flux density of SN 1993J at 15.25 GHz.

2 Observations and Results The Ryle Telescope — the upgraded 5-km Telescope (Ryle 1972) — is an Earth-rotation synthesis instrument with 8 antennas on an east-west baseline. The observations reported here are of Stokes parameter I + Q, and were centred on 15.25 GHz with a bandwidth of 0.28 GHz. Most of the observations of the SN were tracks of less than 12 hours obtained during intervals between longer observations of other targets. The nucleus of M81 is a strong radio source which lies ss 3 arcmin from SN 1993.L so we used the longer baselines to resolve the SN from the nucleus. The observations were calibrated using short observations of 130954+658, which were interleaved with the observations of the SN, plus longer observations of 3C48 made on most days to define the flux density scale. 3C48 was assumed to have a flux density of 1.7 Jy, which is consistent, with the scale of Baars et al. (1977). The calibrated observations of SN 1993J are shown in Fig. 1. In addition to the data presented in Fig. 1. SN 1993J was observed, but was undetected at < 0.5 mJy, on five earlier days before April 4. The early part of the evolution of the radio emission from SN 1993J, together with the upper limits, are shown in Fig. 2.

3 Discussion The preferred model (e.g. Weiler et. al. 1986) for previously observed RSN is the simple 'mini-shell' model (Chevalier 1982), where non-thermal radio

D. Green & G. Pooley: Ryle Telescope observations of SN 1993J

10

205

IB

daj (from 1M3 lurch 27.5)

Fig. 2. The flux density of SN 1993.7 at. 15.25 GHz, and upper limits from early observations.

emission is generated by the interaction of the SN with the surrounding circumstellar material, which also provides a low-frequency turnover to the radio spectra due to free-free absorption. The flux density, S\ of the radio emission at a given frequency varies with time from the SN explosion, tf, as 5 a / / 3 e~ r , with the optical depth, r. varying as r <x ts. So the radio emission rises rapidly as the optical depth decreases, tending to a, simple S oc tl3 behaviour at later times, when the absorption becomes unimportant. Fig. 3 shows the time variation expected for this model for parameters (3 and 6 typical of type II and type Ib RSN. There are several discrepancies between our observations of SN 1993J at 15.25 GHz and the basic 'mini-shell' model. The switch-on of the emission appears to be sharp, whereas the 'mini-shell' model shows a more gradual turn-on. The steady rise in emission lasts for several times the delay between the explosion and the radio switch-on, whereas in the model the duration of this steady rise is comparable with the time delay. There are variations on timescales of weeks from day 60 onwards from the underlying turnover and fading of the emission. The "mini-sheir model can be made to fit the data reasonably well only if the time origin is taken to be many days before the SN explosion date. Lundqvist (these proceedings) presents a modified model for the emission which does not require the time origin to be regarded as a free parameter, and this does fit the data better. This model is based on a circumstellar medium density varying as r" 1 ' 5 rather than the r~2 expected

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D. Green & G. Pooley: Ryle Telescope observations of SN 1993J

1.6 time (arbitrary unita)

Fig. 3. Plot of expected radio emission from a supernova using typical parameters from previously observed RSN and the simple 'mini-shell' model. The solid curve is for type II RSN with (3 — —0.7 and b — —2.4, and the broken curve for type Ib RSN with f3 — —1.4 and 6 — —2.6. for a constant velocity wind in the standard 'mini-shell' model. However, the observed sharp switch-on may still be a problem for this revised model. References Baars, J.W.M., Genzel. R., Pauliny-Toth, U.K. & Wil.zel, A. (1977). Aslron. Astrophys., 61, 99. Chevalier, R.A. (1982). Astrophys. J., 259, 302. Filippenko, A.V. fc Matheson, T. (1993). IAU Circular No. 5787. Pooley, G.G. & Green, D.A. (1993a). IAU Circular No. 5751. Pooley G.G. & Green D.A. (1993b). Mon. Not. R. ash: Soc, 2C4, L17. Ripero J. (1993). IAU Circular No. 5731. Ryle M. (1972). Nature, 239, 435. Weiler K.W., Sramek R.A., Panagia N.. van der Hulst J.M. k Salvati M. (1986). Astrophys. J., 301, 790. Weiler K.W., Sramek R.A., Van Dyk S.D. & Panagia N. (1993). IAU Circular No. 5752. Wheeler J.C. ef al. (1993). Astrophys. J., 417, L71.

The Early Radio Emission from SN 1993J Kurt W. Weiler1, Schuyler D. Van Dyk1-2, Richard A. Sramek 3 , Nino Panagia 4 ' 5 and Michael P. Rupen 3 1 Remote Sensing Division, Code 7215, Naval Research Laboratory, Washington, DC, 20375-5351, USA 2 Naval Research Laboratory/NRC Cooperative Research Associate 3 National Radio Astronomy Observatory, P. 0. Box 0, Socorro, NM 87801, USA 4 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA; affiliated with the Astrophysics Division, Space Science Department of ESA 5 University of Catania, Italy

Abstract After initial upper limits on the 3.6 and 2 cm radio emission from SN 1993J established by the Very Large Array (VLA) on 1993 March 31 only three days after optical discovery, the supernova was first detected in the radio range with a flux density of 0.8 mJy at 1.3 cm wavelength on 1993 April 2. This makes it the earliest epoch and highest frequency that any Type II supernova has ever been detected in the radio. Since that time, regular monitoring has been done with the VLA at 1.3, 2, 3.6, 6, and 20 cm wavelengths to obtain the most, detailed, multifrequency radio light curves ever established for any supernova. First analysis of this initial data set reveals that while the evolution of the radio emission from SN 1993J is regular in both time and frequency, it is not well described by the previously successful modified Chevalier model. A model in which the external and internal absorption obey the same power law evolution with time and with a slower temporal decline rate than predicted by the Chevalier model gives a good description of the data.

1 Introduction SN 1993J in M81 (NGC3031) was discovered by Francisco Garcia. Diez of Lugo, Spain at magnitude V = l l ^ S on 1993 March 28.91 (R.ipero 1993). By 1993 April 30 when it reached maximum apparent brightness of V = 10!71?, SN 1993J had become the brightest supernova. (SN) in the northern hemisphere in almost 60 years (since SN 1937C in IC4182). Shock breakout has been estimated to have occurred, to within a few hours, on 1993 March 28.0 (Wheeler et al. 1993) and early optical spectra showing lines of ionized hydrogen indicated that it was a Type II supernova (Garnavich & Hong 1993; Filippenko 1993). Due to its proximity (3.6 Mpc; Freed man et al. 1994) and the fact that 207

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K. W. Weiler et al: SN 1995'J - early radio emission

all Type II supernovae (SNe) are expected to be radio emitters (Weiler et al. 1989), we made very early attempts to detect radio emission from SN 1993J. Observations with the Very Large Array (VLA) on 1993 March 31, only three days after discovery, at 3.6 and 20 cm yielded no detection (Sramek et al. 1993), but two days later, on 1993 April 2, the first detection of radio emission from the SN was obtained with the VLA of 0.8 mJy at 1.3 cm wavelength (Weiler et al. 1993). This was confirmed by the detection of 2 cm emission by Pooley and Green (1993a) with the Ryle Telescope in Cambridge, UK on 1993 April 5 and by the VLA at 1.3, 2, and 3.6 cm on 1993 April 8 (Van Dyk et al. 1993a). Such early detections at such high frequencies have never previously been obtained for any radio supernova (RSN). Since these first detections we have been regularly monitoring the radio flux density from SN 1993J at 1.3, 2, 3.6, 6, and 20 cm with the VLA and Pooley and Green (1993b) have been observing frequently at 2 cm with the Ryle Telescope. Additional radio observations at 3 mm with the IRAM (Radford et al. 1993) and Owens Valley (Phillips & Kiilkarni 1993a,b) millimeter interferometers, at 6 cm with the Westerbork Synthesis Radio Telescope in the Netherlands (Strom et al. 1993), and at 0.9 cm with the Effelsberg 100 m telescope of the Max Planck Institute for Radioastronomy in Bonn, Germany (W. Reich, private communication) have been made.

2 Results In Figure 1 we plot the time evolution of the flux density of SN 1993J for the five radio bands measured with the VLA. The solid lines are the best fit model light curves discussed below.

3 Discussion The "mini-shell" model of Chevalier (1981, 1982, 1984) has been shown by Weiler et al. (1986), with the modification of Weiler et al. (1990), to adequately describe all previously known radio supernovae. This model invokes the external generation of relativistic electrons and enhanced magnetic field necessary for synchrotron radiation by the shock wave from the SN explosion interacting with a relatively high-density envelope of matter surrounding the presupernova star. This cocoon is presumed to arise from mass loss in a dense stellar wind from the SN precursor or its companion. As shown by Weiler et al. (1986, 1990), and Van Dyk et al. (1993b), the

K. W. Weiler et al.: SN 1993J - early radio emission

209

1.5

1.2

1.4 1.6 1.8 log (Days Since 1993 Mar 25)

2.2

Fig. 1. The radio light curves for SN 1993J at 1.3 cm (triangles), 2 cm (squares), 3.6 cm (pentagons), 6 cm (stars), and 20 cm (stars).

generalized formulation for the model flux density, 5, the external absorption, r, and the internal absorption, r', for a RSN can be written as:

where r =

-2.1 ±-\ 5 GHz7 VI day/

(2)

and T'

= A',

(3)

,5 GHz; Vl day/ ' The scaling parameters A'i, A'2, and A'3 formally correspond to the flux density, external absorption, and internal absorption, respectively, at 5 GHz one day after the date of explosion, to. The absorption is assumed to be

210

K. W. Weiler et al: SN 1993J - early radio emission Table 1. Model Parameters for SN 1993J Kx

a

0

K2

K3

6 = 6'

t0

7.8 x 102

= -0.8

-0.3

1.1 x 103

1.8 x 104

-2.0

27 Mar '93

purely thermal, ionized hydrogen with frequency dependence v 2 1 and time dependence 8 and 8' for the external and internal absorption, respectively. The emission from the RSN is assumed to be non-thermal, synchrotron radiation having spectral index a and decreasing with time index ft. If the Chevalier model with its assumption that the dense, external cocoon (density p oc r~2) is established by a constant mass loss rate (A/), constant velocity wind (w) from a red supergiant progenitor is accepted, the relation 8 = a —ft— 3 should hold. Also, Weiler et al. (1990) have shown that under the assumption of an outgoing shock/reverse shock where the density p oc r~3 describes the mixed, internal absorbing/non-thermal emitting interaction region between the shocks, 6' = 58/3 should hold. These assumptions have been successful in describing the gross properties of the radio emission for all previously known RSNe. Unfortunately, this modified Chevalier model does NOT describe the radio emission from SN 1993J. No fitting of parameters to Equations 1, 2, and 3 with the above constraints on 8 and 8' produces a good description of the available radio data. In particular, the very sharp early turn-on of the 1.3 and 2 cm emission and the flat, slow decline of the more recent data are not well reproduced. The best parameter fits also produce an unrealistic explosion date, t0, as much as 15 days before the relatively well known shock breakout date of 1993 March 28. The requirements of the Chevalier model of8 = a-ft-3 and 6' = 58/3 must therefore be relaxed. When this is done, the best fits to the data are obtained with 6' « 8 and 8 / a —ft— Z. If the minimum xled fitting process is applied with only the constraints of 8 = 8' and a = -0.8 (Most of the available measurements are still optically thick so that the value of a is poorly determined and has little impact on the fit; we chose a = —0.8 as a typical value for Type II SNe.), a reasonable fit is obtained and yield the parameter values listed in Table 1. The model curves from these values are shown as the solid lines in Figure 1.

K. W. Weiler et al.: SN 19923 - early radio emission

211

4 Conclusions From detailed, early observations of SN 1993J at multiple radio wavelengths, we have established that previously successful models provide a poor description. To determine whether this is an intrinsic difference, or is due to the much better data set at much earlier times and shorter wavelengths available for SN 1993J, must await further measurement. However, it is clear that the presently available data deviates from the Chevalier model in two ways: (a) 6' ^ 58/3 and (b) 6 ^ a — (3 - 3 . Although interpretation is still preliminary, the first difference (a) suggests that the thermal absorbing material is not both internal and external as previously proposed, but may be entirely external with a clumpy or filamentary structure and a range of optical depths. The second difference (b) suggests that the assumed p oc r~2 dependence of the circumstellar density in the Chevalier model is not applicable and a p oc r~ 15 dependence is more likely (for an undecelerated shock). Such a density profile may imply a changing M/w (stellar mass loss rate to stellar wind velocity) ratio or an unusual geometry for the circumstellar material. These modelling consequences are explored in Van Dyk et al. (1993c).

References Chevalier, R.A. (1981). Astrophys. J., 251, 259. Chevalier, R.A. (1982). Astrophys. J., 259, 302. Chevalier, R.A. (1984). Ann. N. Y. Acad. Sci., 422, 215. Freedman, W.L. et. al. (1994). Astrophys. J., in press. Filippenko, A.V. (1993). IAU Circular No. 5731. Garnavich, P. & Hong, B.A. (1993). IAU Circular No. 5731. Phillips, J.A. & Kulkarni, S.R. (1993a). IAU Circular No. 5763. Phillips, J.A. & Kulkarni, S.R. (1993b). IAU Circular No. 5775. Pooley, G.G. & Green, D.A. (1993a). IAU Circular No. 5751. Pooley, G.G. & Green, D.A. (1993b). Mon. Not. R. astr. Soc, 264, LI7. Radford, S., Neri, R., Guilloteau, S. & Downes, D. (1993). IAU Circular No. 5768. Ripero, J. (1993). IAU Circular No. 5731. Sramek, R.A., Van Dyk, S.D., Weiler, K.W. & Panagia, N. (1993). IAU Circular No. 5743. Strom, R.G., Boonstra, A.J., Braun, R., de Bruyn, A.G., L Foley, A.R. (1993). IAU Circular No. 5762. Van Dyk, S.D., Weiler, K.W., Rupen, M.P., Sramek, R.A., fc Panagia, N. (1993a). IAU Circular No. 5759. Van Dyk, S.D., Sramek, R.A., Weiler, K.W. & Panagia, N. (1993b). Astrophys. J., 409, 162. Van Dyk, S.D., Weiler, K.W., Sramek, R.A., fc Panagia, N. (1993c), Astrophys. J. Lett., 419, L69. Weiler, K.W., Sramek, R.A., Panagia, N., van der Hulst, J.M. & Salvati, M. (1986). Astrophys. J., 301, 790. Weiler, K.W., Panagia, N., Sramek, R.A., van der Hulst, J.M., Roberts, M.S. & Nguyen, L. (1989). Astrophys. J., 336, 421.

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Weiler, K.W., Panagia, N. & Sramek, R.A. (1990). Aslrophys. J., 364, 611. Weiler, K.W., Sramek, R.A., Van Dyk, S.D., k, Panagia, N. (1993). IA U Circular No. 5752. Wheeler, J.C. et al. (1993). Astrophys. J. Lett., 417, L71.

The circumstellar gas around SN 1987A and SN 1993J Peter Lundqvist Stockholm Observatory, S-133 36 Saltsjobaden, Sweden

Abstract The observational evidence for circumstellar gas around SN 1987A and SN 1993J is discussed along with interpretations of these observations. For SN 1987A we focus on its ring and for SN 1993.1 we mainly concentrate on its radio and N V A 1240 emission.

1. Introduction The circumstellar gas (CSG) around supernovae (SNe) provides information on the mass loss history of the dying star. When the SN explodes, the CSG is ionized by the radiation from both the SN and the gas shocked by the expanding ejecta. By looking at spectral signatures from the ionized CSG at increasingly large radii, we may peer deeper and deeper back through time into the mass loss history of the pre-SN. The recent bright and well-studied Type II SNe 1987A and 1993J have given us an unprecedented chance of doing so. Here we briefly discuss the CSG of these two SNe. 2. SN 1987A Light curves for the narrow UV emission lines from SN 1987A (Sonneborn et al. 1994) show that the emission starts ~ 70 days after the outburst with a roughly linear increase in strength until day ~ 400, followed by a gradual decline up to day ~ 1000 when the lines start to fall below detectability. In the optical there is very good information on the spatial flux distribution from observations with the NTT (Wampler et al. 1990; Wang & Wampler 1992) and the HST (Jakobsen et al. 1991; Plait et al 1994); the emission mainly comes from a patchy, elliptically shaped ring with semimajor and semiminor axes of ~ 0.830 (corresponding to ~ 6.2 X 1017 cm at 50 kpc) and ~ 0.605 arcsec, respectively. Even before the imaging observations, it was realized that the lines must be circumstellar since they were very narrow ( < 30 km s"1) (Fransson et al. 1989; Wampler & Richichi 1989) and indicated high N/O and N/C ratios (Fransson et. al. 1989). Models (e.g., Lundqvist & Fransson 1991) show that the source of excitation of the lines was the EUV/soft X-ray burst accompanying the shock 213

214

P. Lundqvist: Circumstellar gas around SN 1981A and SN 1993J 1

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outbreak from the SN, and that the evolution of the lines provide information on the characteristics of the burst. Lundqvist, Fransson k Blondin (1994) focus on fitting the light curves in greater detail, using more elaborate calculations for the burst (Ensman & Burrows 1992). Lundqvist et al. find that most of the ring has to be optically thick to the outburst in order for all light curves to start rising simultaneously around ~ 70 days. Using the 500M12 model by Ensman & Burrows (which has a peak effective temperature of ~ 7 x 105 K), the innermost part of the ring is heated to ~ 2 x 105 K and ionized up to C 5+ , N 5+ and O 6+ . There is a gradual decline in temperature and ionization away from the inner edge of the ring. The subsequent phase of cooling and recombination is governed by the density of the ring, and to model the bulk of the line emission the electron density, n e , has to be ~ 4 x 104 cm"3. To account for the emission after day 1000, a low-density component with ne ~ 1.5 x 104 cm"3 needs to be added. Preliminary results are shown in Figure 1, where the total ionized mass

P. Lundqvist: Circumstellar gas around SN J987A and SN 1993J

215

of the ring is 0.03 M 0 , 90 % of the mass being confined to the high-density component. Best fits are obtained if the low-density component is truncated just after the N4+-zone, in order not to overproduce NIV] A 1486. The He/H-ratio = 0.3, by number, and the overall metallicity is 0.3 times solar with C/N/O = 1/3.5/2.3. All lines were calculated with light travel time effects taken into account. For N V A 1240 (the only resonance line in Figure 1), it is also necessary to include scattering outside the ring to extend the modeled emission beyond ~ 900 days. In Figure 1, the scattering medium has been approximated by a cylinder with inner radius 8 X 1017 cm, with thickness 3x 1017 cm and with height 1 x 1018 cm. This roughly simulates the densest part of the wind exterior to the bipolar nebula in the hydrodynamic calculations by Blondin & Lundqvist (1993), and also gives good fits to the light curve. The line center optical depth through the cylinder in the equatorial plane is 2. The scattering gas is assumed to expand radially at 10 km s"1, and it is too dilute to contribute to any line emission. It is interesting that the imaging observations, the light curve model and the results by Blondin & Lundqvist all agree if the high-density component is confined to the equatorial plane, and is coated by a partially optically thick low-density component. The lobes extending further away from the equatorial plane are optically thin so that the wind exterior to the lobes is ionized at least up to N 4+ . Most light curve models use the assumption that the ring is smooth, thin and circular. If the modeled ring has either a finite width and/or height as in Lundqvist (1991), or if one allows for the patchy structure seen in the HST images (Plait et al. 1994) this only marginally changes the shape of the model light curves at times less than ~ 500 days, and even less at later times. The density estimates in these models are also rather insenstive to whether the ring is intrinsically elliptical or circular. The HST images are compatible with both these geometries. However, an intrinsically elliptical ring may introduce errors into a distance estimate to the LMC of the type carried out by Panagia et al. (1991). They assume a circle-like structure. In order to quantify this uncertainty we extend the analysis by Dwek & Felten (1992) to construct the sweeping of the ring by the light echo paraboloid, assuming a. tilted ellipse instead of a circle. We also assume that the turn on of the UV line emission occurs when the light echo paraboloid starts intersecting parts of the ring, and that the peak of the light curve occurs when the entire ring is enclosed by the paraboloid. One may then estimate the distance to the LMC, as long as one is also able to fit the shape and size of the observed ring. Allowing for the turn on occurring at epochs in the range 60 - 80 days and for the peak

216

P. Lundqvist: Circumsiellar gas around SN 1981'A and SN 1993J

between 400 - 450 days (cf. Fig. 1), we get a distance to the LMC in the range 43 - 55 kpc. The ratio of minor to major axis in these calculations is ^ 0.9, so even a rather mild intrinsic ellipticity causes almost a factor of ~ 2 larger uncertainty in the distance estimate than found by Panagia et al. . Only with high-resolution observations (both spatially and spectroscopically) when the ejecta shock reaches the ring may we somewhat reduce this uncertainty. It should be emphasized that the estimated ring velocity obtained from line profiles, ~ 10.3 km s"1 (Crotts & Heathcote 1991), would not be affected by such a small ellipticity as the one just discussed. 3. SN 1993J SN 1993J was first observed on 28 March, and quickly showed evidence for CSG both in optical (Cumming et al. 1994) and UV (Sonneborn 1994) spectra. On 5 April X-rays and, shortly thereafter, gamma rays were detected. This radiation is thought to be a result of the SN ejecta interacting with CSG (Fransson, Lundqvist & Chevalier 1994, and references therein). Another signature of such interaction was the detection of radio emission (Pooley & Green 1993; Weiler et al. 1994). The first detection occurred at 1.3 cm only five days after the optical outburst, and then later at longer wavelengths. It was quickly realized that the standard Chevalier (1982) model, in which the SN ejecta run into a stellar wind with an r~2 density dependence, and which has been very successful in accounting for radio emission from previous SNe (e.g., Chevalier 1984), could not explain the observed radio light curves for SN 1993J (Fransson et al. 1994). A slower density decrease with radius, more like an r~15 law, fits the data better. This is shown in Figure 2, where we have compared models with the observations by Pooley & Green at A = 2 cm. The model parameters and their values are discussed below. The importance of the unshocked CSG in Chevalier's model is that it provides free-free absorption producing an optical depth, rfr, that falls below unity shortly before the time of the observed maximum. Assuming a constant temperature and composition throughout an r~2 wind we get 8.2 x 10 A where M_5 is the mass loss rate in units of 10~5 MQ yr" 1 , v\ the wind velocity in units of 10 km s"1, T$ the wind temperature in 105 K, V4 the maximum ejecta velocity at time t, in 104 km s"1, and where we have evaluated the Gaunt factor for A = 2 cm and T5 = 1, as well as used a typical shock thickness of 30 % of the shock radius. The observations roughly give

P. Lundqvisi: Circumstellar gas around SN 1981A and SN 1993.1 1

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Tff — 0.5 at 2 cm on 22.5 April (Fransson et al. 1994), which translates into M-5/v\ w 4 Tg 4 , assuming ny\e/nu = 0.1, V4 (25.5 days) = 2, and taking limb darkening into account. For an r~15 wind, M/v would be lower than this estimate at radii smaller than ~ 4 X 1015 cm, but higher further out in the wind. The r ~ 1 5 wind appears to be present at least until day ~ 60. This suggests a variation of M jv on a time scale of a few hundred up to one thousand years, depending on the wind velocity. Inside ~ (1 — 2) x 1015 cm we cannot say anything about the mass distribution from the radio observations. The observed decline of the radio flux after maximum is roughly a /.~ 025

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P. Lundqvist: Circumsiellar gas around SN 1981'A and SN 1993J

and the spectral index a « -0.8 (Weiler et al. 1994). Chevalier's model would, for this spectral index, predict a t~08 for a freely expanding shock, and an even faster decline if the shock is retarding. Changing the radial dependence of M/v to M/v oc r 0 5 gives the right decline. However, a timevarying M/v is not the only way to produce a slow decline of the radio flux. This can also be made for constant mass loss at constant wind speed if the mass was lost from the progenitor preferentially in a plane to form a disk-like structure. A disk model is motivated by the fact that many groups (see Podsiadlowski 1994, and references therein) argue that the SN was a member of a binary system, and that it lost a substantial fraction of its envelope prior to the explosion. The mass not captured by the companion may have been focussed to the orbital plane of the system. Assuming a disk with a thickness that increases with radius as r^, though still retaining other assumptions inherent in Chevalier's model, we are able to model the observed flux decline for a freely exanding shock, if we choose = 0.4 and a = —0.8. Because the density in this model (assuming M/v to be constant with time) roughly falls off with radius as oc r~l~^, a value of close to 0.4 gives a radial density dependence similar to the r~15 model in Figure 2. If the disk is viewed nearly edge-on the <j> = 0.4 model would thus provide a good fit also to the early part of the radio light curve. If not, the radio maximum occurs too early (cf. Fransson et al. 1994). More detailed models have to consider tilt angle effects for the optically thick part of the light curves, and also take into account the fact that the circumstellar shock in a disk model is expanding faster away from the disk plane than within the disk. This increases the shock area, and also the amount of radiation produced. In order to test the disk geometry, one should also consider line profiles. Given the observed line widths (see below), we can state immediately that the disk is not being viewed face on, unless it is very thick. Our estimate for M/v is sensitive to the wind temperature close to the circumstellar shock. For the spherically symmetric case, Fransson et al. (1994) find that this is ~ (2-4) xlO 5 K, using the burst model by Shigeyama et al. (1994) together with hydrodynamic simulations for the ejecta/wind interaction. The main uncertainty in this calculation is the time of formation of the circumstellar shock. For a disk-like structure the detailed geometry is also important. In order to map the CSG around SN 1993J and to gain understanding about the time of shock formation, it is important to consider the narrow emission lines, in particular the strong NV A 1240 line. This was already present when IUE was pointed at the SN for the first time, on 30.2 March.

P. Lundqvist: Circumstellar gas around SN 1987A and SN 1993J

219

The line then declined by nearly two orders of magnitude in a week. During the same epoch, the line first indicated a velocity of ~ 103 km s" 1 , decreasing to ~ 100 km s"1 (see Sonneborn 1994, for details.) To explain both the width of the line and its high luminosity, Fransson et al. (1994) argued that the gas forming the line was preaccelerated by the radiation from the SN outburst; collisional excitation is not enough, photoexcitation boosted by a velocity gradient in the wind is also needed. It is interesting to note that N V A 1240 drops roughly at the same rate as the continuum near the line. One problem with N V A 1240 is that it shows that N 4+ has to be present at least until day 4, even when light echo effects are allowed for (Fransson et al. 1994). However, the calculations for spherical symmetry by Fransson et al. show that as soon as the circumstellar shock forms, i.e., no later than after ~ 0.5 days, the wind is completely ionized and the temperature is raised to ~ ( 2 - 10) x 105 K. The X-ray observations on day 7 confirm this; the circumstellar X-ray optical depth is negligible. Although hydrodynamic models invoking spherical symmetry show that the gas should be fully ionized, there are two explanations as to why N 4+ could be present. Either the shock does not form until after day 4, or the geometry is such that the densest part of the CSG subtends a rather small part of the volume swept up by the ejecta. In the first model, no ionizing flux except that from the photosphere is produced prior to day 4, and in the second the flux of ionizing radiation is lower than in the spherically symmetric case. We believe that the latter model is the most appropriate, because with only the photospheric radiation included, the gas recombines in less than ~ 2 days to ionization stages lower than N 4+ , resulting in no NV A 1240 emission. In particular, a model with a disk-like structure is attractive since it should be possible to tune this so that it takes ~ 4 days for the ionization front to move radially through the disk, hence producing a N4+-zone persisting this long. The shock in this model also produces the continuum radiation around 1240 A needed to photoexcite the line. In the delayed shock formation scenario, there is no such continuum emission. Once the circumstellar shock has formed, the ionizing radiation produced by the shocked gas is a mix of inverse Compton scattered and free-free emission. The calculations by Fransson et al. (1994) show that it is difficult to avoid producing so much inverse Compton scattered radiation that the models are incompatible with the observations. A disk structure will act to decrease the importance of inverse Compton scattered relative to free-free emission (Fransson et al. 1994), so this may also argue for a disk geometry. Finally, it should be noted that a model with blobs instead of a disk shares many of the virtues of a disk model. However, having the high density gas

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distributed radially to form a disk makes it easier to have a long period of high N 4+ abundance. Acknowledgements I am grateful to John Blondin, Roger Chevalier, Robert Cumming, Claes Fransson and Phil Plait for stimulating collaboration, and to David Green, Toshikazu Shigeyama and George Sonneborn for access to their results prior to publication. Support by the Swedish Natural Sciences Research Council is kindly acknowledged. References Blondin, J. M., k Lundqvist, P. (1993), ApJ, 405, 337 Chevalier, R. A. (1982), ApJ, 259, 302 Chevalier, R. A. (1984), Ann. N. Y. Acad. Sci., 422, 215 Crotts, A. P. S., k Heathcote, S. R. (1991), Nature, 350, 683 Cumming, R. J., Meikle, W. P. S., Walton, N. A., k Lundqvist, P. (1994), this volume Dwek, E., k Felten, J. (1992), ApJ, 387, 551 Ensman, L., k Burrows, A. (1992), ApJ, 393, 742 Fransson, C , Cassatella, A., Gilmozzi, R., Kirshner, R. P., Panagia, N., Sonneborn, G., k Wamsteker, W. (1989), ApJ, 336, 429 Fransson, C , Lundqvist, P., & Chevalier, R. A. (1994), in preparation Jakobsen, P., et al. (1991), ApJ, 369, L63 Lundqvist, P. (1991), in Proc. ESO/EIPC Workshop, SN 1981A and Other Supernovae, ed. I. J. Danziger, k K. Kjar (Garching: ESO), 607 Lundqvist, P., k Fransson, C. (1991), ApJ, 380, 575 Lundqvist, P., Fransson, C , k Blondin, J. M. (1994), preprint Menzies, J. W. (1991), in Proc. ESO/EIPC Workshop, SN 1981'A and Other Supernovae, ed. I. J. Danziger, k K. Kjar (Garching: ESO), 209 Panagia, N., Gilmozzi, R., Macchetto, F., Adorf, H.-M., k Kirshner, R. P. (1991), ApJ, 380, L23 Plait, P., Lundqvist, P., Chevalier, R. A., k Kirshner, R. P. (1994), ApJ, submitted Podsiadlowski, Ph. (1994), this volume Pooley, G. G., k Green, D. A. (1993), MN, 264, L17 Shigeyama, T., Suzuki, T., Kumagai.S., Nomoto, K., Saio, H., k Yamaoka, H. (1994), ApJ, 420, 341 Sonneborn, G. (1994), this volume Sonneborn, G., et al. (1994), in preparation Wampler, E. J., k Richichi, A. (1989), A&A, 217, 31 Wampler, E. J., Richichi, A.,k Baade, D. (1989), in IAU Coll. 120, Structure and Dynamics of the Interstellar Medium, ed. G. Tenorio-Tagle, M. Moles, k J. Melnick (Berlin: Springer), 180 Wampler, E. J., Wang, L, Baade, D., Banse, K., D'Odorico, S., k Gouiffes, C , (1990), ApJ, 326, L13 Wang, L. (1991), AkA, 246, L69 Wang, L., k Wampler, E. J. (1992), AkA, 262, L9 Weiler, K. W., Van Dyk, S. D., Sramek, R. A., Panagia, N., k Rupen, M. P. (1994), this volume

X-ray emission from the collision of supernova ejecta with circumstellar matter: SN 1987A k SN 1993J Tomoharu Suzuki, Toshikazu Shigeyama and Ken'ichi Nomoto Department of Astronomy, University of Tokyo, Bunkyo-ku, Tokyo 113, Japan

1 Introduction If a supernova progenitor has undergone significant mass-loss then the expanding supernova ejecta will eventually collide with this circumstellar material (CSM). Shock waves arising from the collision will compress and heat both the ejecta and the CSM. The emission from the shocked material depends strongly on the density distributions of the ejecta and the CSM, thereby providing important information about the nature of the CSM.

2 SN 1987A Images from the European Southern Observatory (ESO) (Wampler et al. 1990) and the Hubble Space Telescope (HST) (Jakobsen et al. 1991) revealed the presence of a ring-like structure at ~ 6 x 101' cm from SN 1987A. The outermost part of the supernova, ejecta is expanding at ~ 104 km s" 1 (Shigeyama & Nomoto 1990) and so is expected to collide with the ring ~ 10 years after the explosion.

2.1 Hydrodynamical model The progenitor of SN 1987A went through a red supergiant (R.SG) phase, and then contracted to a blue supergiant (BSG) before the explosion (for reviews, see Arnett et al. 1989, Hillebrandt & Hoflich 1989, Podsiadlowski 1992, and Nomoto et al. 1993a). This evolutionary scenario implies that the SN 1987A environment was formed as follows: the progenitor blew a stellar wind with a velocity ~ lOkms" 1 and a mass loss rate ~ lO~ 5 M0yr~ a during the RSG stage, and with corresponding values of ~ 550 km s" 1 and ~ lO" 6 M 0 yr- 1 during the BSG stage (Lundqvist & Fransson 1991). Consequently, the fast BSG wind struck the slow RSG wind, and a shock wave arising from this collision is propagating outward through the RSG wind. Dense regions formed behind the shock wave due to radiative cooling, resulting in the high-density nebula. Formation of the observed ring must be 221

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due to an effect with an aspherical geometry, for example, the rotation of the progenitor or the presence of a magnetic field (Chevalier & Luo 1994; Washimi et al. 1994). Our model consists of four components: the ejecta, the BSG wind, the RSG wind, and the ring (see Suzuki et al. 1993a for details). We adopt model 14E1 of Shigeyama & Nomoto (1990) for the supernova ejecta, in which the density distribution in the outer envelope is given by p oc r~ 86 . We assume the density distributions of the BSG and RSG winds to be 3 16 17 />BSG(0 = 9amucm" at 3 x 10 cm < r < 5.7 X 10 cm and pRSG(r) = 3 17 2 3 17 3 x 10 (r/10 cm)~ amucm~ at r > 5.7 x 10 cm, respectively. The ring lies in the RSG wind, touching the contact surface between the spherical RSG and BSG winds. Here we adopt 5.7 x 1017cm for the distance from the supernova to the ring (Jakobsen et al. 1991; Panagia et al. 1991). We choose O.O5M0 and 2.4 x 104amucm~3 respectively for the total mass and density of the ring (Lundqvist & Fransson 1991). The expanding ejecta collides first with the BSG wind (the first collision). A shock wave generated by this collision propagates outward through the BSG wind and collides with the ring and the RSG wind (the second collision). We therefore calculate the shock propagation from the first collision to the second collision using a one-dimensional spherical Lagrangian Piecewise Parabolic Method (PPM) (Colella & Woodward 1984), and from the second collision using two-dimensional cylindrical Smoothed Particle Hydrodynamics (SPH) (e.g., Benz 1990). After the second collision three shock waves compress the ejecta, the RSG wind, and the ring, respectively (Fig. 1). When the shock wave reaches the edge of the ring at / ~ 36 yr, the ring material is most strongly compressed and the X-ray emission reaches maximum, as will be described below.

2.2 X-ray emission We calculate the X-ray emission due to thermal bremsstralilung from the shocked material with the assumption that the shocked material is completely ionized. Figure 2 shows the X-ray light curve. X-rays from the ring are dominant because of its high density. The luminosity increases monotonically as the mass of the shocked region increases. At t ~ 36 yr, when the shock wave reaches the edge of the ring, the ring is compressed most strongly so that the luminosity attains its maximum of ~ 1037ergs~1. Afterwards the luminosity decreases because the ring expands and cools adiabatically. The luminosity from the ejecta and the BSG wind also increases monotoni-

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critical information for constraining ring formation models (see Luo & McCray 1991, Wang k Mazzali 1992, and Lundqvist 1992).

3 SN 1993J SN 1993J is characterized by relatively early observations of X-ray emission. We have performed detailed hydrodynamical modelling of the collision between the ejecta and the CSM and have calculated the X-ray emission to confirm the emission mechanism and to provide some constraints for model parameters which are still uncertain (Nomoto et al. 1993b).

3.1 Hydrodynamical model We assume that the CSM was formed by a steady wind from the progenitor and that its density distribution is p = /9o(?"/2xl014cm)~2 at r > 2xl0 14 cm. The outermost layer of ejecta expands homologously, so that its velocity distribution is v = ve^{r/2 x 1014 cm) and its density distribution is p = p e d ge (r/2 X 10 14 cm)~ n . In this study pedge = 3po is assumed (see Suzuki et al. 1993b). The collision forms two shock waves: a forward shock propagating into the CSM and a reverse shock moving back into the ejecta. Because of higher densities in the ejecta than in the CSM, the X-ray luminosity is dominated by the shocked ejecta (Fig. 3).

3.2 X-ray emission X-rays from SN 1993J have been observed with ROSAT at 0.1 - 2.4 keV (Zimmerman et al. 1993a,b,c) and with ASCA at 1 - lOkeV (Tanaka et al. 1993). The observed luminosities at 0.1 - 2.4keV and 1 - lOkeV are Xi_io ~ 5x lO^ergs" 1 (Tanaka et al. 1993) and £0.i-2.4 ~ 1.6x 10 39 ergs" 1 (Zimmerman et al. 1993b), respectively. The luminosity observed with ROSAT is gradually decreasing (Zimmerman et al. 1993c). The calculated X-ray luminosities £0.1-2.4 and £i_io and the ratio Xi_io/£o.i-2.4 are shown in Fig. 3. To summarize, for a larger wedge a "d smaller n, the electron temperature of the reverse shocked ejecta is higher and thus the emitted X-rays are harder. To be consistent with the relatively hard X-rays observed with ROSAT

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and ASCA, the ejecta are required to have a. high expansion velocity e.g. v edge ~ 5 X 104kms~1 and a shallow density gradient e.g. n ~ 8. It is important to construct a consistent hydrodynamical model starting from a more realistic configuration for the progenitor, i.e. the atmosphere with mass-loss should be smoothly connected to the CSM. The optical and X-ray light curves and spectra based on such a model would provide more accurate constraints on the still uncertain structures of the progenitor and the CSM. Continued observation of X-rays with ASCA and R.OSAT will be very valuable.

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Acknowledgements We would like to thank Kuniaki Masai and Hiroshi Itoh for discussions on the X-ray emission and Lifan Wang for discussions on the observations and models of circumstellar matter around SN 1987A. We also thank the ASCA team, Kuniaki Masai, and Claes Fransson for useful discussions on SN 1993J.

References Arnett, W. D., Bahcall, J. N., Kirshner, R. P., k Woosley, S. E. (1989). ARA&A, 27, 629. Benz, W. (1990). Numerical Modeling of Nonlinear Stellar Pulsation: Problems and Prospects, ed. J. R. Buchler (Kluwer Academic Publishers, Dordrecht), p.269. Chevalier, R. A. k Luo, D. (1994). Astrophys. J., 421, 225. Colella, P. k Woodward, P. R. (1984). J. Comput. Phys., 54, 174. Hillebrandt, W. k Hoflich, P. (1989). Rep. Prog. Phys., 52, 1421. Jakobsen, P., et al. (1991). Astrophys. J., 369, L63. Lundqvist, P. (1992). PASP, 104, 787. Lundqvist, P. k Fransson, C. (1991). Astrophys. J., 380, 575. Luo, D. k McCray, R. (1991). Astrophys. J., 379, 659. Luo, D., McCray, R. k Slavin, J. (1994). Astrophys. J., (in press). Masai, K. k Nomoto, K. (1994). Astrophys. J., in press. Nomoto, K., Shigeyama, T., Kumagai, S., Yamaoka, H. k Suzuki, T. (1993a). Supemovae, (Les Houches Summer School, COURSE X, Session LIV), ed. J. Audouze, et al. (Elsevier Science Publishers B.V., Amsterdam), in press. Nomoto, K., Suzuki, T., Shigeyama, T., Kumagai, S., Yamaoka, H. k Saio, H. (1993b). Nature, 364, 507. Panagia, N., GilmozziL R., Macchetto, F., Adorf H.-M. k Kirshner R. P. (1991). Astrophys. J., 380, L23. Podsiadlowski, Ph. (1992). PASP, 104, 1. Shigeyama, T. k Nomoto, K. (1990). Astrophys. J., 360, 242. Suzuki, T., Shigeyama, T. &; Nomoto, K. (1993a). Astron. Astrophys., 274, 883. Suzuki, T., Kumagai, S., Shigeyama, T., Nomoto, K., Yamaoka, H. k Saio, H. (1993b). Astrophys. J. Lett., 419, L73. Tanaka, Y. k the ASCA team (1993). IAU Circular No. 5753. Wampler, E. J., Wang, L., Baade, D., Banse, K., D'Odorico, S., Gouiffes, C. k Tarenghi, M. (1990). Astrophys. J., 362, L13. Wang, L. k Mazzali, P.A. (1992). Nature, 355, 58. Washimi, H., Mori, M. k Shibata, S. (1994). (preprint). Zimmerman, H. U. et al. (1993a). IAU Circular No. 5748. Zimmerman, H. U. et al. (1993b). IAU Circular No. 5750. Zimmerman, H. U. et al. (1993c). IAU Circular No. 5766.

Observations of Interstellar and Intergalactic gas towards SN 1993J in M81 D. L. King 1 , G. Vladilo2-3, M. Centurion 3 , K. Lipman 4 , S. W. Unger5 and N. A. Walton 5 1

Royal Greenwich Observatory, Madingley Road, Cambridge CBS OEZ, U.K. Observatorio Astronomico di Trieste, Via G.B. Tiepolo 11, 34131 Trieste, Italy 3 Instituto de Astrofisica de Canarias, 38200 La Laguna, Tenerife, Spain 4 Institute of Astronomy, Madingley Road, Cambridge, CBS OHA, U.K. 5 Royal Greenwich Observatory, Apartado 321, Santa Cruz de La Palma, 38780 Tenerife, Spain 2

1 Introduction Supernovae are important for the study of several astrophysical problems - nuclear processing in stellar interiors, distance scale determinations, and the chemical enrichment of the interstellar medium have all been explored. Additionally, they may be used as background probes of interstellar gas by studying, for example, Nal and Call lines, sampling the gas in the host galaxy, the Milky Way halo gas, and any intervening intergalactic gas. SN 1987A allowed the detailed study of gas towards and within the LMC to a distance D~0.05 Mpc (Vidal-Madjar et al. 1987, de Boer et al. 1987). With the unusually bright SN 1993J in M81 it is now possible to extend the search for interstellar/intergalactic absorptions beyond the local group of galaxies, out to the distance of the M81 group at a distance of D~3.25 Mpc. In this paper we present a preliminary study of high resolution optical interstellar spectra towards SN 1993J. The observations are described in section 2. The origin of the absorption lines, which fall into three distinct groups are discussed in section 3.

2 Observations The observations of SN 1993J were obtained during the nights 1993 April 48, using the Utrecht Echelle Spectrograph on the William Herschel Telescope at the Observatorio del Roque de los Muchachos, La Palma. The detector was a cooled Tektronix CCD of 1024 x 1024, 24 /.im pixels; the slit width was 1", resulting in a resolution of 5 and 6 km s"1 in the blue and red respectively. Preliminary results concerning the Call and Nal lines are discussed here a detailed analysis of these and other data will be presented subsequently 227

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(Vladilo et al. 1993, 1994). In Figure 1, the Call and Nal observations are shown. At least 11 components are identifiable in the Call data, fewer in the Nal data, covering the range -135 km s"1 < v/s7. < +165 km s"1. It can be seen that the absorptions fall into three distinct groups, labelled A-D, E-G and H-J in Figure 1.

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3 Identification of absorbing regions Routly and Spitzer (1952) demonstrated that the Nal/Call column density ratio in the interstellar medium can be used as a discriminant of the prevailing conditions. Specifically, Ca + is much more easily incorporated onto dust grains than Na° in the interstellar medium in Galactic disks, where the gas is at rest velocity. If the grains are disrupted however, the Ca + is released into the gas phase. Such conditions are found in high velocity gas - often distant from the disk. Thus we have Ar(NaI)/./V(CaII) > 1 in rest velocity gas, and JV(NaI)/W(CaII) < 1 in high velocity gas (Hobbs 1983). We use the ratio of the equivalent widths as an estimator of the Nal/Call ratio. This quantity has a value of ~ 2.2 in the disk of our Galaxy (Vidal-Madjar et al. 1987). 3.1 Interstellar gas in M81 Observations at 21 cm of M81 show that in the direction of the supernova HI is evident in the range —155 km s"1 < v\sr < —115 km s"1, peaking at visr = -135 km s"1 (Rots and Shane 1975), agreeing with components A and B in Figure 1. The ratios of the Nal/Call equivalent widths are 0.9 and 1.9 respectively, indicating an origin in rest velocity gas. It is possible that components A, B and C are contaminated by Galactic high-velocity clouds (Muller et al. 1963). An extended HI cloud - HVC C - covers much of the northern Galactic sky close to M81, where negative velocities v/sr < —100 km s"1 are expected. The most complete survey of HVCs (Hulsbosch & Wakker 1988) shows some with velocities close to components A and B, though not in the line of sight to SN 1993J. From this and the equivalent width ratios we conclude it is most likely that these components originate in interstellar gas at rest in the disk of M81. 3.2 Interstellar gas in the Galaxy The component G in Figure 1 at v/sr ~ 0 km s~l with an equivalent width ratio of 2.0 is most readily identified with local interstellar gas at rest velocity, although Genova et al (1990) do not find Mgll detections in this direction within 30 pc of the sun. Our data include Galactic stars in the general direction of the supernova; in more detailed work to follow we show that the local gas lies beyond ~ 100 pc, and out to ~ 500 pc, local gas is seen around rest velocity only. The equivalent width ratios of components E and F are significantly less than 1, indicating an origin in shocked gas. Rots and Shane (1975) showed

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there is no HI emission at these velocities at the supernova location. We identify these components with Galactic intermediate-velocity clouds similar to those studied by de Boer et al. (1990) 3.3 High Positive velocity gas The components HIJK are obviously not at rest with respect to Galactic interstellar gas; their velocity differs from the rotation curve of M81 at the position of the supernova by more than 200 km s"1, indicating that they are not formed in rest velocity gas in M81 either. Additionally, our observations of nearby foreground stars reinforce arguments against their formation in the disk of our Galaxy. Surveys of high velocity gas do not find HI at high positive velocities close to the supernova, which suggests their origin is not in Galactic high-velocity clouds. Components I and J have equivalent width ratios of 2.5 and 1.5 respectively, indicating they are formed in unshocked gas. Therefore, as they are unlikely to be formed in the disk of our Galaxy, or the disk of M81, they must be formed in a disk-like structure in the intergalactic medium between our Galaxy and M81. A normal galaxy existing between us and M81 would not escape detection; similarly a search for isolated intergalactic clouds found none in the M81 group (Lo and Sargent 1979). However, H I surveys have found large amounts of intergalactic gas linking M81 with NGC 3077 (Cottrel 1976), M82 (Cottrel 1977) and NGC 2976 (Appleton and van de Hulst 1988). These investigations show that this gas is likely to be the result of tidal interactions between M81 and its companion galaxies. We thus suggest that components HIJK originate in tidally stripped gas within the M81 group. Further evidence for locating these components within the M81 group of galaxies comes from similar velocities found in a filament linking M81 with NGC 2976, passing 25' west of SN 1993J (Appleton & van de Hulst 1988). References Appleton, P.N. &; van de Hulst, J.M. (1988). Mon. Not. R. astr. Soc, 234, 957. Cottrel, G.A. (1976). Mon. Not. R. astr. Soc, 174, 455. Cottrel, G.A. (1977). Mon. Not. R. astr. Soc, 178, 577. de Boer, K.S., Grewing, M., Richter, T., Wamsteker, W., Gry, C. & Panagia, N. (1987). Astron. & Astrophys., 177 , L37. de Boer, K.S., Morras, R. & Bajaja, E. (1990). Astron. & Astrophys., 233, 523. Genova, R., Molaro, P., Vladilo, G. & Beckman J.E. (1990). Astrophys. J., 355, 150. Hobbs, L.M. (1983). Astrophys. J., 265, 817. Hulsbosch, A.N.M. & Wakker, B.P. (1988). Astron. & Astrophys. Suppl., 75, 191. Lo, K.Y. & Sargent, W.L.W. (1979). Astrophys. J., 227, 756. Muller, C.A., Oort, J.H. & Raimond, E. (1963). C.R. Acad. Sci.. Paris, 257, 166.

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Rots, A.H. fc Shane, W.W. (1975). Astron. & Astrophys., 45, 25. Routly, P.M. & Spitzer L. (1952). Astrophys. J., 115, 227. Vidal-Madjar, A., Andreani, P., Cristiani, S., Ferlet, R., Lanz, T. & Vladilo, G. (1987). Astron. & Astrophys., 177, L17. Vladilo, G., Centurion, M., de Boer, K.S., King, D.L., Lipman, K., Stegert, J., Unger, S.W.U. & Walton, N.A. (1993). Astron. & Astrophys. Letters 280, Lll. Vladilo, G. et al. (1994). in preparation.

Mass Loss from Late Type Stars I. Cherchneff1 and A. G. G. M. Tielens2 1

Physics Department, New York University, 4 Washington Place, New York, NY 10003 - USA 2 Theoretical Studies Branch, MS 245-3, NASA Ames Research Center, Moffett Field, CA 94035-1000

Abstract Physical processes involved in mass loss from late type stars are reviewed, including the formation of an extended atmosphere, chemical nucleation and growth of dust grains, and radiation pressure driven winds. Extensive numerical and analytical studies show that shock waves are a viable mechanism to lift material above the photosphere of AGB stars where radiation pressure on newly formed dust can drive a cool wind. Atmospheres of RGB stars are permeated by limited strength acoustic shock waves and the force associated with them drives an outflow once the radiative cooling timescale becomes long compared to the dynamical timescale. This leads in a rather natural way to Reimer's law. Non-radial pulsations are likely important for protoplanetary nebula formation. Stardust formation is a chemical process regulated by thermodynamic as well as kinetic effects. Detailed models for C-stardust formation, based upon the extensive chemical literature on sooting flames, suggest that nucleation takes place close to the stellar photosphere, while the main chemical growth occurs at much larger distances (~ 2 — 3R.). Radiation pressure on dust coupled by friction to the gas determines the physical charateristics of AGB winds, but plays no role in RGB winds.

1 Introduction Mass loss from late type stars is a ubiquitous phenomenon which has important ramifications for the further evolution of the star. Typically, the nuclear burning timescale of a giant is 10~7 M 0 /yr. For comparison, the mass loss rate on the Asymptotic Giant Branch (AGB) varies from 10~6 to a few times 10~4 M©/yr. Hence, the envelope mass is most affected by the mass loss rate. Late type giants are also a main source of interstellar gas, injecting collectively about lM Q /yr in the interstellar medium. In contrast, massive stars inject only about 50% of that; about half on the main sequence and half as type II supernovae. Finally, late type giants are an important factory for Stardust, contributing 5 x 10~3 M©/yr to the galaxy. Despite its importance, the detailed mechanism of mass loss from late type stars is poorly understood, mainly due to the large number of physical processes involved and their inherent complexity. Nevertheless, driven by observations and simple theoretical arguments, a paradigm has evolved over the last decade. Mass loss is thought to be initiated by atmospheric pro232

/. Cherchneff & A. G. G. M. Tielens: Mass loss from late type stars

233

cesses that create an extended (~ 2 R») atmosphere in which dust nucleates and condenses. Radiation pressure on the dust couples to the gas through friction and drives a cool and gentle wind. Hence, in order to understand mass loss, we have to understand the processes that levitate the atmosphere above the photosphere, the nucleation and condensation of dust grains, and the radiation pressure forces on dust and its coupling to the gas. A comprehensive "ab-initio" theory combining all three processes is still a long way off, and generally, most studies attempt to isolate a specific process using a simplified prescription for the other processes involved. For example, dust nucleation has been studied in detail adopting a scale height approximation for the density distribution of the extended atmosphere and a uniform stationary stellar wind model. While such an approach is only approximate, much progress has been made in our understanding. This paper reviews the physical processes involved in mass loss from late type stars with an emphasis on AGB stars for which our understanding has most evolved. In § 2, we review the characteristics of mass loss by late type stars, including the origin of the extended atmosphere. Pulsational shock waves are thought to cause the levitation of the atmospheres of AGB stars, while acoustic waves are probably very important for red giants (§ 3). Dust nucleation and condensation is reviewed in § 4. Finally, the characteristics of radiation pressure driven winds are discussed in § 5. 2 Mass Loss from Late Type Stars Table 1 summarizes characteristic mass loss rates and outflow velocities associated with late stages of stellar evolution as well as (our opinion on) the origin of the extended atmosphere from which the wind originates. Except for massive supergiants, the various types of objects are different phases in low mass star evolution. These stars (earlier than M5III) evolve onto the red giant branch (RGB) once they have exhausted the H in their cores and start H-shell burning. Their mass loss rate is quite small compared to other late evolutionary phases, but the mass lost may amount in total to O.IM© (Rood 1973), probably during a brief Mira-like phase at the tip of the RGB (Willson 1989). Observed outflow velocities are low; 10% of the escape velocity. Likely, mass loss is initiated by acoustic heating of the outer atmosphere and the formation of a chromosphere (see § 3.2). During the second ascent onto the giant branch, the Asymptotic Giant Branch (AGB), the mass loss rate is considerably higher and mass loss dominates the further evolution of the star. This class of objects contains optically visible Miras and the more heavily obscured OII/IR stars. These

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Table 1. Mass Loss from Late Type Stars Object

Spectral Type

<M5 III M5-8 III M -B I KI RCr B MI Supergiants RGB AGB PPN

M

Vco

M 0 /yr

km/s

< 10- 7 1 0 - 6 _ 1 Q -4 1Q-3 a

10"6

io- 7 -io- 5

Extended atmosphere

acoustic 4-40 radial pulsations 4-40 30-300* non-radial pulsations non-radial pulsations 200c 10-40 non-radial pulsations ?

Notes: a) ~0.2M© is lost in 10-100 yr at the start of the PPN phase. Moderate mass loss may occur during this phase as well, b) Outflow velocity along the poles. Disk outflow velocity (~ lOkm/s). c) Outflow velocity of the "squirts". The shell coasts at ~20 km/s.

late type giants are pulsating variables surrounded by extensive dust and gas shells, coasting at moderate velocities (~10 km/s). The extended atmosphere is likely created by pulsational driven shock waves. Some evolution in the mass loss rate is expected during the ascent of the AGB since the pulsational period increases which sets the density scale height of the extended atmosphere (cf., § 3.1), but whether all stars eventually end the AGB phase as OH/IR stars is still an open question (cf., Habing 1990). Mass loss may also be punctuated by thermal pulses which temporarily interrupt the outflow. At the end of the AGB, the mass in the stellar envelope is reduced to ~ 10~2 M.0 and the star evolves at constant luminosity towards the planetary nebula phase. Objects in this transition phase are generally called protoplanetary nebulae (PPN). Present day mass loss rates are often difficult to determine. The wind velocities along the poles can reach fairly high values, ~200 km/s, due to the steep pressure gradient in that direction associated with the asymmetric density distribution. In general, while the AGB phase is characterized by spherical symmetric winds - as evidenced for example by OH maser studies - many PPNe are surrounded by massive (~ 0.2 M s ) circumstellar toroids. Similarly, many planetary nebulae show toroidal structures near the central star surrounded by a spherical halo (Balick et al. 1992). We opinionate that these toroids represent the superwind - the last stage of mass loss at the tip of the AGB. While the AGB wind is spherically symmetric and driven by stellar radial pulsations, when the envelope's mass becomes small enough (~O.2M0), the star becomes unstable against non-radial pulsations (Soker and Harpaz 1992). Perhaps, while the star readjusts itself on a Kelvin-Helmholz timescale (10-100yr; Willson 1989),

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almost all of the remaining envelope is ejected in a superwind. At the end of this stage, the left-over white dwarf core surrounded by a thin (~ 1O~2M0) envelope will appear as an F-G supergiant (Willson 1989); the beginning of the PPN stage. Other proposed models for the origin of these toroids include the transfer of orbital angular momentum from a close companion to the giant's atmosphere (cf., Livio in these proceedings) and gravitational interaction of the giant's outflow with a companion (Morris 1990). The surface composition (He- and C-rich, H-deficient), effective temperature and luminosity identifies R Cr B stars as PPNe. The presence of large circumstellar dust shells (Gillett et al. 1986) supports this connection. Models for R Cr B stars include resuscitation of a white dwarf to giant dimensions either through the merging of two white dwarfs or through a final helium flash after the H-shell has been extinguished (Renzini 1990). Alternatively, R Cr B stars are sometimes considered to be population II PPNe (ie., low mass ( < 1 M Q ) progenitor), similar to the RV Tauri variables with which they share some observational characteristics. The final helium flash model is perhaps the most likely and hence only about 10% of the AGB stars will go through this phase. R Cr B stars show pulsations with a period of ~40 days, modulated on a timescale of ~1000 days. On irregular intervals, their visual light curves go through pronounced minima, probably connected to extinction caused by the condensation of dust grains in small ejected clouds (Fayeef1987; Feast 1990). The localized nature of the ejection suggests that some form of non-radial pulsation is involved. Indeed, the very small (< 10~2 M0/yr) envelope mass would allow only highly non-radial pulsations. M-supergiants, descendants from massive stars, are know to lose mass at a high rate and with modest outflow velocities. Like AGB stars, they are also surrounded by extensive circumstellar dust and gas shells. Pulsational variability is present albeit not as pronounced as for the AGB stars. The "chromosphere" of supergiants may reflect shock propagation, as in AGB stars, or acoustic heating as in RGB stars (Dupree et al. 1990; Judge and Stencel 1991). Various pieces of evidence indicate that the mass loss has an episodic character (cf., Goldberg 1986), occuring on a ~ lOOyr timescale. In contrast to giants, supergiants often show distinct non-spherical mass loss. 3 Levitating the stellar atmosphere 3.1 Shock waves in AGB Stars The sound waves generated by the pulsations in the interior will steepen to shocks near the surface of the star due to the steeply decreasing den-

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/. Cherchneff & A.G.G.M. Tielens: Mass loss from late type stars

sity distribution. Since the same amount of energy has to be transferred (pAu3=constant), the wave amplitude, Av, has to increase. Once, these shock waves reach the photospheric surface, their energy, Eo, can be radiated away, Erad, and/or lead to a stellar wind, Ew\ i.e., Eo = Erad + Ew. If the density is very low, radiative cooling timescales are very long and the shock waves are adiabatic. The heated gas can then only lose its energy through expansion. Hence, adiabatic models lead to heavy mass loss. If radiative cooling timescales are very short, on the other hand, the shock will be isothermal and such models are characterized by very little mass loss (Bowen 1988). For a star in hydrostatic equilibrium, the density scale height is very small (a fraction of a stellar radius). Pulsations will thus lead to energy deposition in the outer atmosphere where the densities are too low to cool radiatively. As a result, the gas will expand and material is moved from the inner atmosphere to greater heights. Eventually, this will lead to a steady state in which the density above the photosphere is enhanced considerably; the extended atmosphere, which radiates away much of the pulsational energy. A parcel of gas accelerated upwards by a shock will gradually decelerate under influence of the stellar gravity and eventually fall back towards the star. This cycle will repeat itself with the next pulsations. The instantaneous velocity of the gas is generally very large, a few tenth of km/s. However, close to the photosphere, the net outward velocity averaged over a cycle is very small, a few m/s, and thousands of cycles are required to move the gas into the extended atmosphere. Hence, a given fluid element is shocked many times at almost the same height in the atmosphere. The shock waves in the atmosphere will approach but not exceed a limiting maximum amplitude which decreases with height in the atmosphere (Wilson and Bowen 1986). Essentially, this reflects momentum conservation. Each gas parcel accelerated by the shock and decelerated by gravity will have to return to exactly the same position. Time-averaged analysis can give some insight in the scaling of the density distribution with stellar and pulsational para.meters( Wilson and Bowen 1986). For isothermal shocks, the density scale height is given by

where H0(r) is the static scale height (kTr2//tGM) and j 2 is the average ratio of kinetic energy to gravitational energy of a parcel of gas. Figure la shows the shock amplitude in units of the escape velocity as a function of stellar pulsational properties (Willson and Bowen 1986). Comparison with models

/. Cherchneff & A.G. G.M. Tielens: Mass loss from late type stars i.

1

l

Av = gPl

' (a)

X 1 i

/ / /MIRA / .• MODELS .

_

/ .• Radial modes _ 0

"E

-1

1

*

-

o

—. ——-

_

To

1 10-1

5

10-1 7

StaticX . Photosphere \ ~ I I 2

•;

I (b) -

I

S 10-13 --

-3

log [Q(r/R.r 3/2 ]

1

10-11

Q 1

I

^10"9 -kl

«

l

/ /Ballistic -

-

if 1 <

i

237

3

Static

4

^

^

U_ 5

I ~ 6

R(10i3cm)

Fig. 1. (a) The shock amplitude in units of the escape velocity as a function of stellar pulsational parameters. The solid curves show the expected relationship using various simple analytical models. The dashed curve shows the results of detailed numerical calculations (Willson and Bowen 1986; Bovven 1988). (b) Calculated density distributions for various piston amplitudes at 2, 4, 6, 8, and 10 km/s (bottom to top; Bowen 1988). Note the convergence of r0 to the limiting photospheric radius, r,, with increasing amplitude.

show that eq. (1) with constant 7 provides a reasonable approximation. Typically, 7 2 ~ 0.4 and the dynamic scale height is a few times larger than the static one. Since the density depends exponentially on the scale height this leads to density enhancement of many orders of magnitude in the extended atmosphere. As the star ascents the AGB, the density in the extended atmosphere will increase because both Ho and 7 increase. For isothermal shocks, the limiting velocity amplitude as well as the density scale height do not depend much on the driving amplitude of the piston, but the actual density does. Both the velocity amplitude and the scale length are limited by momentum conservation. However, when the piston amplitude is increased, the position, r0, in the atmosphere where this limited velocity amplitude is reached shifts inwards and hence the zero point in the density distribution increases (Fig. lb). 3.2 Acoustic heating in red giants and supergiants

The turbulent convection zone in (super)giants produces pressure fluctuations which act as sources for the generation of acoustic waves. These waves grow to large amplitudes due to the steep density gradient and form hydrodynamic shocks which heat the gas. For a 3000K giant, the acoustic energy flux, Fm, is theoretically estimated to be ~ 107 erg cm"2 s"1 peaking at

238

/. Cherchneff & A. G. G. M. Tielens: Mass loss from late type stars

about jg/irCs ~ 5 x 10"4 Hz (A ~ 1O10 cm; Bohn 1984; Ulmschneider 1991). Like the thermal pressure gradient force, the mechanical energy flux exerts an outward force on the atmosphere which is given by em (l/Xm + 2/r) where em is a characteristic energy density (= Fm/vm with vm a characteristic speed) and Am the dissipation length. Hence, as for shocks driven by pulsations, these acoustic shocks will levitate the atmosphere. For an isothermal atmosphere in hydrostatic equilibrium, the density scale height is given by

where Ho is the static scale height, e* the thermal energy density, and 7^ is the ratio of the mechanical energy density to the gravitational energy density. Not surprisingly, the latter equation is very reminiscent of that derived for the pulsational case (c.f., Pijpers and Habing 1989). Again, the balance between the increase in the shock strength resulting from the decrease in density with the increase in the dissipation leads to a limiting shocks strength and in that limit the dissipation length becomes equal to Ho (Ulmschneider 1991). Hence, if em ~ et, the atmosphere becomes very extended. Using the values quoted above, we find that em/et ~ 3 x 10 13 /TI. Thus, ii\ the photosphere, this plays little role. But, assuming that the chromospheres of red giants and supergiants are due to acoustic heating, then at the base of the chromosphere em/et is large and the atmosphere is extended. Assuming that the mechanical flux in the chromosphere scales with that in the photosphere, we find gM ~ 4irrlFm, since cooling is dominated by expansion. Using Fm ~ T^jj (Bohn 1984), we recover Reimers law (M ~ L/g). Essentially, the acoustic flux taps directly the convective flux which carries the stellar luminosity. Extensive models for stellar winds driven by acoustic waves in RGB stars have been developed by Pijpers and Hearn (1989) and Pijpers and Habing (1989).

4 Stardust Formation Stardust formation is a ubiquitous and efficient process in late type red giants. We will concentrate here on the formation of carbon Stardust, which has been studied in greater detail than silicates. Early studies were based upon thermodynamics and classical nucleation theory, but such studies are fundamentally flawed. Classical nucleation theory was developed to describe the growth of droplets by condensation of water molecules in the atmosphere and, typically, it predicts dust condensation when the vapor is supersatu-

/. Cherchneff & A.G.G. M. Tielens: Mass loss from late type stars

239

rated. In contrast, the monomer for C-soot formation (ie., C-atoms) is not even present in the outflow from C-rich giants (Donn and Nuth 1985). Instead, C2H2 is the precursor molecule and chemical bonds have to be broken during the nucleation and condensation process and kinetic factors play an important role. Therefore, the chemical pathways which convert C2H2 into solid dust grains have to be identified and that is now generally recognized within the astrophysical literature (Frenklach and Feigelson 1989; Cherchneff et al. 1992). In C-rich giants, photospheric C is mainly in the form of CO, C2H2 and HCN and the chemistry that leads to C-soot formation is very similar to that in fuel-rich hydrocarbon flames, except for an overabundance of H. It is then possible to tap the wealth of information on soot formation in the chemical literature and detailed chemical schemes have been developed (Frenklach and Feigelson 1989; Cherchneff et al. 1992). These results have recently been reviewed (Tielens 1990; 1993) and a short summary will suffice here. Essentially, the formation of the first aromatic ring from acetylene initiates C-soot formation. This is the chemical equivalent of nucleation in earlier studies of Stardust formation. Once the first ring is formed, subsequent chemical growth takes place through alternating steps of H-abstraction (ie., formation of a radical site) followed by acetylene addition. After two such steps, the next aromatic ring will "close" (ie., cyclization) and this process can repeat itself. Hence, chemical growth takes place through the formation of larger and larger fused PAHs (Stein 1978). At a later stage, coagulation also becomes important and clusters/platelets are formed consisting of randomly stacked, small PAHs. Simultaneously, further chemical growth will take place, in particular at aromatic plane peripheries, leading to cross-linking of aromatic planes by tetrahedrally bonded carbon (ie., aliphatic hydrocarbon chains). Some cross-linking may also occur randomly dispersed in the aromatic planes. Finally, these platelets can cluster to form a spherical soot particle. This chemical growth scheme gives rise to a narrow temperature (~ 200 K) window around 1000 K in which actual nucleation can take place in a C-rich outflow (Fig 2; Frenklach and Feigelson 1989; Cherchneff et al. 1992). Essentially, for decreasing temperature (ie., further out in the flow) the chemical equilibrium between H addition and IT abstraction from naphthalene will shift towards the latter. Likewise, acetylene addition to naphthalene is favored over acetylene loss at lower temperatures. Eventually, for T<800 K, the endothermicity of the H abstraction reaction shuts this reaction down and further growth stops. Thus, the combination of kinetic and thermodynamic factors give rise to this window of "opportunity" for

240

/. Cherclmeff & A.G.G. M. Ttelens: Mass loss from late type stars

-15 500

900

1300

1700

T(K)

Fig. 2. Key reaction rates for soot nucleation in a typical outflow model. Rather than distance, the ordinate is labeled by the temperature at that point, to illustrate the presence of the nucleation window around 1000K.

chemical nucleation. While similar chemical steps are involved, the thermodynamic factors are more favorable for large]- PAHs and their "window of opportunity" is much larger. Hence, if they could form, they would grow big very rapidly close to the photosphere. The calculated yields for circumstellar PAH formation and hence carbon Stardust formation are always very small for realistic circumstellar envelope models (< 10" 5 ; Cherchneff et al. 1992). Basically, the low T required for initial PAH nucleation locates the growth regime far from the photosphere where the densities are low and growth is slow compared to the expansion timescale. Moreover, the yield is sensitive to the acetylene density (Y~ n{Q,iYL-i)A) and the density this far from the photosphere is very sensitive to the details of the models (ie.. stellar pulsation, temperature, mass loss rate). In contrast, observations show efficient soot formation for a variety of conditions. Table 2 compiles typical conditions for two red giants spanning the plausible range of parameter space. Realizing that expansion occurs at near the sound speed in the major condensation region, the growth in size within an expansion timescale is Art ~ 2e

R n 10 1 0 cm- 3 10 1 4 cm

(3)

with e the reaction efficiency of C 2 H 2 colliding with a grain. Hence, substantial grain growth requires e near unity (cf., Draine 1986), particularly for the warmer C-giants (Table 2). As discussed above, e is expected to be quite small because only a. small fraction of the time a. grain will have a radical site available for reaction and these reactions proceed with low efficiency.

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Table 2. Physical Conditions in AGB Photospheres and Condensation Zones

Te// P 2 K dyne cm" 2000 3000

105 102

Ho

L. Lo

R. Ro

R©

6000 3000

650 200

18.4 2.6

M" M©/yr 10~4

io-

7

" ' a3 cm" 3 x 1010 3x 108

Notes: a) Typical mass loss rates assumed, b) Density at the condensation point (r=3R») calculated from mass loss rate and sonic velocity.

Much of this problem resides with the nucleation and growth of the first clusters and various ways have been suggested to circumvent it. Small PAHs will be cooler than dust grains (ie., the inverse greenhouse effect) and this will shift the nucleation zone closer to the denser photosphere (Cherchneff et al. 1991). Alternatively, as a result of the dynamical motions generated by the shocks, a parcel of gas spends a considerable portion of its time at larger (up to 50% ) distance from the star than its nominal distance. The cooler temperatures there may favor PAH growth and, if large enough, these newly formed molecules might survive the next shock. Close to the photosphere, a typical parcel makes ~ 1000 of these excursions and their cumulative effect could be important (Cherchneff et al. 1992). Finally, it has been suggested that carbon-stardust formation might occur by growth on large, preexisting nuclei, for example small SiC grains (Caldwell et al 1993). Thermodynamics will then rapidly drive the condensation process close to the photosphere. While this suggestion does not address the nucleation process itself, models show that carbon soot formation might be very efficient under a wide va'riety of conditions.

5 Dusty Winds Driven by Radiation Pressure The structure of dust driven outflows has been studied extensively (Kwok 1975; Tielens 1983; Netzer and Elitzur 1993, Habing et al. 1993). Dust and gas are momentum coupled in the outflow; ie., the momentum gained by the dust from the radiation field is transferred to the gas. Some insight in the structure of circumstellar shells can be gleaned from a simple analysis of the momentum equation.

£ =
r

2

(4)

242

/. Cherchneff & A. G. G. M. Tielens: Mass loss from late type stars

where F is the ratio of the radiation pressure force to the gravity. For a constant F, this integrates to v(r) — Co + Uoo(l - r o /r)°' 5 with UQO = veac(r0) (F — 1) ' , and Co and r0 the sound velocity and inner radius. Hence, the acceleration zone is very small; v(2ro) ~ 0.7uoo. For typical dust parameters and assuming complete condensation of all condensibles, F ~4 and hence v^ ~ 2ve3c(r0). From observations, r0 ~ 3iE* and hence v^ ~20 km/s. In order to drive an outflow by radiation pressure on dust, F has to exceed unity integrated over the flow. For a constant F, this implies a minimum momentum flux in the flow, Mv > 5 x 10~8 M©/yr km/s. The exact value for this limit is to some extent (factor 2) sensitive to other physical processes in the flow (ie., heating/cooling; growth/sputtering; Tielens 1983; Kwok 1975; Gail and Sedlmayr 1986). However, the principle remains: when the momentum flux is smaller than this value, gas-grain collisions do not transfer enough momentum to the gas to drive the flow. Comparing with observations, we conclude that dust has a. dominant influence on the structure of outflows from AGB giants as well as M supergiants. In contrast, RGB stars have momentum fluxes considerably less than this value and dust does not play a role in driving their outflows. Their outflow is likely driven by deposition of acoustic energy in the upper atmosphere (cf., Pijpers et al. 1989ab). When the radiation pressure force is much larger than gravity (F ^ 1), the momentum equation can be integrated to yield Mv^ = TL*/C. Hence, for high mass loss rates, M scales directly with the average radiation pressure optical depth of the dust shell. In using this well know equation, it is sometimes assumed that r is less than unity, based upon the mistaken notion that the momentum in the flow has to be less than that in the radiation field. Actually, the total momentum (a vector) in the flow as well as in the radiation field are identically zero and r can well exceed unity. Moreover, the "momentum" in the radiation field doesn't change in the shell since Z»/c is constant to a high degree of precision {TV^^C < 1). In this equation, r is the radiation pressure efficiency averaged over the spectrum and averaged over the flow. Hence, for C-dust which has a steep spectral dependence for Q, r can be much less than the NIR dust optical depth through the shell; ie., the photons escape through frequency space. The Planck averaged opacity of silicates is independent of the temperature of the radiation field for T > 200 K (ie., Q(2/zm)~ Q(20/tm)) and the NIR optical depth is a good indicator of T. Extreme OH/IR stars (eg., OH26.5) are characterized by values of r ~ 10. These simple minded concepts are well preserved in more complete models (Kwok 1976; Tielens 1983; McGregor and Stencel 1992; Netzer and

/. Cherchneff & A.G.G. M. Tielens: Mass loss from late type stars 243 20

i

i i mill

IRC 10011

1 i i mill

104

1 i ii

IRC 10011

Heating -

A

10 3

15

VISCOUS

O

-1

'. \> \ -2

3 S

.

f

adiabatic

if

-3 1014

\l "

Cooling '.

H2O

o 1016

r(cm)

1014

'io° 1 0 1 6 r (cm)

Fig. 3. The calculated structure of the circumstellar shell around IRC 10011 (M = 10~ 5 M o /yr; Justtanont et al. 1994). The right hand side shows the gas temperature, T, and velocity, v, and the dust drift velocity, VdHjt- The latter is calculated for a grain size of 2500 A and scales with the square root of the grain size. The left hand side shows the variation of the heating and cooling terms in the energy balance. The sum of the A,'s is the power in the dependence of T on r (i.e., T~r A ). Note that A; greater (smaller) than zero indicates net. heating (cooling) by a process. Near the photosphere, HoO, CO and H:» can tap directly into the stellar photon field and the gas is quickly heated to about 2000 K (see text for details). Over most of the envelope the energy balance is dominated by adiabatic cooling, HoO rotational cooling, and viscous heating. Elitzur 1993; Habing et al. 1994). Figure 4 shows the results of a calculation for the circumstellar envelope of IRC 10011 (Justtanont et al. 1994). Besides momentum, the dust-gas coupling also transfers energy and this viscous heating can be an important heating source for the gas. Cooling occurs through adiabatic expansion as well as through molecular line cooling. The energy equation of the gas can be written as 4 ^ . .. dlnT (5) dl nr 3 with A, = Td/Tc(i), the expansion timescale over the cooling/heating timescale of process i per unit mass (R.odgers and Glassgold 1991). Hence, heating

244

/. Cherchneff & A.G.G.M. Tiekns: Mass loss from late type stars

(cooling) processes will increase (decrease) the exponent in the temperature law from -4/3, characteristic of adiabatic cooling. Figure 4 shows the various heating and cooling processes in a model for IRC 10011 (Justtanont et al. 1994). In the outer parts of the flow, adiabatic cooling dominates the energy balance. For r< 1015 cm, in the acceleration zone, H2O rotational cooling is important, steepening the temperature law. Viscous heating is of importance as well through much of the envelope. Very close to the inner boundary, vibrational pumping by stellar photons of H2O, followed by radiative decay, results in a net rotational excitation of the molecule. Collisional deexcitation then heats the gas and this gives rise to the initial sharp rise in the temperature. Clearly, the exponent in the temperature law varies with distance from the star due to the interplay of these processes (Fig. 4). Similarly, this exponent will vary from source to source, since viscous heating as well as rotational cooling will increase in importance with increasing mass loss rate. The temperature gradient is particularly important for the interpretation of molecular line observations. While in the past such analysis was often based upon the temperature law derived for the carbon star IRC 10216, the importance of temperature law variations from source to source are now well recognized (Sahai 1990; Rodgers and Glassgold 1991; Kastner 1992; Justtanont et al. 1994; Groenewegen 1994). Realizing that populating higher CO levels requires warmer and denser gas, the temperature structure of an envelope can be well probed by measuring a large number of CO transitions (Justtanont et al. 1994). Since this implies gas closer to the central star, such a study simultaneously probes the mass loss history of the star. This may be of particular interest for the extreme OH/IR stars which started their superwind phase < 103 yr ago. Indeed, the CO 1-0 transition measures by necessity the preceeding AGB wind since the superwind will not expand to the region where the 1-0 line is formed until t>> 104 yr (Justtanont et al. 1994). This may well be the origin of the large discrepancy between IR- and CO-based mass loss rates for OH/IR stars (Heske et al. 1990).

Acknowledgements We gratefully acknowledge many stimulating discussions with Harm Habing, John Barker, Kay Justtanont, and Chris Skinner. Theoretical studies of interstellar dust at NASA Ames are supported through NASA grant 39920-01-30 from the Astrophysics Theory Program.

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References Balick et al. (1993), in Planetary Nebulae, eds. R. Weinberger and A. Acker, (Kluwer, Dordrecht), 131 Bohn, H.U., 1984, Astron Ap, 136, 338 Bowen, G.W., 1988, ApJ, 329, 299 Cadwell, B.J., Wang, H., Geigelson, E.D., and Frenklach, M., 1993, in prep Cherchneff, I., Barker, J.R., and Tielens, A.G.G.M., 1991, Ap J, 377, 541. Cherchneff, I., Barker, J.R., and Tielens, A.G.G.M., 1992, Ap J, 401, 269 Donn, B. and Nuth, J., 1985, Ap J, 288, 187 Draine, B. T., 1986, in Interrelationship between Circumstellar, Interstellar and Interplanetary Dust, ed. J. Nuth, NASA CP-2403, 19 Dupree, A.K., et al. 1990, in Confrontation between Stellar Pulsations and Evolution, eds. C. Cacciari and G. Clementini, (Kluwer, Dordrecht),468 Feast, M.W., 1990, ASP conference Series, 11 , p. 538. Frenklach, M., and Feigelson, E.D., 1989, Ap J, 341, 372 Gail, H.-P and Sedlmayr, E., 1986, Astron Ap, 161, 201 Gillett, F.C., Backman! D.E., Beichman, C , and Neugebauer, 1986, ApJ, 310, 842 Goldberg, L., 1986, in The M-Type Stars, eds. H.Johnson &; F.Qnerci, NASA SP-492, 245 Gordreich, P. and Scoville, N., 1976, Ap J, 205, 144 Groenewegen, M., 1994, A& A, submitted Habing, H.J., 1990, in From Miras to Planetary Nebulae, eds. M.O. Mennessier and A. Omont, (Editions Frontieres, Montpellier), p.16. Habing, H.J., Tignon, J., and Tielens, A.G.G.M., 1993, Astron Ap, submitted. Heske, A. Forveille, T., Omont, A., van der Veen, W.E.C.J., and Habing, H.J., 1990, A& A, 239, 173 Judge, P.G. and Stencel, R.E., 1991, ApJ, 371, 357 Justtanont, K., Skinner, C.J., and Tielens, A.G.G.M., 1994, ApJ, submitted Kastner, J.H., 1992, ApJ, 401, 337 Kwok, S., 1975, Ap J, 198, 583 McGregor, K.B. and Stencel, R.E., 1992, Ap J, 397, 644 Morris, M., 1990, in From Miras to Planetary Nebulae, eds. M.O. Mennessier and A. Omont, (Editions Frontieres, Montpellier), p.520. Netzer, N. and Elitzur, M., 1993, Ap J, 410, 701 Pijpers, F.P. and Habing, H.J., 1989, Astron Ap, 215, 334 Pijpers, F.P. and Hearn, A.G., 1989, Astron Ap, 209, 198 Renzini, I., 1990, in Confrontation betxueen Stellar Pulsations and Evolution, eds. C. Cacciari and G. Clementini, (Kluwer, Dordrecht), 549 Rodgers, B. and Glassgold, A.E., 1991, 382. 606 Rood.R.T. (1973), ApJ, 184, 815 Sahai, R., 1990, ApJ, 362, 652 Soker, N. and Harpaz, A., 1992, PASP, 104, 923 Stein, S.E., 1978, J Phys Chem, 82, 566 Tielens, A.G.G.M., 1983, Ap J, 382, 606 Tielens, A.G.G.M., 1990, in Carbon in the Galaxy, ed. J.Tarter et al. NASA CP-3061, 59 Tielens, A.G.G.M., 1993, in Dust and Chemistry in Astronomy, eds. T.J. Millar and D. A. Williams, (Univ Press, Bristol), 103 Ulmschneider, P., 1991, in Mechanisms of Chromospheric and Coronal Heating, eds. P.Ulmschneider, E.R. Priest, and R. Rosner, (Dordrecht, Kluwer), 328 Willson, L.A., and Bowen, G.W., 1986, in Cool Stars Stellar Systems and the Sun, ed. M. Zeilik and D.M. Gibson, (Berlin, Springer Verlag), 385 Willson, L.A., (1989), in The Use of Pulsating Stars in Fundamental Problems of Astronomy, ed. E.G. Schmidt, (Kluwer, Dordrecht), 63

Kinematics and structure of the molecular envelopes around stars on the AGB and beyond Hans Olofsson Stockholm Observatory, S-13336 Saltsjobaden, Sweden

Abstract This review discusses the kinematics, the overall spatial structure, and the more detailed structure of envelopes around AGB-stars, post-AGB objects, and PNe, as deduced from molecular radio line emission. A possible scenario for the evolution of a circumstellar envelope as the star evolves from an AGB-star to the white dwarf stage is presented.

1 Introduction Red giant stars on the asymptotic giant branch (AGB) lose considerable amounts of matter in a slow wind, and a circumstellar envelope (CSE) of gas and dust is formed. The CSE gradually becomes thicker as the star evolves towards the end of the AGB. The mass loss decreases substantially as the star leaves the AGB and the CSE detaches itself from the star. Eventually, the central post-AGB object becomes hot enough to ionize the inner regions of the remnant AGB-CSE and a planetary nebula. (PN) forms. Thus, the AGB-CSE provides a common link through this evolutionary sequence, and hopefully much can be learnt about the late stages of stellar evolution as well as the metal enrichment of the interstellar medium through the study of its properties. Unfortunately, space does not permit a discussion of CSEs around supergiants (see e.g. Knapp & Woodhams 1993). 2 The physical characteristics of a "standard" CSE The number density distribution of H2 in the case of a smooth CSE, expanding radially with a constant velocity v^, that is formed by a steady mass loss rate M, and that consists mainly of molecular hydrogen, is given by, 2

i.e., for M = 10~5 MQ yr" 1 we have nn2(10ucm) « 109 cm" 3 and ??.H2(1017cm) « 103 cm" 3 , hence there is a considerable range in the number density. It 246

H. Olofsson: Kinematics and structure of circumstellar envelopes

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appears that the mass loss rate decreases considerably in a relatively abrupt way as the star leaves the AGB. This can be described, to a first approximation, as a detached CSE expanding away from the star. The number density at the inner radius of the CSE will decrease rapidly, and so will the extinction through the CSE (see e.g., Olofsson 1993). The kinetic temperature of the gas is mainly determined by the cooling due to the adiabatic expansion and by the heating due to the streaming of dust particles through the gas. A number of more or less uncertain cooling and heating mechanisms can be added (Truong-Bach et al. 1990). Huggins et al. (1988) found from a relatively detailed treatment, including for instance molecular line cooling, that in a C-rich CSE, formed by a mass loss rate of 10~5 MQ yr~*, the following relation applies (2)

Jura et al. (1988) showed that at large radii T^r) oc A/~ 05 , i.e., we expect thick CSEs to be significantly cooler than thin CSEs. The radiation field emanates from three major sources. The central star will have characteristics, which, depending on its evolutionary state, falls in the range between a. luminous (sa 1O4L0) and cool (w 2500 K) red giant and a much less luminous (< 1O2L0) but hot (< 105K) (pre-)white dwarf. The dust in the CSE contributes scattered starlight as well as its own heat radiation. Finally, the interstellar radiation field is of particular importance during the AGB evolution since it determines, through photodissociation, the sizes of the molecular envelopes. The density distribution of a. molecule does not, in general, follow the gas density distribution. Some molecules are of photospheric origin, but most of them are produced in the envelope by chemical reactions (see e.g., Millar 1988). At sufficiently large radii, the molecules become photodissociated. As a result of this, most molecular brightness distributions tend to be shell-like in spherical CSEs, but the peak radius differs from molecule to molecule, and furthermore it scales roughly as M 0 5 . The dominating dissociating radiation field depends on the evolutionary status of the object. It is the interstellar UV field that causes the photodissociation on the AGB. Beyond the AGB the UV radiation from the central star becomes increasingly important, and eventually it completely dominates the dissociation. The strong radial dependences of the physical conditions in a CSE make emission from different molecules and transitions emanate from different regions. Once the radiative transfer and the chemistry are mastered, efficient

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^

r[cm]

SiO masers, HCN masers < > H2O masers < > OH main lines >• SiO thermal, SiS, NaCl,... *• HCN,... < > CN,HC3N,... < > OH 1612 MHz > CO

Fig. 1. Radial regimes probed by different molecular emissions. The sizes of the molecular envelopes scale with the mass loss rate. use can be made of this in the exploration of the CSE. In this review we are limited to a discussion of the kinematics and the density distribution within the photodissociation radius of CO. All other species with lines in the radio regime have significantly smaller photodissociation radii (see e.g., Olofsson 1993). Fig. 1 gives a rough estimate of the spatial regions probed by different molecular emissions. At present there are 47 molecular species detected at radio wavelengths in CSEs [see Olofsson (1993) and add C O + and NaCN].

3 Kinematics of the molecular gas 3.1 AGB-CSEs We will now proceed to discuss the kinematics of the CSEs as deduced from the molecular radio line emission. The CSEs around AGB-stars, post-AGB objects, and PNe are discussed separately. The terminal velocity. A number of quite extensive studies of the terminal velocities of CSEs around AGB-stars exists today. Lewis (1991) analyzed the OH and SiO maser data on O-rich stars. The result is a i^-distribution that peaks at «15 k m s " 1 , and stretches over about an order of magnitude, from « 3 to 30 k m s " 1 . Olofsson et al. (1993b) obtained a i^-distribution for CSEs around bright carbon stars. It peaks at «12.5 k m s " 1 and has, when compared to that of the O-rich stars, a lower number of l o w - ^ sources and a higher number of high-u^ sources. These general trends agree with the results of Nyman et al. (1992), Kastner et al. (1993), and Loup et al. (1993). Recently, Lindqvist et al. (1992) showed that for a sample of galactic centre OH/IR-stars the latitude distribution of the high-u^ sources is more concentrated to the galactic plane than that of the low-t^ sources.

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A similar result for C-rich CSEs was presented by Barnbaum et al. (1991). Thus, it seems that VQC increases with the main sequence mass. Acceleration. The region over which acceleration occurs is presumably quite small, and hence there are severe observational difficulties. Most of the detailed observational information comes from maser emission. Mclntosh et al. (1989) concluded that the SiO masers of R Cas are located within « 3 x 1014 cm of the star, and they show no discernable velocity pattern. Nyman & Olofsson (1986) found that the total velocity coverage of the SiO maser emission was essentially the same for all stars in their study, and it may be independent of the terminal velocity of the CSE. These results suggest that the SiO maser emission originates in a region where the gas has not reached the terminal velocity, and it is even possible that ordered outflow motion is not dominating. At a. sligthly larger scale, 1014 - 1015 cm, most information is obtained from the H2O 22 GHz maser emission. Bowers et al. (1993) and Yates (1992) have recently provided evidence for acceleration of the gas in this region in four AGB-CSEs. The terminal velocity appears to be attained at « 1015 cm. Bowers (1992) found little support for an increase in the velocity range between the H2O and OH masers, and Bowers et al. (1989) concluded from OH maser maps that the expansion proceeds at nearly constant velocity at « 1015 cm. On the other hand, Lucas et al. (1992) showed that the size of the "thermal" SiO emitting region as a function of the line of sight velocity did not follow the trend expected for a constant velocity outflow. They claimed that v^ is not yet attained at a radius of « 5 x 1015 cm. However, Sahai & Bieging (1993) managed to model similar data assuming fairly rapid acceleration. Multiple winds. There is some evidence that multiple winds exist in the case of AGB-objects. Bujarrabal & Alcolea (1991) presented CO(2-1) spectra of R Cas and \ Cyg, which appear to show low-level emission wings extending about 5 k m s " 1 on both sides of the "normal" AGB-CSE feature. Possibly, there is a second, spatially confined, somewhat faster wind. Olofsson et al. (1990, 1992) found three C-stars, each one having two winds produced at different epochs. The old winds have expansion velocities of «15 k m s " 1 , while the present winds expand at « 5 km s" 1 . In this context we also mention two cases where the line profiles differ markedly from the expected one. The CO spectra of V Hya (a C-star) and n 1 Gru (an S-star) have prominent wings that suggest wind velocities of at least 35 and 25 k m s " 1 , respectively (Kahane et al. 1993; Sahai 1992). There is evidence in both cases that the high-velocity wind is bipolar. An explanation in terms of binary systems seems possible, and for the moment we prefer to regard these objects as exceptions.

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envelopes

3.2 Post-AGB CSEs The remnant AGB-CSE. Most of the information on the remnant AGBCSEs of post-AGB objects stems from CO observations. This is mainly due to the fact that these lines are by far the strongest, and do not weaken significantly as the object evolves. No change is expected in the kinematics of the remnant AGB-CSE as the star leaves the AGB, and this is also borne out by the observations (Bujarrabal et al. 1992; Nyman et al. 1992; Volk et al. 1993; Omont et al. 1993). An interesting point is that the outflow velocities of the detected post-AGB sources appear to be higher than those of most AGB-sources [see e.g., Nyman et al. (1992) and Volk et al. (1993)], but this may be a selection effect. Post-AGB winds. The most notable observational characteristic of a post-AGB CSE is the appearance of broad, low-intensity line wings in the CO spectra in particular. This suggests the presence of a high-velocity wind. The best example is CRL618. A projected outflow velocity approaching 200 kms" 1 [UOO(AGB-CSE)R;18 kms" 1 ] has been inferred for this apparently bipolar wind (Cernicharo et al. 1989; Gammie et al. 1989; Neri et al. 1992). CRL2688 is a post-AGB object that appears to be in an earlier evolutionary stage than CRL618. The NH3 data of Nguyen-Q-Rieu et al. (1986) clearly delineate an expanding disk-like feature in this source, and HCN and other molecular line observations suggest that this whole structure is slowly rotating by «1.2 kms" 1 at a radius of ss 6 X 1016 cm (Bieging & NguyenQ-Rieu 1988a; Nguyen-Q-Rieu & Bieging 1990). CO and HCN observations of Young et al. (1992) show the existence of a high-velocity outflow that may even consist of two winds with projected velocities of w40 km s"1 (possibly bipolar) and «100 kins' 1 [^(AGB-CSEJ^IS kms" 1 ]. In a recent study Jaminet et al. (1992) question some of these results. Another example of a probable post-AGB object is OH231.8+4.2. It has a. clear bipolar high-velocity outflow (Morris et al. 1987; Lindqvist et al. 1993). The projected outflow velocity is salOO kms" 1 . Bowers (1991) analysed the OH 1667 MHz emission and concluded that the outflow velocity increases from «10 kms" 1 in the equatorial plane to ss200 kms" 1 along the polar axis. Other examples of probable post-AGB objects with high-velocity winds are IRAS19500-1709 (Bujarrabal et al. 1992), IRAS22223+4327 (Omont et al. 1993), and IRAS23304+6147 (Woodsworth et al. 1990). Peculiar sources. A number of peculiar sources, whose locations in the scenario of stellar evolution are uncertain, have recently been detected. A classification as post-AGB objects do not seem unreasonable, and their kinematics certainly makes them worthy of a discussion. Among these sources

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we find high-latitude stars with supergiant-like spectra of class B to G. Recently, van der Veen et al. (1993) increased the number of such objects detected in circumstellar CO emission to 11, and to this we can add six sources detected only in OH (te Lintel Hekkert et al. 1992). Most of the sources exhibit normal AGB-CSE CO line profiles. However, HD101584 has a remarkably broad CO line with, possibly, seven individual peaks that may imply the presence of four winds with projected outflow velocities of «20, «40, «90, and «130 kms" 1 (van der Veen et al. 1993). An OH 1667 MHz maser emission map of the source shows a bipolar outflow with a projected velocity of «40 kms" 1 (te Lintel Hekkert et al. 1992). High-velocity (>40 kms" 1 ) OH emission is present in an additional four of these sources (te Lintel Hekkert et al. 1992). In this category of peculiar objects we also place four sources with very broad OH maser features (Av «150 to 200 kms" 1 ) (te Lintel Hekkert et al. 1988). One of these and two other objects have remarkable H2O 22 GHz maser emission (Likkel et al. 1992). The spectra consist of two groups of very narrow features that are separated by «100 to 250 kms" 1 . 3.3 PN-CSEs The remnant AGB-CSE. The evolution of the object from the tip of the AGB to the PN-stage is relatively fast, and the CO molecules in the AGBCSE normally manage to survive (as opposed to most other molecular species). The result is that for the young PNe the CO emission is still completely dominated by the AGB-CSE emission. Typical examples are NGC7027 (Jaminet et al. 1991) and IR.AS21282+5050 (Likkel et al. 1988), but many other examples are found (see e.g., Huggins & Healy 1989). The widths of the line profiles are comparable to those found towards AGB-stars. PN-induced kinematics. It appears that the development of the HII-region has, in most cases of detectable CO emission, not substantially affected the kinematics of the CSE. There are two notable exceptions. In Ml-17 the CO(2-1) line profile resembles that of a post-AGB object with a second moderate-velocity wind of «40 kms" 1 , while in BD+30°3639 the CO(21) line profile appears normal, but it suggests an outflow velocity of «50 kms" 1 , i.e., it is much higher than a "normal" AGB-CSE terminal velocity (Bachiller et al. 1991). Other less clear-cut examples of PN-induced kinematics exist. There is evidence of a higher outflow velocity along the major axis in the case of four bipolar nebulae, NGC2346 (Bachiller et al. 1989b), NGC3132 (Sahai et al. 1990), NGC6781 (Bachiller et al. 1993), and IC4406 (Sahai et al. 1991). It seems reasonable that these relatively high velocities

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are an effect of the interaction with the evolving nebula. This process is likely to have been initiated already during the early evolution. The CO(10) map data of NGC7027 by Bieging et al. (1991) and Graham et al. (1993) show that the major part of the AGB-CSE is still unaffected and expands at a velocity of «15 kms" 1 , but the CO(3-2) data of Jaminet et al. (1991) suggest a higher-velocity outflow, «23 kms" 1 , along the polar axis of the nebulosity. There is no evidence in this source or in any other PNe of the high-velocity wind seen in post-AGB objects. 4 The density structure of the molecular gas 4.1 AGB-CSEs In this section we discuss the overall spatial structure of the CSEs around AGB-stars, post-AGB objects, and PNe. Short discussions of the mass loss history of individual AGB-stars, and the detailed density structure ("dumpiness") of AGB-CSEs are also given. The inner region. The density structure in the region where the mass loss is initiated, < 1014 cm, may be probed by the SiO maser emission, but it is difficult to infer a density distribution from a maser brightness distribution. The SiO map of R Cas obtained by Mclntosh et al. (1989) shows an elongated structure, and so does the earlier map of Lane (1984) but the position angles are almost orthogonal to each other. However, this is an area where we may expect, in the near future, new observational data (see e.g., Colomer et al., 1992). The innermost region of the CSE is probed by the H2O 22 GHz maser emission. A high-quality map towards W Hya shows emission coming from a ring- or shell-like structure of size 3 x 1014 cm (Reid k Menten 1990). Bowers et al. (1993) and Yates (1992) obtained data for four stars and concluded that there are suggestions of asymmetric structures. In particular, the maser distributions are asymmetric relative to the estimated stellar position. Also C-rich stars have been shown to exhibit maser emission, in this case due to HCN. The HCN(0,2°,0, J=l-0) emission of IRC+10216 (Lucas k Guilloteau 1992) and CIT6 (Carlstrom et al. 1990) originate from regions within « 1015 cm of the stars, but the emission is not resolved. In the region 1015 to 1016 cm lies the OH masers in CSEs formed by moderate mass loss rates. Bowers k Johnston (1988) and Bowers et al. (1989) studied five objects and concluded that axi-symmetric density distributions would fit the data. Chapman et al. (1991) made a detailed study of U Ori, and concluded that the best, but by no means the only, explanation would be a toroidal distribution of radius « 1015 cm of the OH masers. The same radial regime is also probed by

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thermal emission from species like SiO, SiS, and NaCl. In total nine stars have been mapped in the SiO(u = 0, J = 2 - 1) line, and the brightness distributions appear circularly symmetric (the maps are based on a limited UV-coverage), implying that the CSEs are largely spherical at this size scale (Lucas et al., 1992; Sahai & Bieging 1993). Only in the CSE of IRC+10216, do we have information from many molecules. The SiS(5-4) map of Lucas et al. (1993a) suggests an elongated structure with PAw 30°. A NaCl(7-6) map also suggests an elongated structure, but now at a PAss -20° (Cernicharo et al. 1993), and this is also the orientation of the CS(2-1) emission (Guelin, these proceedings). The position angle of the SiS emission is close to that of the 2.2 fim emission (PA« 20°; Dyck et al. 1987). Thus, the NaCl and CS data may outline the density structure, while the SiS emission follows the light distribution, much like a bipolar optical nebulosity. Apparently, there is an axi-symmetric structure within a radius of 1016 cm of this CSE that on the largest scales appears to have an overall spherical symmetry. The extended envelope. On the spatial scale > 1016 cm most of the information comes from CO data, and OH 1612 MHz maser emission in the case of OH/IR-stars. Bujarrabal & Alcolea (1991) mapped the CO(2-1) emission of 10 AGB-CSEs (their beam corresponds to 7 x 1016 cm at the average distance of the stars 380 pc). The brightness distributions appear circular in all cases except one (V Cyg) where an elongated structure may be present. The extended CSE (radius > 5 x 1017 cm) around IRC+10216 shows no deviations from circular symmetry (Huggins et al. 1988; Truong-Bach et al. 1991). Olofsson et al. (1990, 1992, 1993a) have detected four carbon stars with large CO-emitting shells, Fig. 2. These CSEs, although by no means "normal", appear to have an overall spherical symmetry. In these and in many other studies, the relatively simple kinematics of a CSE is used to infer the three-dimensional structure. Several dozens of OH/IR-stars have been mapped in OH 1612 MHz maser emission (Bowers et al. 1983; Herman et al. 1985). It seems that the CSEs of these stars have an overall spherical symmetry. Once again, the only CSE thoroughly investigated is that of IRC+10216. The region within (1 — 5) x 1016 cm have been mapped in lines of e.g. HCN, HNC, HC3N, HC5N, C2H, C3H, C4H, C5H, C6H, C3N, MgNC, and SiC2 (see e.g., Bieging et al. 1984; Bieging & Nguyen-Q-R.ieu 1988b, 1989; Gensheimer et al. 1992; Takano et al. 1992; Bieging & Tafalla 1993; Dayal & Bieging 1993; Guelin et al. 1993; Guelin, these proceedings). There is to a first approximation an overall spherical symmetry in these data, but the emissions are very patchy, and they may not be centred on the star, possibly indicating a binary system (Guelin et al. 1993). Fig. 3 illustrates that different molecular emissions emanate from different regions of a CSE.

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60 30

b)

0 -30-60

Fig. 2. Maps of CO(J = 1 — 0) intensity towards S Set integrated over nine velocity ranges from blue (a) to red (i) (Olofsson et al. 1992). The size of the envelope as a function of velocity follows that expected for a. spherical shell expanding at a constant velocity. The scale is in arc seconds.

The mass loss history. We will only discuss direct evidence of variations in the mass loss of individual stars, and we are limited to a time scale of the order 104 years or less since most molecular line emissions are confined to radii less than a few 10 1 ' cm. Full-proof evidence in the circumstellar molecular line emission of varying mass loss, appears to exist only for the four carbon stars, R Scl, U Ant, S Set, and TT Cyg, Fig. 2. CO maps and modelling provide good evidence for present mass loss rates of a few 10~ 8 M0yr~ x , and mass loss rates in the past in excess of 10~5 M^yr'1 (Olofsson et al. 1990, 1992, 1993a; Bergman et al. 1993; Eriksson & Stenholm 1993; Yamamura et al. 1993). The high mass loss rate epoch seems to have been short, and it may have been triggered by a He-shell flash. These stars appear to be normal AGB-stars (apart from being C-stars), and hence

H. Olofsson: Kinematics and structure of circumstellar envelopes

i

i

i

255

*i

Fig. 3. Examples of molecular emissions observed towards IRC+10216, a) SiS(J = 6-5) (Bieging k Tafalla 1993), b) HC3N( J = 11-10) (Bieging k Tafalla 1993), c) MgNC(iV = 8 - 7 ) (Guelin et al. 1993), and d) C4H(/V = 1 0 - 9 ) (Guelin et al. 1993).

we can expect that a fair fraction of all AGB-stars occasionally have this highly episodic mass loss behaviour. The dumpiness. Whether the CSEs are smooth or clumpy is important not only for the understanding of the mass loss mechanism, but also for the excitation of the molecules, for the radiative transfer of the molecular line emission, and for the ability of the molecules to survive. Observations of SiO masers suggest spot sizes of a few 1012 cm (Colomer et al. 1992), but the difference between a maser spot size and the true clump size may be considerable (Goldreich k Keeley 1972). Assuming a clump size of 10 13 cm and IIH2 ~ 1010 cm" 3 (the SiO masers are quenched at higher densities, Lockett & Elitzur 1992) leads to a clump mass of « 10~7 M 0 . Measured

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spot sizes for the H 2 0 masers are « 1013 cm at a typical distance of 5 x 1014 cm from the star (Spencer et al. 1979), while the characteristic OH maser spot size in U Ori is « 4 x 1014 cm at a radius of 1015 cm (Chapman et al. 1991). The extended, shell-like brightness distributions observed towards IRC+10216 are very patchy, and so are those of the detached CO-shells observed by Olofsson et al. (1992, 1993a), Figs 2 and 3. Bergman et al. (1993) have shown that it is possible to model the brightness distributions in the latter cases as due to emission from a shell consisting of «1000 randomly distributed, identical, spherical clumps. Similar clumpy structures are found in the OH 1612 MHz maser maps.

4.2 Post-AGB CSEs The inner region. The high-velocity wind in CRL618 is apparently a bipolar outflow, oriented along the axis of the optical nebula, that runs into the remnant AGB-CSE, at a radius of a few 1016 cm (Neri et al. 1992). Very high spatial resolution information on this object is provided by the NH3 data of Martin-Pintado et al. (1993). In a recent study of CR12688 Jaminet et al. (1992) proposed that the high-velocity flow is not bipolar in the usual sense. The complicated structure may instead be the result of multiple mass loss events. However, the IICN data, of Bieging & Nguyen-Q-Rieu (1988a), the HC3N data of Nguyen-Q-Rieu & Bieging (1990), the NH3 data of Nguyen-Q-Rieu et al. (1986), and the SiS data, of Lucas et al. (1993b) suggest that the 40 km s"1 wind is bipolar and oriented along the optical bipolar nebula. The concentration of molecular gas in the equatorial plane of the optical bipolar nebulosity, may be a toroidal structure of radius w 6x 1016 cm (Nguyen-Q-Rieu et al. 1986; Bieging & Nguyen-Q-Rieu 1988a; Lucas et al. 1993b). The high-velocity outflow of OH231.8+4.2 is distinctly bipolar out to a radius of « 1017 cm, and aligned with the optical bipolar nebulosity, as shown in the CO data of Lindqvist et al. (1993). The extended envelope. The remnant AGB envelope of CRL618 appears to be spherically symmetric in its inner regions. The HCN data of Neri et al. (1992) suggest a dense core (< 8x 1016 cm) which is largely spherical. On the other hand, the extended parts of the CSE, out to 5 x 1017 cm, is sligthly elongated (Phillips et al. 1992). Truong-Bach et al. (1990) and Jaminet et al. (1992) provided CO maps of the extended envelope of CRL2688 that suggest no major departures from spherical symmetry. The CO observations of Bujarrabal et al. (1992) of the extended CSEs of four post-AGB objects suggest envelopes that are largely spherically symmetric, although there is a tendency for elongated structures in two cases.

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Peculiar sources. Little information exists on the morphology of the CSEs that surround the high-latitude stars with supergiant-like spectra. The OH 1667 MHz maser map of HD101584 of te Lintel Hekkert et al. (1992) provides evidence of a distinct bipolar outflow. A velocity gradient along the bipolar axis is interpreted as due to an interaction between the fast wind and a remnant CSE. The kinematics of the sources with remarkable high-velocity outflows traced by OH and/or H 2 O emission strongly suggests that there is a bipolar structure. This is also borne out by the OH observations of IRAS16342-3814 and W43A (te Lintel Hekkert et al. 1992; Diamond et al. 1985), and the H 2 O observations of W43A (Diamond & Nyman 1988).

4.3 PN-CSEs Young PNe. By far, the most well studied young PN is NGC7027. At the large scale, > 10 1 ' cm, the CO envelope appears very patchy, but roughly spherically symmetric (Bieging et al. 1991). However, there is a well defined symmetry axis at PA« 150° at a sligthly smaller scale (Bieging et al. 1991; Graham et al. 1993), which is aligned with that of the ionized gas. The CO emission peaks on either side of the HII-region along its minor axis, and so does the HCO+ emission (Likkel 1992; Deguchi et al. 1992). The high-quality HCO + (1-0) data of Cox et al. (1993) show a prominent toroidal structure. Another young PN, IRAS21282+5050, shows an equatorial density enhancement much like that of NGC7027 (Shibata et al. 1989). Evolved PNe. It appears that the more evolved PNe have molecular gas density distributions that can be described as more and more extreme versions of that of NGC7027 as the objects evolve. In NGC2346 CO observations suggest that the equatorial density enhancement is more prominent than in NGC7027, and it appears that the molecular gas in the polar regions constitutes the wall of a cavity (Bachiller et al. 1989b). M4-9 could be a similar object (Forveille & Huggins 1991), as well as the bipolar PNe, NGC3132, NGC6072, NGC6563, and IC4406 studied in CO by Sahai et al. (1990, 1991) and Cox et al. (1991), and Ml-92 studied in OH by Seaqvist et al. (1991). Also NGC6720 (the Ring nebula) probably has this structure (Bachiller et al. 1989a). Bachiller et al. (1993) have studied three PNe, NGC6781, NGC6772, and VV47, that may outline a continuation of this evolutionary sequence. NGC6781 is probably somewhat more evolved than NGC2346 and most of the molecular gas in the polar regions is gone. There is considerable small scale structure in the envelope. M2-51 may also fall in this category (Forveille & Huggins 1991). In NGC6772 the equatorial emission is even more dominating, and the fragmentation of the gas has de-

258

H. Olofsson: Kinematics and structure of circumstellar envelopes

veloped further. VV47 is the extreme example where only fragments of the equatorial density enhancement remain. NGC7293 (the Helix nebula) is, due to its proximity and large spatial extent, a perfect candidate for the study of the small scale structure of the molecular gas in evolved PNe. Forveille & Huggins (1991) showed that the molecular gas is fragmented into knots and filaments. Huggins et al. (1992) found small globules of molecular gas (mass > 5 x 10~6 Af© and nn2 > 105 cm" 3 ) that have managed to survive inside the ionized cavity of the nebula. 5 A possible scenario for the evolution of a CSE In this section we present a somewhat speculative, but possible, scenario for the evolution of a CSE, as the central object evolves from an AGB-star towards the white dwarf stage. It seems that the mass loss occurs in the form of lumps rather than as a smooth wind, and it may be episodic on a time scale ranging from that of the optical period up to the time between thermal pulses. A crude estimate of the number of lumps emitted per unit time is given by

i.e., only a few lumps per cycle may be ejected. The number of lumps within the H 2 0 and OH maser regions (» 5 X 1014 and « 5 x 10 t5 cm, respectively) would then be of the order 10 and 100, respectively. It is clearly difficult to determine the density structure of the inner regions of the CSE from the H2O maser observations if these lumps are randomly ejected, and this may explain why the brightness distributions are often asymmetric with respect to the stellar position. The OH main line masers should give a more balanced picture of the structure, but it is very likely that maser emission gives an enhanced view of for instance density contrasts. The lumps quite rapidly reach their terminal velocity, although there may be a spread in the radii at which this occurs, while the spread in terminal velocities is small. The velocity field is probably essentially isotropic. The lumps expand, due to internal motion, as they recede from the star, an effect that may possibly be reflected in the increasing spot size as one goes from the SiO to the H2O and finally the OH masers. The number of lumps within a region sampled by a typical single dish CO observation (sa 5 x 1016 cm) is of the order 1000. With this resolution both the brightness distributions and the line profiles will appear smooth. It is only in the high-resolution interferometer data, that the dumpiness becomes discernible (although one should be cautious when interpreting these data, since the observations are less sensitive to extended

H. Olofsson: Kinematics and structure of circumstellar envelopes

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structures), and it is particularly apparent in the external parts of the CSE and if the emitting material is distributed in a thin shell (for chemistry reasons or due to an episodic mass loss). It appears from the observations of the extended CSEs that the AGB mass loss must be close to isotropic. Soon after the star leaves the AGB, a high-velocity wind starts to blow. It may be inherently bipolar or it may be channelled by a compact equatorial density enhancement. There is evidence for such a structure on larger scales around young PNe. This may indicate that for some reason the AGB mass loss becomes axi-symmetric during the very last but intense mass loss period. The high-velocity wind during the post-AGB evolution and the development of an HII-region during the PN formation will further enhance the density contrast. The chemistry and the increased penetration of the UV light along the polar axis may make the molecular line brightness distributions more concentrated to the equatorial plane than is the actual density distribution. Eventually, there is molecular gas left only in a. fragmented ring or in some cases in a shell. The final structure possibly depends on the initial density contrast in the AGB-CSE. The fragments consist of individual lumps (or groups of lumps) in the remnant AGB-CSE, that have now been compressed by the pressure of the developing nebulosity. Acknowledgements I am grateful to the Swedish Natural Science Research Council (NFR) for travel support. References Alcolea, J., Bujarrabal, V., 1992, A&A, 253, 475 Bachiller, R., Bujarrabal, V., Martin-Pintado, J., Gomez-Gonzalez, J., 1989a, A&A, 218, 252 Bachiller, R., Huggins, P.J., Cox, P., Forveille, T., 1991, A&A, 247, 525 Bachiller, R., Huggins, P.J., Cox, P., Forveille, T., 1993, A&A, 267, 177 Bachiller, R., Planesas, P., Martin-Pintado, J., Bujarrabal, V., Tafalla, M., 1989b, A&A, 210, 366 Barnbaum, C , Kastner, J.H., Zuckerman, B., 1991, AJ, 102, 289 Bergman, P., Carlstrom, U., Olofsson, H., 1993, A&A, 2G8, 685 Bieging, J.H., Chapman, B., Welch, W.J., 1984, ApJ, 285, 656 Bieging, J.H., Nguyen-Q-Rieu, 1988a, ApJ, 324, 516 Bieging, J.H., Nguyen-Q-Rieu, 1988b, ApJ, 329, L107 Bieging, J.H., Nguyen-Q-Rieu, 1989, ApJ, 343, L25 Bieging, J.H., Tafalla, M., 1993, AJ, 105, 576 Bieging, J.H., Wilner, D., Thronson, Jr., H.A., 1991. ApJ, 379, 271 Bowers, P.F., 1991, ApJS, 76, 1099 Bowers, P.F., 1992, ApJ, 390, L27 Bowers, P.F., Claussen, M.J., Johnston, K.J., 1993, AJ, 105, 284 Bowers, P.F., Johnston, K.J., 1988, ApJ, 330, 339

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Bowers, P.F., Johnston, K.J., de Vegt, C , 1989, ApJ, 340, 479 Bowers, P.F., Johnston, K.J., Spencer, J.H., 1983, ApJ, 274, 733 Bujarrabal, V., Alcolea, J., 1991, A&A, 251, 536 Bujarrabal, V., Alcolea, J., Planesas, P., 1992, A&A, 257, 701 Bujarrabal, V., Gomez-Gonzalez, J., Bachiller, R., Martin-Pintado, J., 1988, A&A, 204, 242 Carlstrom, J.E., Welch, W.J., Goldsmith, P.F, Lis, D.C., 1990, AJ, 100, 213 Cernicharo, J., Guelin, M., Lucas, R., 1993, in prep. Cernicharo, J., Guelin, M., Martin-Pintado, J., Penalver, J., Mauersberger, R., 1989, A&A, 222, LI Chapman, J.M., Cohen, R.J., Saikia, D.J., 1991, MNRAS, 249, 227 Colomer, F., et al., 1992, A&A, 254, L17 Cox, P., Guilloteau, S., Bachiller, R., Huggins, P.J., Omont, A., Forveille, T., 1993, in prep. Cox, P., Huggins, P.J., Bachiller, R., Forveille, T., 1991, A&A, 250, 533 Cox, P., Omont, A., Huggins, P.J., Bachiller, R., Forveille, T., 1992, A&A, 266, 420 Dayal, A., Bieging, J.H., 1993, ApJ, 407, L37 Deguchi, S., Izumiura, H., Nguyen-Q-Rieu, Shibata, K.M., Ukita, N., Yamamura, I., 1992, ApJ, 392, 597 Diamond, P.J., Norris, R., Rowland, P., Booth, R., Nyman, L.-A., 1985, MNRAS, 212, 1 Diamond, P.J., Nyman, L.-A., 1988, in The Impact of VLBI on Astrophysics and Geophysics, eds M.J. Reid and J.M. Moran, Dordrecht, Kluwer, p.253 Dyck, H.M., Zuckerman, B., Howell, R.R., Beckwith, S., 1987, PASP, 99, 99 Eriksson, K., Stenholm, L., 1993, A&A, 271, 508 Forveille, T., Huggins, P.J., 1991, A&A, 248, 599 Gammie, C.F., Knapp, G.R., Young, K., Phillips, T.G., Falgarone, A., 1989, ApJ, 345, L87 Gensheimer, P.D., Likkel, L., Snyder, L., 1992, ApJ, 388, L31 Goldreich, P., Keeley, D.A., 1972, ApJ, 174, 517 Graham, J.R., et al., 1993, AJ, 105, 250 Guelin, M., Lucas, R., Cernicharo, J., 1993, A&A, 280, L19 Herman, J., Baud, B., Habing, H.J., Winnberg, A., 1985, A&A, 143, 122 Huggins, P.J., Bachiller, R., Cox, P., Forveille, T., 1992, ApJ, 401, L43 Huggins, P.J., Healy, A.P., 1989, ApJ, 346, 201 Huggins, P.J., Olofsson, H., Johansson, L.E.B., 1988, ApJ, 332, 1009 Jaminet, P.A., Danchi, W.C., Sandell, G., Sutton, E.C., 1992, ApJ, 400, 535 Jaminet, P.A., Danchi, W.C., Sutton, E.C., Russell, A.P.G., Sandell, G., Bieging, J.H., Wilner, D., 1991, ApJ, 380, 461 Jura, M., Kahane, C , Omont, A., 1988, A&A, 201, 80 Kahane, C , Audinos, P., Barnbaum, C , Morris, M., 1993, in Mass Loss on the AGB and Beyond, ed. H. Schwarz, ESO Conference and Workshop Proceedings No. 46, p.437 Kastner, J.H., Forveille, T., Zuckerman, B., Omont, A., 1993, A&A, 275, 163 Knapp, G.R., Woodhams, M., 1993, in Massive stars: their lives in the interstellar medium, eds J.P. Cassinelli and B. Edward, ASP Conference Series, Vol. 35, p.199 Lane, A.P., 1984, in VLBI and Compact Radio Sources, ed. R. Fanti, Dordrecht, Reidel, p.329 Lewis, B.M., 1991, AJ, 101, 254 Likkel, L., 1992, ApJ, 397, L115 Likkel, L., Forveille, T., Omont, A., Morris, M., 1988, .4 0.4, 198, Ll Likkel, L., Morris, M., Maddalena, R.J, 1992, A&A, 256, 581 Lindqvist, M., Habing, H.J., Winnberg, A., 1992, A&A, 259, 118 Lindqvist, M., Nyman, L.-A., Olofsson, H., Winnberg, A., 1993, in prep. Lockett, P., Elitzur, M., 1992, ApJ, 399, 704 Loup, C , Forveille, T., Omont, A., Paul, J.F., 1993, A&AS, 99, 291 Lucas, R., Cernicharo, J., 1989, A&A, 218, L20

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Lucas, R., Cernicharo, J., Guelin, M., Kahane, C , 1993a, in prep. Lucas, R., Guilloteau, S., 1992, A&A, 259, L23 Lucas, R., Omont, A., Kahane, C , 1993b, in prep. Lucas, R., et al., 1992, A&A, 262, 491 Martin-Pintado, J., Gaume, R., Bachiller, R., Johnston, K. , 1993, ApJ, 419, 725 Mclntosh, G.C., et al., 1989, ApJ, 337, 934 Millar, T.J., 1988, in Rate Coefficients in Astrochemistry, eds T.J. Millar and D.A. Williams, Kluwer, Dordrecht, p. 287 Morris, M., Guilloteau, S., Lucas, R., Omont, A., 1987, ApJ, 321, 888 Neri, R., Garcia-Burillo, S., Guelin, M., Cernicharo, J., Guilloteau, S., Lucas, R., 1992, A&A, 262, 544 Nguyen-Q-Rieu, Bieging, J.H., 1990, ApJ, 359, 131 Nguyen-Q-Rieu, Winnberg, A., Bujarrabal, V., 1986, A&A, 165, 204 Nyman, L.-A., et al., 1992, A&AS, 93, 121 Nyman, L.-A., Olofsson, H., 1986, A&A, 158, 67 Olofsson, H., 1993, in Molecular Opacities in The Stellar Environment, IAU Coll. No. 146, ed. U.G. J0rgensen, in press Olofsson, H., Bergman, P., Eriksson, K., Gustafsson, B., 1993a, in prep. Olofsson, H., Carlstrom, U., Eriksson, K., Gustafsson, B., 1992, A&A, 253, L17 Olofsson, H., Carlstrom, U., Eriksson, K., Gustafsson, B., Willson, L.A., 1990, A&A, 230, L13 Olofsson, H., Eriksson, K., Gustafsson, B., Carlstrom, U., 1993b, ApJS, 87, 267 Omont, A., Loup, C , Forveille, T., te Lintel Hekkert, P., Habing, H., Sivagnanam, 1993, A&A, 267, 515 Phillips, J.P., Williams, P.G., Mampaso, A., Ukita, N., 1992, A&A, 260, 283 Reid, M.J., Menten, K.M., 1990, ApJ, 360, L51 Sahai, R., 1992, A&A, 253, L33 Sahai, R., Bieging, J.H., 1993, AJ, 105, 595 Sahai, R., Wootten, A., Clegg, R.E.S., 1990, A&A, 234, Ll Sahai, R., Wootten, A., Schwarz, H.E., Clegg, R.E.S., 1991, A&A, 251, 560 Seaquist, E.R., Plume, R., Davis, L.E., 1991, ApJ, 367, 200 Shibata, K.M., Tamura, S., Deguchi, S., Hirano, N., Kameya, O., Kasuga, T., 1989, ApJ, 345, L55 Spencer, J.H., et al., 1979, ApJ, 230, 449 Takano, S., Saito, S., Tsuji, T., 1992, PASJ, 44, 469 te Lintel Hekkert, P., Habing, H.J., Caswell, J.L., Norris, R.P., Haynes, R.F., 1988, A&A, 202, L19 te Lintel Hekkert, P., Chapman, J.M., 1993, in prep. te Lintel Hekkert, P., Chapman, J.M., Zijlstra, A.A., 1992, ApJ, 390, L23 Truong-Bach, Morris, D., Nguyen-Q-Rieu, 1991, A&A, 249, 435 Truong-Bach, Morris, D., Nguyen-Q-Rieu, Deguchi, S., 1990, A&A, 230, 431 van der Veen, W.E.C.J., Trams, N.R., Waters, L.B.F.M., 1993, A&A, 269, 231 Volk, K., Kwok, S., Woodsworth, A.W., 1993, ApJ, 402, 292 Woodsworth, A.W., Kwok, S., Chan, S.J., 1990, A&A, 228, 503 Yamamura, I., Onaka, T., Kamijo, F., Izumiura, H., Deguchi, S., 1993, PASJ, 45, 573 Yates, J.A., 1992, in Dusty Discs, ed. P.M. Gondhalekar, RAL-92-084, p.92 Young, K., Serabyn, G., Phillips, T.G., Knapp, G.R., Giisten, R., Schulz, A., 1992, ApJ, 385, 265

Circumstellar Shells of Long-Period Variables Dust Formation and Optical Appearance A. Gauger, A. J. Fleischer, J. M. Winters and E. Sedlmayr Institut fur Astronomie und Astrophysik, Technische Universitat Berlin

1 Introduction Most long-period variables (LPVs) are surrounded by extended dusty circumstellar shells accompanied by considerable mass loss, which is often large enough to influence their evolution and to provide substantial amounts of chemically processed material to the ISM (Lafon & Berruyer 1991). Due to the improved instrumental capabilities at infrared and longer wavelengths a large amount of observational data is available for these objects (for a review see Habing 1990). On the other hand based on the progress in theoretical dynamical modelling of LPV atmospheres achieved during the last years (e.g. Bowen 1988, Fleischer et al. 1992) a detailed comparison of theory and observations seems to become accessible now.

2 The modelling method 2.1 Dynamical model calculations Our approach for the dynamical modelling of circumstellar dust shells of LPVs, which is described in detail in Fleischer et al. 1992, comprises the explicit solution of the hydro-equations in spherical symmetry adopting Lagrangian coordinates, the treatment of radiative transfer by the Eddington approximation for a spherical grey atmosphere according to Lucy (1976), the description of postshock cooling of the gas either by limiting cases (e.g. isothermal shocks) or by cooling functions, and the consistent detailed treatment of the formation, growth and destruction of carbon grains by means of a moment method (cf. Gauger et al. 1990). Starting with an initially dust free hydrostatic atmosphere the interior pulsation is simulated by a sinusodial motion of the inner boundary (piston), and the time evolution of the model is followed until a cyclic pattern has developed. The models are determined by the four stellar parameters mass M*, luminosity i», effective temperature T», and abundances of elements £,-, especially the carbon abundance ec, and by the period P and velocity amplitude Aup of the piston motion. 262

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t I 0.0 P

'

>

263

- t = 0.6 P

no o

0

2

4

6

8

10 0

2

4

6

8

10

Fig. 1. Radial structure of a typical dust, shell model at. two phases. Upper panel : gas temperature Tgas (solid), radiation temperature Trad (dotted) and velocity u (dashed). Lower panel : density p (dashed) and degree of condensation fCOnd (dashed). RQ is the initial stellar radius.

2.2 Radiative transfer

calculations

To study the optical appearance of the dust shell model we solve the frequency-dependent stationary radiative transfer problem for a given radial structure, i.e. at fixed instants of time, since the approximate treatment of radiative transfer in the hydrodynamical calculations do not yield spectra or lightcurves. The radiation field is calculated by means of a Feautrier-like iteration between the solution for the monochromatic intensity in (p,z) coordina.tes and the solution of the equations for the moments of the intensity. Grey absorption and thermal emission of the gas and frequency-dependent absorption and thermal emission of the dust are considered, and radiative equilibrium is assumed for the gas component and the dust grains, respectively. The dust extinction is calculated from the small particle limit of Mie theory, assuming the grains to be spherical and to consist of pure amorphous carbon. More details are given in Winters et al. 1993a.

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3 Results and discussion Figure 1 shows the radial structure of a typical dust shell model with the parameter values M* = 1M 0 , L, = 1O4L0, T* = 2600 K, ec/so = 1.8, P = 650 d, and Aup = 2km/s. In contrast to our previous calculations time-dependent postshock relaxation is considered assuming the source function of the shocked gas to be given by the Stefan-Boltzmann law at the gas temperature Tgas. Then the cooling rate per unit mass is proportional to Kgas(Tgas — T*a(i), where Traci is the radiation temperature and ngas is the gas opacity, which is assumed to be constant for simplicity. This rate leads to a more efficient cooling of the gas than the cooling functions introduced by Bowen (1988). It can be seen from Fig. 1, that due to the effective cooling Tgas differs only slightly from Tra(i except for the spikes at the shocks. In fact, the dynamical shell structure (velocity, density, etc.) is rather identical to the structure of a model discussed in Fleischer et al. (1992), which has been calculated with equal parameters but assuming isothermal shocks. A general result of our models is, that dust formation is confined to a certain phase during the period which results in a shell-like distribution of the dust as indicated by the degree of condensation fcond- In this model grain formation starts at about phase 0.4. When the density is enhanced by the passage of a shock at phase 0.5 effective grain growth takes place increasing fCOnd and the radiative acceleration which eventually leads to the amplification of the shock wave. In addition the new dust layer blocks the outgoing radiation and therefore heats up the material behind, producing the corresponding step in the temperature structure. Figure 2 shows the synthetic infrared lightcurves of the dust shell model with isothermal shocks. Despite the overall variation due to the varying luminosity there occur intermediate maxima at shorter wavelengths. The increase of the flux starting at phase 0.4 is caused by the decreasing optical depth due to the expansion of the dust shell, whereas the steep decline after phase 0.6 results from the formation of the new dust layer in the inner parts and the corresponding increase of the optical depth. The strength of these maxima decreases to longer wavelengths and they disappear for A > 2^im. At longer wavelengths additional maxima arise at phase 0.7, which are caused by the thermal emission of the energy absorbed at shorter wavelengths by the new innermost dust layer. This effect produced by the periodic formation of discrete dust layers may provide an explanation for similar features present in several observed

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0 0.5 1 1.5 2 time [P]

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I I I I | I I I I I I I I I I I I I 1

0 0.5 1 1.5 2 time [P]

Fig. 2. Synthetic infrared lightcurves at different wavelengths. Fluxes are given in magnitudes relative to the mean magnitudes. Note the different scaling of the ordinates.

lightcurves (e.g. Le Bertre 1992). A more thorough discussion will be given in a forthcoming paper (Winters et al. 1993b). Acknowledgements This work has been partly supported by the Deutsche Forschungsgemeinschaft, Schwerpunkt 'Theorie kosmischer Plasmen' (grant Se 420/5-4). References Bowen G.H., 1988, ApJ, 329, 299 Fleischer A.J., Gauger A., Sedlmayr E., 1992, A&A, 242, 321. Gauger A., Gail H.-R, Sedlmayr E., 1990, AfcA, 235, 345. Habing H.J., 1990, The evolution of Red Giants to White Dwarfs. In: Mennessier M.O., Omont A. (eds.) From Miras to Planetary Nebulae: Which path for stellar evolution ? Editions Frontieres, Gif sur Yvette, p. 16. Lafon J.-P.J., Berruyer N , 1991, AfcAR, 2, 249. Le Bertre, T., 1992, AfcAS, 94, 377. Lucy L.B., 1976, ApJ, 205, 482. Winters J.M., Dominik C , Sedlmayr E., 1993a, A&A, submitted. Winters J.M., Fleischer A.J., Gauger A., Sedlmayr E., 1993b, AfcA, submitted.

Observation of circumstellar shells with the IRAM telescopes Michel Guelin, Robert Lucas and Roberto Neri IRAM, 300 Rue de la Piscine, FS84O6 St. Martin d'Heres, France

Abstract Recent observations of circumstellar shells at arc second resolutions (i.e., 100 R, in the case of IRC+10216) reveal clumpy structures, asymmetries and jets. The most recent maps of such objects, observed with the IRAM Plateau de Bure interferometer and Pico Veleta telescope, are presented.

1 Introduction The dusty envelopes of late type stars are fascinating objects on their own; they are also interesting for what they teach us about IS chemistry. From their velocity field and density profile, we can study the mass loss during a crucial phase of stellar evolution: since the gas expands at nearly constant velocity in all, but the innermost envelope, the velocity maps yield a 3-D view of the molecule spatial distribution; the distributions of the different molecular species show how photochemical, molecule-molecule and grainsurface reactions proceed with time in a well behaved environment. The closest massive envelopes lie a few hundred parsecs away and have small angular sizes. The construction of large millimeter-wave interferometers, in particular the IRAM Plateau de Bure interferometer (Guilloteau et al. 1992), has provided a. major breakthrough in their investigation. 2 Molecular emission in IRC+10216 The most remarkable and probably closest massive envelope surrounds the bright IR object IRC+10216 (CW Leo). Its outer radius, observed in the mm lines of 12CO, the most abundant molecule and the best shielded from photodissociation after H2, is R = 3' (Guelin & Cernicha.ro, in preparation). Its warm 'core', where about 50 different molecules are detected, has a radius of ~ 15". The IRAM Plateau de Bure interferometer is ideally suited to study the IRC+10216 core. Its primary field has a FWHP of 55" at 90 GHz and its synthesized beam can be as small as 2". The large effective area of the instrument (460 m2 since April 1993) makes it sensitive enough to map 266

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with a few arcsec resolution scores of molecular lines. The possibility to observe simultaneously two 500 MHz-wide IF bands with a good frequency resolution, allows one to study 5-10 lines in a single observing session. A programme of mapping the A 3mm lines of the chemically most significant molecules is under way at Bure (Lucas et al. 1992, Guelin et al. 1993, Lucas 1994, Lucas et al., in preparation). 14 molecules have been observed to date: CO, CS, C 3 S, SiO, SiS, SiC 2 , HC 5 N, the carbon-chain radicals C n H, n — 2-6, and the metal compounds NaCl and MgNC. The synthesized beam HPW range from 3" to 10" (for CCH and C 5 H, two of the very first maps observed with the Bure interferometer). In addition, vibrationally-excited lines of C4H (Lucas 1994) and HCN (Lucas & Guilloteau 1992) ha.ve been observed. In a complementary programme, higher energy 2 mm and 1.3 mm lines of HC 3 N, C 4 H, SiC 2 , SiS, etc have been mapped with the IRAM 30-m telescope. The HC3N 1.3 mm data illustrate the importance of the excitation conditions in the derivation of molecular column densities (Audinos et al. 1994). Figure 1 shows the brightness distributions observed for the N= 10-9 C4H line in adjacent 3 kms" 1 wide channels. A point-like continuum source, coincident in position with the compact IR source and the central star, was removed from the maps. The source position is denoted by a cross at the center of the field. The IRC+10216 envelope is known to expand with a fairly constant radial velocity (-14.5 kms" 1 ), so that each velocity-channel sees the emission from a conical sector axed on the line of sight. The lowest and highest velocities arise from two narrow cones, respectively on the near and rear sides of the envelope. The velocities close to vsys = —26.5 kms" 1 arise from the vicinity of the meridian plane. The maps of Figure 1 have the appearance expected for a hollow spherical shell: the extreme velocity-channel maps show unresolved disk-like sources, whereas the median velocity-channels, which represent the line brightness distribution across the meridian plane, show broad ring-shaped sources. It is worth noting that the CjH sources, although almost perfectly circular, are not centred on the star (cross), but 2-3" more to the east. The 3 mm lines of several other molecules and radicals (HC5N, C3H, MgNC,...) show very similar brightness distributions (see e.g. Fig. 2). They arise mostly from a few arcsec thick shell with a radius of ~ 15". The velocity-channel maps have all the same clumpy appearance (the bright clumps lying often at the same positions) and show all intensity minima to the N and to the SE.

268

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13 17 00

13 16 40 •

13 18 20 S 47 58

S 47 5«

ra, dec (J2000)

Fig. 1. Velocity-channel maps of the C4H N= 10-9 line emission in IRC+10216, observed with the Plateau de Bure interferometer. The synthesized beam has a FWHP of 5 x 5". A point-like continuum source of 66 mJy, coinciding with the central IR source (cross) has been removed from the maps (Guelin et al. 1993).

The ring, or hollow shell appearance of the brightness contours of Fig 1 is not common to all molecules. One third of the species mapped so far show centrally-peaked distributions. Figure 3 shows the brightness distribution of the J=5-4 line of SiS. The SiS source is much more compact than the C4H and C3H sources and fits inside the 'hole' in the center of these latter. It is also elongated in the NS or NE-SW direction, i.e. more or less in the direction joining the minima of the C4H 'ring'. The differences between Fig. 2 and Fig. 3 cannot be explained by the envelope structure and/or line excitation effects. Obviously, the 'central hole' in Fig. 3 is not due to lack of gas and the central peak of Fig. 4 is not just an excitation artifact: it is observed as well in the easy to excite CO molecule (dipole moment fi = 0.11 D), as in the hard to excite species SiS and NaCl (/* = 10 D); it is not observed in C 4 H (fi= 0.7 D), HC 5 N (/t = 4 D) and MgNC (n = 5D). These differences must result from chemistry.

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IS It 00

UU(0

Fig. 2. The intensity distributions of the C 3 H 2 n 1 / 2 , J = 9/2-7/2 line (solid contours) and the C4H, N= 10-9 line (dotted contours), observed in IRC+10216 with the Plateau de Bure interferometer (Guelin et al. 1993). The intensities are integrated over the velocity interval -34 kms" 1 < v < —21 kms" 1 and the peak intensity of each species is normalized to 1; the contours are equally spaced. The angular resolution is 5" x 5". The maps are not corrected for attenuation by the primary beam (HPBW 50").

SiS &; C4H t«U

• 47 SS » 17 67 ra. (tee (JMOO)

Fig. 3. The intensity distributions of the SiS J = 5-4 line (solid contours) and of the C 4 H N=10-9 line (dotted contours), observed in IRC+10216 with the Plateau de Bure interferometer (Lucas 1994). The intensities (normalized to 1) are integrated over a 3 kms -1 -wide band centred on v,y, - -26.5 kms" 1 . The angular resolution is 3" x 3" for the SiS line, 5" x 5" for the C 4 H line.

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3 Chemistry in IRC+10216 The gas phase chemistry in IRC+10216 has been modelled by several authors (McCabe et al. 1979, Lafont et al. 1982, Glassgold et al. 1986, Nejad & Millar 1987, Cherchneff et al. 1993). Interstellar UV penetrates into the outer envelope (R >> 20R*) and dissociates the 'parent' molecules (CO, HCN, H2C2,...) expelled by the central star. The radicals and ions produced by photodissociation fuel a rich gas-phase chemistry which can form fairly complex molecules, such as the long cyanopolyynes HCnCN, n= 2-6. In the model of Cherchneff et al. (1993a,b), SiS is a 'parent' molecule coming from the stellar atmosphere. Its distribution should peak on the star, as is observed in Fig. 3. C4H is formed from CCH by radical-radical reactions in the outer envelope. It should thus have a shell-like distribution, as is observed in Fig. 2. For a mass loss rate of 2 10~5 MQ yr" 1 and a 'normal' interstellar UV field, Cherchneff et al. predict a radius of the C4H shell of 4-5 1016 cm, in perfect agreement with what is observed. The observed maps are however much too detailed to be fully explained by simple chemical models. According to Cherchneff et al., C3H is formed by radiative association, a much slower process than the C2H+C2H reaction leading to C4H; its abundance should peak at a larger distance from the central star. The great similarity of the C3H and C4H maps, in particular as concerns the extent of the emission, implies either that the shell observed in Fig. 2 marks the edge of the dense envelope, or that the three molecules are formed and destroyed together at a much faster rate than calculated by Cherchneff et al. (1993). Desorption of molecules from grain mantles could offer an interesting alternative to the relatively slow gas phase reactions. Desorption could result from irradiation by interstellar UV photons or, if the molecules are weakly adsorbed, from shocks. In all cases, the bright spots on Fig. 2, which are observed for molecules with very different excitation conditions, must correspond to high density clumps. The bright rings on Fig. 2 appear dimmer to the N and the SW. A similar pattern is observed for about all the molecules showing a shell-like brightness distribution. Very probably, these dark areas correspond to holes in the envelope. The holes define a N-SW axis roughly aligned with the major axis of the SiS source (Fig. 3) and with that of the IR source (Dyck et al. 1987). The outflow from the central star may not be as symmetrical as suggested by the fairly spherical CO outer envelope and may occur preferentially along a N-S axis, roughly perpendicular to the line of sight.

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4 The very fast outflow in CRL 618 Whereas bipolar outflow and shock chemistry are only suggested in IRC +10216, they are obvious in the more evolved stellar envelope CRL 618. CRL 618, which is thought to be a 'proto planetary nebula' at the end of the high mass loss phase, has already developed a compact HII region at the center of its thick dusty envelope. The central star, hidden by dust, has reached the temperature of a BO star; it is surrounded by two optical nebulosities showing high expansion velocities in the Ha and [NII] lines (Carsenty & Solf 1982). The distance to CRL 618, although very uncertain, is estimated to be circa 1.7 kpc. Molecular line observations show, as in IRC+10216, an extended molecular envelope (R ~ 40") with a compact dense core (R - 3") expanding at fairly constant velocity (20 kms" 1 ). Inside the core, one observes a bipolar outflow with a very large velocity (150-200 kms" 1 ) in the lines of half a dozen of molecules (Cernicharo et al. 1989). The envelope core and the fast molecular outflow have been mapped in several molecular lines, including the 2-1 (Cernicharo et al. 1989) and 3-2 (Neri et al., in preparation) lines of CO, using the IRAM 30-m telescope (resolution 12" and 9"). The HCN 1-0 line has been mapped with a 3" resolution with the IRAM Plateau de Bure interferometer (Neri et al. 1992). The high velocity HCN emission arises from a barely resolved source (2"x < 0.5") elongated in the EW direction and aligned with the optical lobes. Like in Ha, the HCN blue lobe is to the E and the red lobe to the W, the outflow axis being probably inclined by 45° with respect to the line of sight (Figure 4). Observing molecules with such large velocities with respect to the surrounding gas (200 kms" 1 versus 20 kms" 1 ) is surprising. It is known that almost all molecules are dissociated in a shock with Au > 50 kms" 1 (Hollenbach & McKee 1980). The high-velocity gas is likely to have been shocked and its molecules dissociated; the observed molecules must have reformed after the shock. The outflow of CRL618, whose geometry and time scales are roughly known, offers thus a text-book case of post-shock chemistry. According to the interferometric data, of Neri et al. (1992), the highvelocity HCN emission arises from several unresolved sources located at 3 10 16 cm from the star. The sources probably result from the impact of a primary outflow on dense envelope clumps (see insert in Fig. 4). The HCN re-formation time is found to be < 50 yr. The relative HCN and CO abundances in the high-velocity gas are consistent with those calculated by Neufeld & Dalgarno (1989) for post-shock chemistry in carbon-rich gas.

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M. Guelin et al.: Observations of circumstellar shells

' , ' '; : . ' : H I ' I R E G I O N n r s V " - ,

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Fig. 4. Morphology of the envelope core and high velocity outflow in the proto-planetary nebula CRL 618, as derived from HCN J = l-0 interferometric data and from Ha measurements. The HII region lies at the center of a compact and rather spherical core, itself surrounded by a low density halo (extending beyond the edges of the figure). A highly energetic flow collimates stellar material along an axis inclined by about 45° on the line of sight. Impacts of the flow on dense envelope clumps induce strong shocks. The high velocity wings of the HCN line arise from turbulent flows behind these shocks (Neri et al. 1992).

It will be interesting to reobserve the HCN line in several years to check whether the outflow geometry and molecular abundances have changed.

5 Other envelopes Other high resolution maps of molecular emission have been made with the Bure interferometer on CRL 2688, another 'proto planetary' C-rich nebula (HC5N, SiS, see Lucas 1994) and on NGC 7027, a planetary nebula (HCO + , see Cox et al. 1994). These maps, like those of IRC+10216 and CRL 618, show clumpy structures and deviations from spherical symmetry. An ambitious programme aimed at measuring the sizes and asymmetries of a large number of thick circumstellar envelopes is being carried out at Bure and Pico Veleta. Fifty stars of various types, known to be surrounded

M. Guelin et al.: Observations of circumstellar shells

273

by conspicuous envelopes, have been observed in the mm lines of 1 2 C 0 (Neri et al., in preparation); a subset of this sample has also been measured in SiO (Lucas et al. 1992). The Bure observations have been made in the 'snapshot' mode. For most objects, the sampling of the UV plane is too coarse to yield images of the molecular emission as detailed as those of Figs. 2 & 3; they allow us however to derive for each envelope the line radial brightness distributions in the different velocity-channels, as well as to measure the degree of ellipticity and the position angles of the envelopes.

References Audinos, P., Kahane, C , Lucas, R. 1994, Astron. Astrophys. in press. Bloemhof, E.E., Danchi, W.C., Townes, C.H. 1985 Astrophys. J. 299, L37. Bieging, J.H., Tafalla, M. 1993 Astronomical J. 105, 576. Carsenty, U., Solf, J. 1982, Astron. Astrophys. 106, 307. Cernicharo, J., Guelin, M., Martin-Pintado, J., Penalver, J., Mauersberger, R. 1989,

Astron. Astrophys. 222, LI. Cherchneff, I., Glassgold, A.E., Mamon, G.A. 1993 Astrophys. J. 410, 188. Cherchneff, I., Glassgold, A.E. 1994 Astrophys. J. in press. Cox, P., Bachiller, R., Guilloteau, S., Forveille, T., Huggins, P. 1994 in 34th Herstmonceux Conference Proceedings, Cambridge, UK, July 1993. Dyck, H.M., Zuckerman, B., Howell, R.R., Beckwith, S. 1987 Publ. Astr. Soc. Pacific, 99, 99. Glassgold, A.E., Lucas, R., Omont, A. 1986 Astron. Astrophys. 157, 35. Hollenbach, D., McKee, C.F. 1980 Astrophys. J. 241, L47. Lafont, S., Lucas, R., Omont, A. 1982 Astron. Astrophys. 106, 201. Lucas, R. 1992 in Astrochemistry of Cosmic Phenomena, Ed. P.D. Singh, (publ. Kluvver: Dordrecht), p. 389. Lucas, R., Guilloteau, S. 1992 Astron. Astrophys. 259, L23. Lucas, R. et al. 1992 Astron. Astrophys. 262, 49. Lucas, R. 1994, in IAU Coll. 140, Astronomy with Millimeter & Submillimeter Wave Interferometry, Hakone, Japan Oct. 1992. McCabe, E.M., Smith, R.C., Clegg, R.E.S. 1979 Nature, 281, 263. Nejad, L.A.M, Millar, T.J. 1987 Astron. Astrophys. 183, 279. Neri, R., Garcia-Burillo, S., Guelin, M., Guilloteau, S., Lucas, R. 1992 Astron. Astrophys. 262, 544. Neufeld, D.A., Dalgarno, A. 1989 Astron. Astrophys. 340, 869.

Morphology and kinematics of Planetary Nebulae Hugo E. Schwarz European Southern Observatory, Casilla 19001, Santiago 19, Chile

Abstract The morphology and kinematics of Planetary Nebulae (PNe) are reviewed based on the available data from the work of Keeler (1908) to the present. The correlations between morphological class and other fundamental parameters are explored. An HR diagram of the different classes is presented. Asymmetrical and special nebulae are considered: bipolars, point symmetry, irregulars, multiple shell objects, ansae, jets, and haloes. The role of binarity and the link with symbiotic stars are shown to be important. Post-PN nebulae are discussed. 1R, CO and radio data on the morphology and kinematics of PNe are briefly reviewed.

1 Introduction The morphology of Planetary Nebulae (PNe) has been studied for many years. Early catalogues of images of PNe have been presented by Keeler (1908) and Curtis (1918). These were the first comprehensive photographic samples of PNe images, and a lot of work on the classification and morphology has been based on these papers. The importance of a nebula's morphology, of which the diameter of the object is the simplest form, lies in the physical parameters that can be derived from it. The diameter of a nebula is related to its evolutionary stage, to its distance and is needed for the modeling of its spectrum. One method of distance determination uses the diameter and the H/3 flux of a nebula, under the assumption of a constant given mass for all PNe (Shklovskii 1956). Also, the mass and density can be calculated from the flux and diameter. Note that the distance is a particularly badly know parameter for nearly all PNe and impacts on all other parameters (see Terzian 1993). The shape or morphology of a PN allows, in principle, the complete history of formation and subsequent mass loss to be determined, albeit that care has to be taken in interpreting this (Mellema, these proceedings). Also the evolutionary stage of the central star (CS) (or its locus in the HR. diagram) can be linked to the morphology (Stanghellini et al. 1993). Determining the shape of a PN is not easy. The size of a nebular image on a photographic plate depends non-linearly on the exposure, and plates 274

H. E. Schwarz: Morphology and kinematics of PNe have a much smaller dynamic range than nebulae. The advent of CCDs which are linear detectors, has partially overcome these problems. However, the size as well as the shape of a PN depend also on the pass band. An extreme example of this is the nebula Ml-16 (Schwarz 1992) which is 94" long and bipolar in the light of Ha and [Nil], 26" and potato shaped in [OIII]500.7nm, and is listed by Acker et al. (1992) as being 3" in diameter. Obviously, great care has to be taken when drawing statistical conclusions about a sample of nebular images. The best is to use a large sample of images, all taken through the same niters and using a linear detector. Until recently, such homogeneous data were not available. The importance of knowing the kinematics of nebulae hardly needs stating: the expansion velocity of PNe (optical or radio data) allows distances to be found (see e.g. Pottasch 1984) when the angular expansion is also measured; the 3-dimensional structure of a nebula may be determined using kinematical information combined with the morphology, and for simple geometries the dynamical age of a PN may be determined. The obtained line ratios and shapes can provide electron densities and temperatures, excitation parameters, and velocity dispersions. Also here the efficient new detectors have allowed higher resolution studies of several fainter nebulae, and especially the spectra of peculiar, high velocity nebulae have allowed constraints to be put on several models of nebulae (e.g. Solf 1992, Corradi and Schwarz 1993a,b). This is important because until very recently only the brightest and largest nebulae have been studied, giving strong observational selection effects. Both the morphology and kinematics of PNe have been studied in other wavelength ranges than the optical: near IR, mm and sub-mm, and classical radio observations have made important contributions.

2 PN morphology 2.1 Optical After having been found by Messier in 1764, Herschel in 1785 first coined the words "Planetary Nebula", simply because of their visual likeness to planetary disks in his telescopes. The first catalogues of photographic images of PNe were published by Keeler (1908) and by Curtis (1918). Other image catalogues were published by Evans & Thackeray (1950), Voroncov-Vel'jaminov (1961), Westerlund & Henize (1967). Later well-known catalogues which collected nearly all previous samples together are those of Perek & Kohoutek (1967), and Acker et al. (1982, 1992). The problem with these data is that

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they come from a variety of sources and lack homogeneity. Intercomparing images and doing statistical studies was therefore difficult or impossible. Some catalogues with homogeneous, narrowband images taken with a CCD are those of Jewitt et al. (1986) (47 images of 47 nebulae), Balick (1987) (223 images of 51 nebulae), and Schwarz et al. (1992, SCM) (507 images of 255 nebulae). Several morphological classifications have been made over the years, using two methods: purely based on the two-dimensional observed shape of the nebula, or attempting to take into account the physics and three dimensional structure of the objects. Stoy (1933) classified a sample of PNe into 6 groups according to the observed shape, the spectrum and the brightness of the central star, and made the first attempt to link the morphology to evolution. Voroncov-Vel'jaminov (1934) used only the observed shape to arrive at 6 classes of morphologically different nebulae.lt is interesting to note that Stoy's system is hardly ever referenced, while that of VoroncovVel'jaminov was extensively used. This is probably due to its simplicity; Stoy's system is complex by comparison. Several other systems have been devised: Gurzadian (1962) used 11 classes of objects based on shape, Hromov & Kohoutek (1968) classified the nebulae into three groups using the inner or main structure (round, elliptical, and effectively bipolar) and three groups using the outer shape (regular archlike closed filaments, torn diverging filaments, and faint regular rings). It is clear from the above that it is not easy to classify, in a simple manner, a sample of nebulae. Complex descriptions or a large number of classes were used. Hromov & Kohoutek also attempted to design one single underlying three dimensional nebula, which could, by various projections, explain all observed shapes. In a study of multiple shell nebulae, Chu et al. (1987) concluded that about 50 % of PNe should (at some time during their life) experience a multiple shell event. Jewitt et al. (1986) found about 70 % of their nebulae to have multiple shells, indicating that between 25 and 60 % or so of all PNe have multiple shells. The images of SCM yield only about 12 % multiple shell objects. The prediction by Tuchman (1983) of 15 % multiple shell nebulae seems to fit the result of SCMs data. Observational selection effects can be important here. The most recent morphological classification is that of Schwarz et al. (1993a). They classified 361 nebulae from Ha images. Stanghellini et al. (1993) and Corradi & Schwarz(these proceedings) discussed the link between morphological class and other fundamental parameters such as CS evolution in the HR diagram, galactic distribution & kinematics, expansion

H. E. Schwarz: Morphology and kinematics of PNe velocity and so on. This classification yielded 4 groups: ellipticals, e, (with inner structure, es, or multiple shells, em); bipolars b, (with multiple events bra); pointsymmetricals (p), (with multiple events, pm), and irregulars, i, the group into which an object is put when it does not fit into any of the other classes. The remaining objects from the catalogue of SCM were either stellar (unresolved) objects or empty fields. Figures la to f show some example images from these classes. Correlations were found between morphological class and various physical parameters. Stanghellini et al. (1993) demonstrated that the classes occupied different loci in the HR diagram, as shown in Figure 2. The bipolars are distributed more evenly over mass than the general class of PNe, and the irregulars are heavily concentrated in one place in the L-T plane. These results are not understood yet, although the latter implies that irregulars are not only of the same mass, but are at the same evolutionary stage. Bipolars are more concentrated towards the Galactic plane, and are predominantly type I PNe (i.e. they are He and N rich, see Peimbert 1978), have an average expansion velocity of 170 kms"1 (about 20 kms"1 for all PNe), show neutral material (data from Huggins & Healy 1989), have a smaller radial velocity dispersion (24 versus 80 kms"1 for ellipticals), and have higher central star temperatures (140.103 versus 75.103 K for the ellipticals). These results are from Schwarz et al. (1993b). Note that Greig (1971, 1972) found that the presence of strong [Nil] and [Oil] lines correlates with the bipolar (binebulous in his terminology) morphology, as do type I nebulae. He suggested that bipolars had GKM giants and symbiotics as progenitors, while "other" nebulae came from LPVs. It is clear that as a class these are very different objects from the general population of PNe. 2.2 PNe in the Magellanic Clouds and globular clusters Blades et al. (1992) have published the first images of four Magellanic Cloud PNe, taken with the HST. SMC N2, and N5 are ring shaped nebulae, while LMC N66 is irregular. LMC N201 is unresolved. Two of their images are shown in Figure 3. Note that features down to 0".06 are resolved. There are only two PNe known in globular clusters: PSl in M15 and IRAS18333-2357 in M22. Schwarz & Claeskens (1993) have imaged PSl, while Borkowski et al. (1993) ha.ve done the same for the M22 PN. PSl was listed by Acker et al. (1992) as having a 1" diameter, but the image in Figure 4a shows a structured nebulaof about 10" in size. Borkowski et al. interpret the shape of their PN (Figure 4b) as due to interaction with the ISM through which M22 is moving with high velocity; in effect a "stripping"

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Fig. 1. Example nebulae of each of the morphological classes of Schwarz et al. 1993a: a. class e; b. class es; c. class em; d. class i; e. class b; f. class p

H. E. Schwarz: Morphology and kinematics of PNe

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Fig. 3. Images of Magellanic Cloud PNe of Blades et al. (1992) taken with the HST. (a) N5 deconvolved, (b) N66 original image. Smallest details are about 0."06 across.

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Fig. 4. Images of (a) PS1 in M15 in the light of Ho from Schwarz & Claeskens (1993), and of (b) IRAS18333-2357 in M22 through an [OIII] filter from Borkowski et al. (1993). The images have the same scale and orientation. These are the only two PNe known in globular clusters process is taking place. Figure 68 (NGC6829) of SCM shows such an interaction with the ISM for a galactic PN. A nebula can be influenced by its environment, and the resulting morphology can be used to infer properties of the local ISM, as well as of the nebula itself. The relation between morphology, evolution and other physical quantities is only recently being explored, especially in Magellanic cloud PNe, but has already yielded interesting results.

2.3 IR, (sub)-mm, and radio The recent advent of IR array cameras has made high resolution morphological studies of PNe possible at wavelength bands between 1.2 and 10/im. Speckle cameras, usually working at 2.2/tm, have imaged the inner structure of nebulae. Since atmospheric seeing is less troublesome at IR. wavelengths and the coherence length necessary for speckle work is longer, higher resolution can be achieved than in the optical. Early IR imaging was done by scanning single pixel detectors across sources. Here I will use data on NGC7027, Ml-16, and the Red Rectangle

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to show the progress of IR, mm, and radio imaging. Becklin et al. (1973) scanned NGC7027 at 10/«n, and showed that the shape was essentially the same as in the radio at 6cm, and unlike its optical appearance. The first array image was published by Arens et al. (1984), and is roughly the same as the scan map but with more detail. A beautiful recent result is that of Graham et al. (1993). Their image shows B17, H2 2.12/jm, and CO(l-O) emission, tracing respectively the ionised (HII) gas, the photodissociation region, and the neutral material. This is the present state of the art in such imaging on PNe. Beginning with the early scans in CO of Knapp et al. (1982), and the synthesised radio map at 4cm of Balick et al. (1973), there are many authors now publishing data at high resolution. The 5 GHz VLA maps of a large sample of PNe by Zijlstra (1989) and Aaquist & Kwok (1990) allow studies of the radio morphology, while in the (sub)-mm range the papers by Olofsson (1993), and Huggins (1993) review the field, and Zuckerman (1993) and Rouan (1993) deal with respectively IR imaging and speckle imaging of PPNe and PNe. Figure 5 shows a representative selection of such data. The discovery of a dust disk in the bipolar nebula M2-9 by Aspin et al. (1988), and the probable CO tori around NGC7027 (Graham et al. 1993) and Ml-16 (Sahai et al. 1993) strengthens the case for the role played by binaries in the forming of bipolars. This type of fundamental question can be addressed by these long wavelength observations: in the IR. the obscuration by dust is no problem and the presence and morphology of neutral material also probes the formation mechanisms.

3 PN Kinematics 3.1 Optical The discussion of the kinematics will be short because a review has recently been given by Schwarz (1993), in which various examples of kinematic data are shown. The parameters that can be derived are: Vex-P, velocity structure in the nebulae, and the radial velocity. These give information on: expansion mechanisms, dynamical ages, distances, and energetics of the nebulae. Models of ionisation fronts and shocks can be strongly constrained by combining the kinematics with the densities, which can be obtained from the same spectra. As example I take He2-lll (images in Figure 6a,b), an extreme bipolar nebula whose nature and evolutionary status were unknown. Long slit spectra show a velocity range of about 700 kms" 1 (Meaburn & Walsh 1989), and a density in the brighter outer lobes of about 300 cm" 3 (Dyson & Schwarz

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H. E. Schwarz: Morphology and kinematics of PNe

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1993). Using the shock explosion and wind models of John Dyson and collaborators, the distance is constrained to be about 500pc, and the shock cooling time is then equal to the dynamical age of the nebula at about 2000 yrs. Note that the previously published distances were between 2.8kpc (Maciel & Dutra 1992) and 4.15kpc (Kingsburgh & English 1992), giving unacceptably high values for the size, ejected mass, the energy in the wind, and the cooling time of the post shock gas. At 4kpc He2-lll is 14pc across! Again here the differences between the images of Figure 6 show the importance of detailed morphological studies. The kinematics of elliptical nebulae is dramatically different from that of bipolars. Weinberger (1989) published a critical compilation of expansion velocities, from which the average Vexj) for PNe is 20 kms" 1 . This value is effectively determined by the class of elliptical nebulae because they form about 70 % of all nebulae. Note that for extended nebulae the measurement of Vexp in the central part, as done by Weinberger can give spurious results: in the case of He2-lll, for instance, V exp =12 kms" 1 from Weinberger, but the echelle spectra of Meaburn & Walsh (1989) and the long slit spectra of Dyson & Schwarz give V e r p =350 kms" 1 ; for Ml.-16 these values are 10 respectively 250 kms" 1 (Schwarz 1992). Clearly, great care has to be taken, especially with bipolar nebulae, over the way in which expansion velocities are measured in PNe.

3.2 IR and aub-mm The IR has, until recently not been very useful for kinematical studies due to the lack of resolving power of IR spectrographs. Modern echelle gratings such as now in use at UKIRT allow velocity studies with long slits to be made. Ml-16 has been shown to have a similar velocity distribution in the K band as in the optical by Aspin et al. (1993), whose spectrum and velocity plot are shown in Figure 7a. Compare this with the optical velocity data of Corradi & Schwarz (1993b) in Figure 7b. It is becoming clear from these observations that an increasing expansion velocity towards the outer parts of a nebula is a general feature of certain classes of PNe. This implies a density distribution which falls of faster than r~ 2 , indicating that the mass loss rate in PNe is not usually constant with time. Published images also show that there are many nebulae showing evidence for discrete multiple mass loss events, or at least for a variable mass loss rate. The morphological and kinematical data combine to demonstrate unambiguously that the classical picture of Red Giant mass loss and PN formation cannot be right. The (sub)-mm observations, mainly of the CO molecule as a tracer of

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hydrogen, does not yet give a clear picture of the internal kinematics of PNe. The ratio between typical object sizes and the spatial resolution of the observations is not large enough, except for interferometric measurements on nebulae in which the CO emission is extended over a large area. The presence of high velocity gas can be inferred from the width and wings of the line profiles. Examples of what can be done are shown in Olofsson (1993) and Huggins (1993), and in their articles in these proceedings. A recent result is the continuum morphology of NGC6302 at 1.3mm observed with the SEST at ESO, shown in Figure 8 (Nyman & Schwarz 1993).

H. E. Schivarz: Morphology and kinematics of PNe

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4 Symbiotics with nebulae It is increasingly realised that symbiotic stars are evolutionarily related to PNe. (e.g Greig, 1972; Schwarz, 1988), and I therefore summarise some results on these stars in this review. Kenyon (1986) defines and discusses symbiotics, while finding charts and spectra have been presented by Allen (1984). Several symbiotics have been discovered to have extensive nebulae

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17 H 1

17H10"10s

Fig. 8. The 1.3mm continuum and superimposed Ha image of NGC6302 from Nyman k Schwarz (1993). associated with them. From Tables 1 and 2 it can be seen that the majority of the nebulae are bipolar in shape both in the optical and in the radio, strengthening the link between bipolarity and binarity, which has been suggested by several authors. Several of the optical nebulae and in at least one case (CIT Cyg) the radio nebulae too, show unusually high expansion velocities. A few 100s of kms""1 in the cases of He2-104, and BI Cru. Several of these nebulae have a Mira as their central star (or part of their central binary systems). He2-104 has a 400d period and BI Cru 280rf. The recent discovery of V417 Cen as a symbiotic star with a surrounding nebula has shown it to have a period of 246d (Van Winckel et al. 1993). An interesting question (which has partially been addressed by Schwarz & Corradi (1992)) is: how can a Mira produce a PN-like nebula and continue to be a Mira? A fast wind from the accretion disk sweeping up material from the dispersed PN produced by the WD could form these symbiotic nebulae: they are strictly speaking not PNe but post-PN nebulae.

287

H. E. Schwarz: Morphology and kinematics of PNe Table 1. Nebulae around

symbiotics-optical.

Discovery

Name

Size

Shape

1922 1983 1984 1989 1991 1992 1992 1993 1993

R Aqr V1016 Cyg HM Sge He2-104 AS201 BlCru He2-442 He2-147 V417 Cen

100" 0.4" 1.5" 75" 13" 150" 6" 7" 30"

b b b b e b jet b e

Dist.(pc)

Method

plate 240 spectrum 3400 spectrum 400 800 spectrum/image image 1800 image 2000 •? image 3400 image 5000 image

Ref. 1 2 3 4 5 6 7 8 9

References: 1 Lampland (1922a,b); 2 Solf (1983); 3 Solf (1984); 4 Schwarz et al. (1989); 5 Schwarz (1991); 6 Schwarz & Corracli (1992); 7 Yudin (1992); 8 Munari & Patat (1993) 9 Van Winckel et al. (1993) Note that in column 4 e is elliptical, b is bipolar, and i is irregular.

Table 2. Nebulae around symbiotics-radio. Discovery 1981 1981 1982 1984 1986 1986 1987 1988 1988 1988

Name AG Peg V1016 Cyg R Aqr HM Sge RS Oph CH Cyg RX Pup Hl-36 SS96 Henl383

Size 1.5" 0.5" 9" 0.5" 0.2" 1.5" 0.25" 5" 0.17" 4"

Shape

Dist.(pc)

Ref.

?

500 3400 240 400-2100

1 2 3 4 5 6 7 8 8 8

b b i

•?

b 9

i C ?

400 1300 7600 ? •?

1 Ghigo k Cohen (1981); 2 Newell (1981; 3 Sopka et al. (1982); 4 Kwok et al. (1984); 5 Porcas et al. (1986); 6 Taylor et al. (1986); 7 Seaquist & Taylor (1987); 8 Taylor (1988)

5 Future In this paper I have reviewed some aspects of the study of the morphology and kinematics of PNe from the observational point of view. It is clear that the major issue at present is the correlation of the morphology with other parameters of the objects. Evolution of the star(s) before, during, and after PN formation is only just being linked to the shape of the nebulae,

288

H. E. Schwarz: Morphology and kinematics of PNe

and the correlations found so far are not understood in anything but the most qualitative manner. Both observations and theory are making major progress: large numbers of nebulae are being imaged in the light of different emission lines, yielding the much needed large and homogeneous samples. Some inescapable conclusions from this work have already challenged the classical picture of PN formation. Multiple mass ejection events, the presence of shocks, jets, ansae, FLIERS (Balick et al. 1993), pointsymmetry, and neutral material (sometimes inside outer ionised haloes), all point to a formation mechanism which is much more complex than previously thought. Theoretical work on shocks, ionisation front propagation, and especially on their interaction with a clumpy medium, and in the geometry of binary systems have made modeling of these asymmetrical morphologies possible. The theory of the effects of radiative cooling on the observed nebular forms is now beginning to be developed, and will have a major impact on the field in the near future. For the moment, however, the observations are still leading the theories, due to the amazing complexity of these nebulae.

6 Acknowledgements My thanks to Robin Clegg for inviting me to give this review, and for financial support. References Aaquist, O.B., Kwok, S., 1990, A& A Supp. 84, 229 Acker, A. et al., 1982, Catalogue of the Central Stars of Galactic PNe. Acker, A. et al., 1992, Strasbourg-ESO Catalogue of Galactic PNe. Allen, D.A., 1984, Proc.ASA 5, 369 Arens, J.F., Lamb, G.M., Peck, M.C., Moseley, H., Hoffman, W.F., Tresch-Fienberg, R., Fazio, G.G., 1984, ApJ 279, 685 Aspin, C , Maclean, I.S., Smith, M.G., 1988, A& A 196, 227 Aspin, C , Schwarz, H.E., Smith, M.G., Corradi, R.L.M., Mountain, CM., Wright, G.S., Ramsay, S.K., Robertson, D., Beard, S.M., Pickup, D.A., Geballe, T.R., Bridger, A., Laird, D., Montgomery, D., Glendinning, R., Pentland, G., Griffin, J.L., Aycock, J., 1993, Ak A 278, 255 Balick, B., 1987, AJ 94, 671 Balick, B., Bignell, C , Terzian, Y., 1973, ApJ 182, L117 Balick, B., Rugers, M., Terzian, Y., Chengalur, J.N., 1993, ApJ 411, 778 Becklin, E., Neugebauer, G., Wynn-Williams, R.,, 1973, ApLetts 15, 87 Blades, J.C. et al., 1992, ApJ 398, L41 Borkowski, K.J., Tsvetanov, Z., Harrington, J.P., 1993, ApJ 402, L57 Corradi, R.L.M., Schwarz, H.E., 1993a, A& A 269, 462 Corradi, R.L.M., Schwarz, H.E., 1993b, Afc A 278, 247 Chu, Y-H., Jacoby, G.H., Arendt, R., 1987, ApJ Supp. 64, 529 Curtis, H.B., 1918, Publ.Lick.Obs. 13, 57 Dyson, J.E., Schwarz, H.E., 1993, in prep.

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289

Evans, D, Thackeray, A., 1950, MN 110, 429 Ghigo, F.D., Cohen, N.L., 1981, ApJ 245, 988 Graham, J.R., Serabyn, E., Herbst, T.M., Matthews, K., Neugebauer, G., Soifer, B.T., Wilson, T.D., Beckwith, S., 1993, AJ 105, 250 Greig, W.E., 1971, A&: A 10, 161 Greig, W.E., 1972, A& A 18, 70 Gurzadian, G.A., 1962, Planetary Nebulae, Moscow. Hromov, G.S., Kohoutek, L., 1968, Bull.Astr.Czech. 19, 1 Huggins, P.J., 1993, in: Mass Loss on the AGB and Beyond, ESO Proc. 46, 365 Ed. H.E.Schwarz Huggins, P.J., Healy, A.P., 1989, ApJ 346, 201 Jewitt, D.C., Danielson, G.E., Kupferman, P.N., 1986, ApJ 302, 727 Keeler, J.E., 1908, Publ.Lick.Obs 8, 1 Kenyon, S.J., 1986, The Symbiotic Stars, CUP. Kingsburgh, R.L., English", J., 1992, MN 259, 635 Knapp, G.R., Phillips, T.G., Leighton, R.B., Lo, K.Y., Wannier, P.G., Wootten, H.A., Huggins, P.J., 1982, ApJ 252, 616 Kwok, S., 1985, AJ 90, 49 Kwok, S., Bignell, R.C., Purton, C.R., 1984, ApJ 279, 188 Lampland, CO., 1922a, Pop.Astr. 30, 162 Lampland, CO., 1922b, PASP 34, 218 Maciel, W.J., Dutra, CM., 1992, A& A 262, 271 Meaburn, J., Walsh, J.R., 1989, A& A 223, 277 Munari, U., Patat, F., 1993, Afc A 277, 195 Newell, R.T., 1981, Thesis, NM Inst. for Mining and Technology. Nyman, L.-A., Schwarz, H.E., 1993, A& A in prep. Olofsson, H., 1993, in: Mass Loss on the AGB and Beyond, ESO Proc. 46, 330 Ed. H.E.Schwarz Peimbert, M., 1978, IAU Symp. 76, 215 Ed. Y.Terzian Perek, L., Kohoutek, L., 1967, Catalogue of Galactic PNe, Academia. Porcas, R.W., Davis, R.J., Graham, D.A., 1986, RS Oph and the recurrent nova phenomenon, 203, Ed M.F.Bode. Pottasch, S.R., 1984, Planetary Nebulae, Reidel, plOO Rouan, D., 1993, in: Mass Loss on the AGB and Beyond, ESO Proc. 46, 155 Ed. H.E.Schwarz Sahai, R., Wootten, H.A., Schwarz, H.E., Wild, W., 1993, ApJ in the press. Schwarz, H.E., 1988, IAU Coll. 103, 123 Eds. J. Mikolajewska et al. Schwarz, H.E., 1992, A& A 264, LI Schwarz, H.E., 1991, Afc A 243, 469 Schwarz, H.E., 1993, in: Mass Loss on the AGB and Beyond, ESO Proc. 46, 223 Ed. H.E. Schwarz Schwarz, H.E., Aspin, C , Lutz, J.H., 1989, ApJ 344, L29 Schwarz, H.E., Claeskens, J.-F., 1993, A& A in prep. Schwarz, H.E., Corradi, R.L.M., 1992, A& A 265, L37 Schwarz, H.E., Corradi, R.L.M., Melnick, J., 1992, A& A Stipp. 96, 23 (SCM) Schwarz, H.E., Corradi, R.L.M., Stanghellini, L., 1993a, IAU Symp. 155, 214 Eds. R.Weinberger, A.Acker Schwarz, H.E., Maciel, W.J., Corradi, R.L.M., Stanghellini, L., 1993b, in prep. Seaquist, E.R., Taylor, A.R., 1987, ApJ 312, 813 Shklovskii, I., 1956, Astr.Zh. 33, 315 Solf, J., 1983, ApJ 266, L113 Solf, J., 1984, A& A 139, 296 Solf, J., 1992, A& A 257, 228 Sopka, R.J., Herbig, C , Kafatos, M., Michalitsianos, A.G., 1982, ApJ 258, L35 Stanghellini, L., Corradi, R.L.M., Schwarz, H.E., 1993, A& A 279, 521

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Stoy, R., 1933, Obs. 56, 269 Taylor, A.R., 1988, IAU Coll. 103, 77 Eds. J.Mikolajewska et al. Taylor, A.R., Seaquist, E.R., Matthei, J., 1986, Nature 319, 38 Terzian, Y., 1993, IAU Symp 155, 109 Eds. R.Weinberger, A.Acker Tuchman, Y., 1983, IAU Symp. 103, 281 Ed. D.R.Flower Van Winckel, H. Schwarz, H.E., Duerbeck, H.W., Fuhrmann, B., 1993, A& A in the press. Voroncov-Vel'jaminov, B.A., 1934, Astrofizika 11, 40 Voroncov-VeFjaminov, B.A., 1961, Astrofizika 38, 75 Weinberger, R., 1989, A& A Supp. 78, 301 Westerlund, B.E., Henize, K., 1967, ApJ Supp. 14, 154 Yudin, B., 1993, IAU Symp. 155, Abstract booklet. Zijlstra, A.A., 1989, PhD Thesis, University of Groningen. Zuckerman, B. 1993, in: Mass Loss on the AGB and Beyond, ESO Proc. 46, 143 Ed. H.E.Schwarz

FLIERs in elliptical Planetary Nebulae M. Perinotto1, B. Balick2'1, Y. Terzian3, A. Hajian3 and A. Maccioni1 1

Dipartimento di Astronomia e Scienza dello Spazio, University of Firenze, Largo E. Fermi, 50125 Firenze, Italy 2 Astronomy Department, University of Washington, FM-20, Seattle, WA 98195, USA 3 Astronomy Department and NAIC, Cornell University, Ithaca, NY 14853, USA

1 Introduction In several elliptical PNe, a number of structures have been observed which are called : rims, shells, caps, ansae, knots, etc. in addition to the haloes. Some are macrostructures (rims, shells, haloes) constituting in a sense the bulk of the nebula itself. They can either be homogeneous or contain themselves smaller structures. Others are microstructures, and some of them qualify as FLIERs (see below). Balick et al. (1993; hereafter paper I) explored the spectra of microstructures in three elliptical PNe : NGC 3242, 7662, and IC 2149. They found that NGC 3242 and 7662 contain pairs of low ionization knots moving at supersonic velocity relative to the ambient gas. They have been named FLIERs (fast low ionization emitting regions). Various interpretations were discussed, but a convincing explanation was not found. Here we present a preliminary study of microstructures in three more elliptical PNe : NGC 6543, 6826 and 7009.

2 Observations The observations were performed using the Palomar 5-m telescope and double spectrograph at dispersions of 2.1 A/pixel in the blue (3400-5150 A) and 3.1 A/pixel in the red (5150-7600 A). The effective resolution corresponds to about 300 km s"1. Along the slit each pixel is 0".58 for the red and 0".78 for the blue. Typical seeing was about 1.5 arcsec. For each nebula the spectrograph slit was oriented to include a pair of microstructures on opposite sides of the central star. The location of the slit in NGC 6543, 6826 and 7009 is shown in the corresponding panels of Fig. 1, where also the nebulae studied in paper I have been reported. 291

292

M. Perinotto: FLIERs in elliptical Planetary Nebulae NGC 3242 [N II]6583A

NGC 6826 [N IIJ6583A

NGC 7009 [N II]6583A

P.A.130>

NGC 7662 [N II]6583A

low-velocity (not FL1EI

Fig. 1. CCD images in [Nil] and [OIII] (insert) of PNe of this paper and of paper I. N is at the top; size of [Nil] images is 72". Key morphological components are identified, and those that qualify as FLIERs are shown in bold italics. The position of the slit is shown for those PNe discussed here.

M. Pennotto: FLIERs in elliptical Planetary Nebulae

293

3 Physical properties of microstructures In all the three objects the microstructures (caps, ansae) exhibit much stronger lilies of the low ionization elements (Oil, SII, Nil, 01 etc.) relative to e.g. the helium lines, than the adjacent nebular regions. We measured the electron temperature and density in the microstructures and in the surrounding regions with good accuracy from the [OIII], [Nil] and [SII] line ratios. The important result is that the temperature and density are in the observed microstructures essentially similar to the values obtained in their surroundings. 4 Chemical properties of the microstructures We then measured the chemical abundances. The result was somewhat surprising. While helium, oxygen, neon, argon, sulfur and chlorine are constant across each nebula, i.e. have the same values within and outside the observed structures, nitrogen is significantly enhanced in the microstructures relative to the surroundings. We used the standard ionization correction factor icf(N + )=O/O + : We find that N/H is in NGC 6543 equal to 1.3 1(T4, 2.8 10~4, 6.0 10~4 in the North rim, cap and ansa, respectively; is 0.83 10~4 and 1.8 10~4 in the South-East rim and ansa of NGC 6826, respectively; is 1.8 10"4, 3.4 10~4 and 10~3 in the West rim, cap and ansa of NGC 7009, respectively. This is an important result for the nature of the studied microstructures. In view of the measured constant density, one might inquire why the mentioned microstructures are prominent in other low ionization lines, like the [SII] lines, whereas the abundance of sulfur appears constant across the nebulae. We think that this should be explained with the role of the ionization structure, or by the presence of shocks. 5 Kinematical Behaviour The kinematical behaviour can be studied with high resolution radial velocity measurements and realistic spatio-kinematical models of the objects. Studies of this type have been made in NGC 6543 by Miranda & Solf (1992), in NGC 6826 by Balick, Preston & Icke (1987) and Reay k Atherthon (1985) and in NGC 7009 again by Balick, Preston & Icke (1987). Although there are difficulties with the spatio-kinematic models, it is rather sure that the caps and ansae here studied are moving supersonically relative to the ambient gas. Consequently, a shocked interface between the supersonic microstructures

294

M. Perinotto: FLIERs in elliptical Planetary Nebulae

and their environments is expected. Indeed, on kinematic grounds, one expects that the supersonic microstructures are the hydrodynamic equivalent of H-H objects associated with young stars. Such objects form bow shocks. A characteristic of bow shocks is that the regions of highest ionizations lie on the upstream side of the bow. However, our observations of the ionization structure observed in the FLIERs appear to shows just the opposite! That is, the lines of lowest ionization, such as [0 I] and [S II] are found slightly further out from the central star than the [N II] and [0 II] lines. 6 Discussion and Conclusions Caps and ansae in the elliptical planetary nebulae NGC 6543, 6826 and 7009 are then low ionization pairs of high velocity microstructures located symmetrically, relative to the central star, along the major axis of the nebula. Thus they can be termed, as similar structures in NGC 3242, 7662 and IC 2149 studied in paper I, as FLIERs. Their electron temperature and density are similar to those of the surroundings areas of the nebulae. They should then have the same thermal pressure as their surroundings; however they should exert a significant kinetic (ram) pressure owing to their relative supersonic speeds. Their chemical abundances are similar to those in the surroundings, with the exception of nitrogen which is significantly enhanced. Doubts may, however, be cast. For instance the abundances have been derived under the assumption of photoionization equilibrium. But in supersonically moving structures, one expects that bow shocks develop, and therefore a contribution from shock heating within the moving features should be present. A comparison of the observed line intensities with those predicted for shocked regions in the literature was however not conclusive. The dynamical time scale of FLIERs (size/sound speed) is typically one order of magnitude greater than the hydrogen-recombination time in these features. This still allows the possibility of photoionization equilibrium, unless there is a thin shocked region. Images with sub-arcsec resolution are needed to test the shock interpretation. What about the origin of the FLIERS? Do they originate from the main structure by some hydrodynamical process in the interacting winds ? If so, two possibilities might be considered. Either they are the result of some type of local instability (then why do they come in pairs so well collimated on the two opposite sides from the central star), or they might be produced within the current interacting winds scenario with density contrast in the

M. Perinotto: FLIERs in elliptical Planetary Nebulae

295

slow wind via some collimating process (single star or double star involving a disk). Also why should they be nitrogen rich ? The alternative is a direct ejection from the central star at a 'subsequent' time relative to the ejection of the main body of the nebula. This could explain the nitrogen abundance. But what is the ejecting process ? Something occurring in a double star with a precession (to explain the twisted nature of some ansae) ? Livio (these proceedings) showed that double systems, even with a very light secondary star , are efficient in producing the density contrast needed to explain the aspherical shape of elliptical (and bipolar) nebulae. None of our three (bright) central stars is known to be a binary system. On the other hand, in the case of a quite light secondary star, and keeping the precession mechanism, we would have something like a solar system, with the rotational axis of the Sun-like star precessing with typical periods of a thousand years (the age of the nebula). Something not precisely supported by the Earth-Sun case. We think that quite a lot of work is urgently needed to explain the mysterious nature of FLIERs. References Balick, B., 1987, A.J., 94, 671. Balick, B., Preston, H.L. & Icke, V., 1987, A.J., 94, 1641. Balick, B. Rugers, M., Terzian, Y. & Chengalur, J.A., 1993, Astrophys. J., 411, 778. Miranda, L.F. & Solf, J., 1992, Astron. Astrohys., 260, 397. Reay, N.K. & Atherton, P.D., 1985, Mon. Not. R. astr. Soc, 215, 233.

Circumstellar Dust in Planetary Nebulae and Proto-Planetary Nebulae Sun Kwok1, Bruce J. Hrivnak2, and Thomas R. Geballe3 1

University of Calgary, Calgary, Alberta, Canada T2N 1N4 Valparaiso University, Valparaiso, Indiana 46383, U.S.A. 3 Joint Astronomy Center, Hilo, Hawaii 96720, U.S.A. 2

1 Introduction Over the last decade, we have come to realize that mass loss on the asymptotic giant branch (AGB) has major effects on the formation of planetary nebulae (PN), and many observable characteristics (e.g. haloes, molecular envelopes) of PN can be traced back to the circumstellar envelopes of their AGB progenitors (Kwok 1982). The large infrared excesses observed in PN are certainly due to the remnant of the AGB envelopes which have cooled as the result of expansion (Kwok 1990, Zhang & Kwok 1991). The detections of the 9.7 /im silicate and the 11.3 /<m SiC features, both commonly observed in AGB stars, provide confirmations to this link between AGB and PN (Aitken & Roche 1982, Zhang & Kwok 1990). However, the infrared spectra of PN also show features not found in AGB stars. The most prominent are the family of features at 3.3, 6.2, 7.7, 8.6, and 11.3 /xm, which are attributed to the PAH molecule (Allamandola et al. 1989). It is clear that these molecules must either be synthesized during the transition from the AGB to the PN phase, or they are produced in the AGB atmosphere but only excited in the PN environment. In either case, it would be useful to study the infrared spectra of young PN and transition objects between AGB and PN (or proto-PN) in order to understand the origin of the PAH features.

2 Infrared spectra of young PN It is well documented that the 3.3 /.im feature observed in the L band window is closely correlated with the presence of the 11.3 ^m feature observed by the IRAS LRS (Jourdain de Muizon et al. 1989), and the 6.2 and 7.7 urn features have been observed in PN from airborne observations (Cohen et al. 1989). The PAH features are particularly strong in PN with large infrared excesses and with [WC 11] central stars (Kwok et al. 1993). Example of this classes of PN are given in Table 1. The [WC] nature of the central star 296

S. Kwok et al: Circumstellar dust in PN and PPN

297

Table 1. PN with [WC 11] nuclei and large infrared excesses Name M 4-18 Vo 1 He 2-113 CPD-56 8032 IRAS 17514-1555 IRAS 21282+5050 NGC 7027

Spectral Type WC 11 WC 11 WC 11 WC 11 WC 11 WC 11 —

IR Features

(/mi)

8.6, 11.3 7.7, 11.3 3.3, 7.7, 8.6 , 11.3 3.3,7.7, 11. 3 — 3.3, 3.4-3.5, 7.7,8.6, 11.3 3.3, 3.4-3.5, 6.2,7.7,8.6, 11.3

indicates that these are carbon-rich objects, and the large infrared excesses suggest these PN have low dynamical ages.

3 Infrared spectra of PPN While the # a n d K band spectra of PPN are dominated by hydrogen recombination lines and CO vibrational- rotational lines, respectively (Hrivnak et al. 1994), the L band spectra of certain PPN show the presence of the 3.3, 3.4-3.5 /tin emission features (Geballe & van der Veen 1990, Geballe et al. 1992). The 3.4-3.5 /tm features are particularly interesting, for the only PN that are known to show these features are BD+30°3639, NGC 7027 and IRAS 21282+5050 (Geballe et al. 1985, deMuizon et al. 1986). Furthermore, the ratio of the strengths of the 3.4-3.5 /tm to 3.3 //,m features has been found to be much larger in some PPN than in PN (or HII regions). It is likely that the 3.4-3.5 /tin features are related to the unidentified emission feature at 21 /tin (Kwok et al. 1989). The 3 PPN (04296+3429,05341+0852, 22272+5435) that show the strongest 3.4-3.5 /tm emission features also possess the 21 /tin feature (Geballe et al. 1992). It should be noted, however, that the 3.4-3.5 fim features in these PPN have different spectral profiles and even peak wavelengths than the corresponding features in PN. Optical spectra have revealed that all of the PPN with the 21 /tm emission feature are F and G supergiants and possess C2 and in most cases C3 absorption (Hrivnak 1993). This indicates that they are all carbon rich. The 6.9 /tm feature, which has been attributed to CH deformation modes in aliphatic groups, is seen in 22272+5435 and also in the 21 /tm feature object 07134+1005 (Buss et al. 1990). Although it is tempting to assign a common origin to the 3.4-3.5, 6.9, and 21 /tm features, the correlation is not perfect. In Table 2 we have summarized relevant spectral properties of these sources; note that some sources have yet to be observed in the 3.3 and 6.9 /tni regions.

298

5. Kwok et al: Circvmstellar dust in PN and PPN Table 2. PPN with Peculiar Infrared Features IRAS name

Spectral Type

04296+3429 05113+1347 05341+0852 07134+1005 20000+3239 22223+4327 22272+5435 22574+6609 23304+6147

GO la G8Ia F F5I G8Ia GO la G5Ia — G2 la

Mol. Feature in Opt. Sp.

IR Features (urn)

C2,C3 C2, C3

3.3, 3.4-3.5, 21 3.3,21 3,3, 3,4-3.5, 21 3.3,6.9,21 21 21 3.3, 3.4-3.5, 6.9,21 21 21

— — Co

cC2,2 C3 —

C2, C3

4 When do the PAH molecules form? It is generally believed that the PAH features in PN and HI! regions are excited by the high ultraviolet radiation background in these objects. This is likely not the excitation source in PPN which mostly have spectral types F and G and emit primarily visible photons. Radiative transfer model fittings to the energy distributions of PPN show that these PPN have a dynamical age of only a few hundred to a thousand years (Hrivnak et al. 1989), which suggests that the PAH molecules must have been formed early in the postAGB evolution, or even in the AGB phase itself. While the 3.3 f.im feature, and to a lesser extent the 3.4-3.5 features are carried over to the PN phase, the 6.9 and 21 /im features are not seen at all in PN. It is possible that the molecules that are carrying these features are destroyed by the ultraviolet radiation or by the fast stellar winds that develop in PN central stars.

5 Speculations The detection of the 3.3, 3.4-3.5, and 21 /(in features in PPN suggests that the responsible molecules are probably synthesized on the AGB. The fact that these features are seen only in carbon-rich PN and PPN suggests that their AGB progenitors are also carbon rich. NGC 7027, one of the PN that shows the 3.3 and 3.4-3.5 /tm features, has one of the most massive central stars found among PN (Tylenda 1989). It has been suggested that 21282+5050 also has a massive nucleus (Kwok et al. 1993). Since the amount of hydrogen surrounding the core when the star leaves the AGB decreases with core mass, these WC stars may have burned up their thin

5. Kwok et al.: Circumstellar dust in PN and PPN

299

hydrogen envelopes entirely during the transition phase. The late spectral classification [WC 11] may imply that the central star has an extended atmosphere as the result of mass loss, and the low effective temperature derives from an extended pseudo-photosphere. Since PN with high mass central stars evolve across the H-R diagram much faster, they will have their circumstellar molecular/dust envelopes largely intact when their central stars have evolved to high temperatures and are emitting large uv fluxes. It is possible that the molecules responsible for these features can only survive in a thick circumstellar envelope environment, a condition that is only provided by massive PN.

References Allamandola, L.J., Tielens, A.G.G.M., fc Barker, J.R. 1989, ApJS, 71, 733 Aitken, D.K., & Roche, P.F. 1982, MNRAS, 200, 217 Buss, R.H., et al. 1990, ApJ, 365, L23 Cohen, M. et al. 1989, ApJ, 341, 246 deMuizon, M., Geballe, T.R., d'Hendecourt, L.B., & Baas, F. 1986, ApJ, 306, L015 Geballe, T.R., & van der Veen, W.E.C.J. 1990, A&A, 235, L9 Geballe, T.R., Tielens, A.G.G.M., Kwok, S., k Hrivnak, B.J. 1992, ApJ, 387, L89 Geballe, T.R., Lacy, J.H., Persson, S.E., McGregor, P.J., & Soifer, B.T. 1985 ApJ, 292, 500 Hrivnak, B. J. 1993, in IAU Symp 151: Planetary Nebulae, in press Hrivnak, B. J., Kwok, S., Volk, K. 1989, ApJ, 346, 265 Hrivnak, B.J., Kwok, S., fc Geballe, T.R. 1994, ApJ, 420, in press Jourdain de Muizon, M., d'Hendercourt, L.B., fc Geballe, T.R. 1989, in Infrared Spectroscopy in Astronomy, ESA SP-290, p. 177 Kwok, S. 1982, ApJ, 258, 280 Kwok, S. 1990, MNRAS, 244, 179 Kwok, S., Hrivnak, B.J., & Langill, P.P.L. 1993, ApJ, 408, 586 Tylenda, R. 1989, in JAU Symp 131: Planetary Nebulae, ed. S. Torres-Peimbert (Dordrecht:Kluwer), 531 Zhang, C.Y., & Kwok, S. 1990, A&A, 237, 479 Zhang, C.Y., k Kwok, S. 1991, A&A, 250, 179

HST Observations of Hydrogen-Poor Ejecta in Abell 30 and Abell 78: Evidence for Mass-Loaded Flows J. Patrick Harrington1, Kazimierz J. Borkowski1, Zlatan Tsvetanov2 and Robin E.S. Clegg3 1

Department of Astronomy, University of Maryland, College Park, MD 20742, U.S.A. 2 Center for Astrophysical Sciences, Department of Physics and Astronomy, Johns Hopkins University, Baltimore, MD 21218, U.S.A. 3 Royal Greenwich Observatory, Madingley Road, Cambridge CBS OEZ, U.K.

1 Introduction Abell 30 and Abell 78 are the best-known members of a small but important class of planetary nebulae (PNe) which are characterized by H-poor, dusty ejecta. Other members of this group include Abell 58 (V605 Aql), IRAS 18333-2357 (in the globular cluster M22) and IRAS 15154-5258. In these objects the H-poor material is surrounded by an outer envelope of normal composition (except for IRAS 18333-2357, where the ram pressure of the ISM would have stripped off the outer envelope: Borkowski et al. 1993a). Clearly, a secondary ejection of highly processed material lias occurred after the loss of the hydrogen envelope of the AGB progenitor. A detailed interpretation was put forward by Iben et al. (1983), who proposed a final helium shell flash after nearly all of the H-rich envelope had been expelled. The H-poor PNe are important because the composition of the ejecta opens a window upon the final phase of AGB nucleosynthesis and dredgeup, and also because the high dust to gas ratio lets us study the physics of dusty plasmas (e.g., gas heating by photoelectrons from grains: Borkowski & Harrington 1991). Here, however, we wish to point out that at least two of these objects also provide an exceptional opportunity to study massloaded flows. Mass-loading occurs when a tenuous, fast wind, as it streams around dense, slow-moving knots, entrains and mixes with bits of the dense material. This accelerates the dense material; it also alters the wind by slowing it and increasing its density.

2 HST Images of A30 and A78: Cometary Structures Perhaps the most compelling evidence of mass-loading is morphological. Images of A30 and A78 obtained with the HST Faint Object Camera (FOC) 300

J. P. Harrington et al: H-poor ejecta in A 30 and A 78

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Fig. 1. Inner regions of A30 and A78 in [O III] A5007. through the [0 III] A5007 filter have revealed numerous "cometary" structures (Figure 1). These structures consist of compact (0.15" - 0.5") knots with radial tails several arcseconds in length. In most cases the tails point directly away from the star, but some show significant deflections. Since the central stars of both A30 and A78 have fast (3600 km/sec) winds, we interpret the tails as material stripped from dense blobs and swept back by the winds. Non-radial tails indicate deflections of the wind flow. For more details of these observations, see Borkowski et al. (1993b). In fact, the winds have probably experienced mass-loading even closer to the central stars than the observed knots. The C IV A15-49 images we obtained show no knots but diffuse emission too intense to be produced by the original stellar wind. We have also detected extended radio emission within a few arcseconds of the central star of A30 (White et al. 1993).

3 HST Spectra of the Central Star of A78 A spectrum of the A78 central star obtained with FIST shows absorption components in the C IV A1549 line due to fast-moving nebular gas.

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J. P. Harrington et al.: H-poor eject a in A 30 and A78 121 Abell 78 Centrol Star 10

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J. P. Harrington et al.: H-poor ejecta in A30 and A78

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Figure 2 shows these features against the stellar P-Cygni profile. The absorption features are shown in more detail in Figure 3. Kaler et al. (1988) had already noted the - 2 3 km/s component of C IV and a - 7 8 km/s component of N V A1242 in WE spectra. Our HST spectrum has the great advantage of a small (0.25" x 0.25") aperture which reduces contamination by nebular emission (at the expense of the loss of half the flux due to the HST's spherical aberration). The individual components give way to a continuous absorption extending to ~ 375 km/s. It would appear that we are seeing parcels of gas from the knots being accelerated to sucessively higher velocities by the stellar wind. The distinct components at lower velocities may indicate that our line of sight crosses several "tails".

4 Summary The picture which emerges of the interaction of the stellar wind with the H-poor ejecta in A30 and A78 apparently involves an extensive mass-loading process which begins in the immediate vicinity (< 4") of the star and continues with the formation of the tails of ablated material seen in the HST [0 III] images. This mass-loaded flow ultimately blows the irregular bubbles in the H-rich envelopes which can be easily seen in the ground-based optical images of A30 (Balick 1987) and A78 (Jacoby 1979). Manchado et al. (1988) found that the H-deficient chemical composition persists out to the rim of this bubble in A78. Since the amount of mass leaving the star is small (M = 2.5 x 10~8 M© yr" 1 ), this also suggests that mass-loading has increased the amount of H-poor material carried outwards. Finally, we have constructed models of the far-IR fluxes of A30 seen by IRAS. We find that a high dust to gas ratio outside the zone of the H-poor knots is required, once again suggesting transport of dusty, H-poor material to the edge of the bubble.

References Balick, B., 1987, A. J., 94, 671. Borkowski, K.J. & Harrington, J.P., 1991, Astrophys. J., 379, 168. Borkowski, K.J., Tsvetanov, Z. & Harrington, J.P., 1993a, Astrophys. J,, 402, L57. Borkowski, K.J., Harrington, J.P., Tsvetanov, Z. &; Clegg, R.E.S., 1993b, Astrophys. J., 415, L47. Iben, I., Kaler, J. B., Truran, W. J. & Renzini, A., 1983, Astrophys. J., 264, 605. Jacoby, G.H., 1979, P.A.S.P., 91, 754. Kaler, J.B., Feibelman, W.A. & Henrichs, H.F., 1988, Astrophys. J., 324, 528. Manchado, A., Pottasch, S. R. k. Mampaso, A., 1988, A&.A, 191, 128. White, S.M., Borkowski, K.J. & Harrington, J.P., 1993, In preparation.

The evolution of the neutral envelopes of planetary nebulae P. J. Huggins 1 , R, Bachiller2, P. Cox3, T. Forveille4 1

Physics Department, New York University, New York NY 10003, USA Centro Astronomico cle Yebes, E-19080 Guadalajara, Spain 3 Observatoire de Marseille, F-13248 Marseille Cedex 4, France 4 Observatoire de Grenoble, B.P. 53X, F-38041 Grenoble Cedex, France 2

1 Introduction Rapid progress has been made over the last few years in the study of the neutral gas in planetary nebulae (PNe), and it is now well established that at least some PNe are surrounded by massive envelopes of neutral gas (see, e.g., the review by Huggins 1993). The neutral envelopes provide a new perspective on the formation and evolution of the ionized nebulae, and allow the study of a range of circumstellar processes with different characteristics than those found in other circumstellar environments. In this paper we summarize the results of a recent survey of millimeter CO emission in PNe to study the molecular component of the neutral gas, and we comment on some of the issues raised bv the observations.

2 The Molecular Gas in PNe Millimeter CO emission has proved to be an especially useful probe of the neutral gas in PNe, since it can be used to determine the structure and kinematics of the envelopes, and to estimate the mass of the molecular component. In order to systematically study the differences between PNe, particularly evolutionary effects, we have undertaken extensive survey work in the 230 GHz CO (2-1) line using the IRAM 30 m and SEST 15 m telescopes. These provide access to both northern and southern PNe, with angular resolutions of 12"-24". Our observations considerably extend the earlier survey work by Huggins & Healy (1989), and are up to a. factor of « 6 times more sensitive, depending on the angular size of the PNe. To date we have observed 93 objects in the survey and have detected circumstellar CO in 23. CO is not detected in 58 to a level of 50-100 mK (in 1 MHz filters); 8 are severely contaminated by galactic CO emission, and 3 are found to be compact. IIII regions (M 1-78, K 4-45, and He 2-77). The results for several of the detected PNe of special interest can be found in Bachiller et al. (1991), Cox et. al (1991), and Bachiller et al (1993a). Examples 304

P. J. Huggins et al.: The neutral envelopes of PNe of newly detected PNe include He2-114, Mzl, Ml-13, and NGC6853 (the Dumbbell). In combination with earlier work, about 150 PNe have now been searched in CO and 45 detected. One important application of the CO observations which we emphasize here is the estimation of the mass of the molecular component of the envelopes (see, e.g., Huggins & Healy 1989). We assume that the CO is fully associated in the molecular gas (with an abundance equal to the lesser of the C/H or O/H ratio in the ionized gas) and this provides a robust lower limit to the actual mass of molecular gas Mm. The results show that there are very large differences between PNe, with Mm ranging from w 1 M e , to ~ l x 1O~3M0. Estimates of Mm for individual PNe are rather uncertain because they depend on the square of the distances to the PNe, and these are in general poorly known. However, if one forms the mass ratio Afm/M,-, where M,- is the mass of the ionized gas, the distance dependence largely cancels, with the result that evolutionary effects and intrinsic differences between PNe become more readily apparent. Fig. 1 shows the variation of the mass ratio Mm/Mi with nebular size and illustrates quite dramatic evolutionary effects. The upper envelope of this figure indicates that there is a class of PNe which first cross the H-R diagram and form compact PNe while still surrounded by massive envelopes of molecular gas (Mm/Mi > 1). As the nebulae expand, the mass of ionized gas increases at the expense of the molecular gas; the two components are roughly comparable when the nebular radius is ss 0.1 pc, and the molecular gas becomes a minor component as the nebulae expand to still larger sizes. Fig. 1 also indicates that there are large, intrinsic differences between PNe, since there are some PNe where the molecular gas is a minor component of the system, even at relatively early stages of their evolution. Comparison with other properties of the PNe indicates that this is correlated with the mass of the progenitor stars. Type I, bipolar, and B type PNe are well represented among those with massive molecular shells, but type C PNe are not. These differences are possibly due to differences in the mass loss rate of the progenitors, the luminosity of the central stars, the evolutionary time scales, or the tendency to form dense condensations, or a. combination of these. Further work is needed to clarify this aspect of the observations. 3 Clumps and the Atomic-Molecular Transition The presence of molecular gas in PNe raises a number of interesting issues, and we comment briefly on two of them here. The first concerns the survival of molecules in the neutral gas. In homogeneous, expanding envelopes the

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•

-

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0.01 =

•

0.001

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gas rapidly becomes diffuse and optically thin, and molecules will be destroyed on a very short time scale by the intense ultraviolet radiation field of the central star. The survival of substantial amounts of molecular gas is only possible with sufficient shielding, and this means that the gas must be highly clumped. In the best studied PNe there is good observational evidence that this is indeed the case. We have recently mapped the compact PN NGC7027 in the 89 GHz H C 0 + (1-0) line with an angular resolution of 2.5" using the Plateau de Bure interferometer (Cox et al. 1993), and find the gas forms a highly clumped envelope around the IIII region; comparison of H C 0 + lines from different upper states also shows that the gas density in the clumps is extremely high (~ 106 cm" 3 ). Similarly, in the much more evolved Helix nebula, the numerous cometary globules around the periphery and within the ionized cavity are found to have dense cores of molecular gas (Huggins et al. 1992). Thus, the clumping of the gas appears to be a fundamental feature of PNe with massive shells of neutral gas. A closely related issue is the distribution of neutral atomic and molecular gas in the envelopes. According to model calculations of photo-dissociation

P. J. Huggins et al.: The neutral envelopes of PNe regions there is expected to be a transition from largely atomic gas adjacent to the ionized regions to molecular gas deeper within the condensations (e.g., Tielens 1993). HI and H2 have been detected in a number of PNe (e.g., Taylor et al. 1990, Webster et al. 1988), but it has not been possible to investigate the HI-H2 transition because of limitations in the data. The CII-CI-CO transition should eventually prove to be a more amenable case. The CO data base is now well established. Fine structure CII emission has been detected in only a few PNe (Ellis & Werner 1985, Dinerstein 1991), but should be widely observable with the ISO satellite, and important progress has very recently been made in the measurement of CI in PNe. Using the Caltech Submillimeter Observatory, we have detected the 609 /mi (492 GHz) 3 Pi - 3 Po transition of CI in the Ring Nebula, NGC6720 (Bachiller et al. 1993b), and Young et al (1993) have detected it in NGC7027. The observations of the Ring are of special interest because of the large amount of CI observed: the number of CI atoms exceeds the number of CO molecules by a factor of about 4. The observations confirm that the mass of the neutral envelope in the Ring exceeds that of the ionized nebula, and indicate that much of the neutral gas resides in the atomic transition region. Further observations of these carbon species should prove valuable in building a detailed picture of the physical state of the gas in the evolving neutral envelopes. Acknowledgements This work was supported in part by a grant from the NSF (to P.J.H.).

References Bachiller, R., Huggins, P. J., Cox, P., Forveille, T., 1991, Astron. Astrophys., 247, 525. Bachiller, R., Huggins, P. J., Cox, P., Forveille, T., 1993a, Astron. Astrophys., 267, 177. Bachiller, R. et al. 1993b, Astron. Astrophys., (in preparation). Cox, P., Huggins, P. J., Bachiller, R., Forveille, T., 1991, Astron. Astrophys., 250, 533. Cox, P., et al. 1993, Astron. Astrophys., (in preparation). Dinerstein, H. L., 1991, Pub. Astron. Soc. Pacific, 103, 861. Ellis, H. B., Werner, M. W., 1985, In Mass Loss From Red Giants, eels. M. Morris & B. Zuckerman, Reidel, Dordrecht, p. 309. Huggins, P. J., 1993, In Planetary Nebulae, eds: R. Weinberger & A. Acker, Kluwer, Dordrecht, (in press). Huggins, P. J., Healy, A. P., 1989, Astrophys. J., 346, 201. Huggins, P. J., Bachiller, R., Cox, P., Forveille, T., 1992, Astrophys. J., 401, L43. Taylor, A. R., Gussie, G. T., Pottasch, S. R., 1990, Astrophys. ./., 351, 515. Tielens, A. G. G. M., 1993, In Planetary Nebulae, eds: R. Weinberger & A. Acker, Kluwer, Dordrecht, (in press). Webster, B. L., Payne, P. W., Storey, J. W. V., Dopita, M. A., 1988, A/on. Not. R. Astron. Soc, 235, 533. Young, K., et al. 1993, Astrophys.

J., (in p r e p a r a t i o n ) .

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Magnetic Shaping of Planetary Nebulae Roger A. Chevalier Department of Astronomy, University of Virginia, P.O. Box 3818, Charlottesville, VA 22903, USA

Abstract As in the case of the solar wind, the magnetic field in the wind from a magnetized, rotating star becomes increasingly toroidal with distance from the star. The strength of the magnetic field can he characterized by IT, the ratio of toroidal magnetic energy density to kinetic energy density in the equatorial plane of the wind. A fast wind shocks against the external medium and creates a bubble whose volume is dominated by shocked gas. The toroidal magnetic field increases in the shocked bubble and can dominate the thermal pressure. Because of the low velocities in the bubble, hydrostatic equilibrium is a good approximation and allows the calculation of the thermal and magnetic pressure in the bubble, as in the model of Begelman &; Li (1992) for the Crab Nebula. The pressure is asymmetric because magnetic tension constrains the flow in the equatorial direction and there are no magnetic effects in the polar direction. The total pressure drives a shell into the surrounding medium, which can be treated in the axisymmetric "thin shell" approximation. If the fast wind is running into a slow wind from a previous evolutionary phase, the interaction shell tends toward motion at constant velocity and the shell structure varies only with polar angle. The structure, which is axisymmetric and extended in the polar direction, depends on 2 parameters: avw/wo, where vu, is the wind velocity and wo is the shell velocity in the polar direction, and A = va/wo, where t>n is the velocity of the slow wind. For small values of A, there is a cusp in the shell in the equatorial plane, i.e. there is an equatorial ring. For larger values of A, the maximum of the surface density moves away from the equator, i.e. a double ring structure. The models should apply to planetary nebulae, if their central stars are sufficiently magnetized and are rotating sufficiently rapidly; the calculated shapes do resemble the observed shapes of planetaries. The model can also be applied to the nebula around SN 1987A, using a low A model with an equatorial ring; the recent, rise in the radio flux from SN 1987A may be related to interaction with a magnetized bubble. In all these cases, the model predicts that X-ray emission from the bubble is concentrated toward the polar axis. Details of the models can be found in Chevalier fc Luo (1994).

References Begelman, M. C. & Li, Z.-Y., 1992, Astrophys. J., 397, 187-195. Chevalier, R. A. & Luo, D., 1994, Astrophys. J., 421, 222-235.

308

The hydrodynamics of aspherical two-wind configurations Vincent Icke Sterrewacht Leiden, Postbus 9513, 2300 RA Leiden, The Netherlands

Abstract Planetary nebulae (PN) are bubbles blown by a tenuous, fast stellar wind into a dense, slow, fossil red giant envelope (RGE). This interactingwinds model is quite complete in spherical geometry (Castor et al. 1975, Weaver et al. 1977, Kwok 1982, Lamers 1983, Kahn 1983). New developments in this sector will be treated by others in this volume. Therefore, my aim is to present some critical remarks which are mainly relevant to aspherical interacting winds and in particular cylindrical ones.

1 Analytical beginnings Balick (1987, 1988) suggested that the interacting-winds model for planetary nebulae, if generalized to two dimensions, might explain the morphologies of nearly all PNs. He supposed that the fossil red giant envelope (RGE) is cylindrically symmetric and that the density is higher at the equator than at the poles. The PN morphology and its evolution is then purely a consequence of the mass distribution in the RGE and the properties of the fast stellar wind. Analytical models for aspherical PNs have been quite successful in describing the propagation of the outer shock of a two-wind configuration. One uses either a snowplow-type approximation (Kahn & West 1985, Soker & Livio 1989) or a generalization of the work of Kompaneets (1960; Balick et al. 1987, Icke 1988, Icke et al. 1989), in which the wind is supposed to generate a uniform pressure inside the expanding bubble. The method is easy to apply and shows the generic features of the outer shocks clearly. It works as follows. The propagation of the strong shock that moves into the RGE is governed by one parameter only, namely the ratio of the pressure Pi behind the shock and the RGE density po ahead of it:

ff = ! ± ! £

(i)

Here 7 is the adiabatic index of the gas; Pi is primarily due to the dynamic pressure of the wind. I adopt spherical coordinates (?•, 0,4>), but suppress the dependence on

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deposited by a stationary supersonic wind, so that to a good approximation P* The function r(6, t) describing the shape of the shock front is a solution of 8r so that

(dx/dO)2} ; x = log(r/r0) ; r = ^ Due to the inverse-square behaviour of the RGE gas density, we have obtained an equation for the propagation of the shock that is particularly easy to solve. Eq.(4) can be decomposed by means of separation of variables: = ET

- J ^E*/A{9)

-lde

+ g(E)

(5)

Here E is an arbitrary constant and g(E) an arbitrary function thereof. From this 'complete integral' we may obtain other solutions of Eq.(4) by taking the envelope of all functions in Eq.(5) (Courant & Hilbert 1968, II , Ch.I §4). If the fast wind is spherically symmetric x must be constant (which we may take to be zero) at r = 0 for all values of E\ therefore g(E) = 0. Eq.(5) can be evaluated by quadrature for any desired value of E (Simpson's Rule is accurate enough). Because it is unknown how the RGE is deposited, and a fortiori how it acquired its flattening, one has to guess what A is; I took A{B) = l-a

+ a exp[/J cos(20) - /?]

(6)

Here a and (3 are constants that describe the flattening of the RGE and the steepness of A(0). Fig.la shows the evolution of shock fronts travelling through the density distributions prescribed by Eq.(6) . Note that the shocks all evolve towards a fixed shape in a finite time, typically of the order of r = 4. This behaviour is generic and is due to the fact that the values of E always span a finite range for a bounded function A{9). These computations are easy and illuminating when implemented interactively on a Macintosh or a workstation (I will be happy to share code with anyone seriously interested). The work by Kahn & West (1985) is somewhat more involved but also lends itself to straightforward applications; the results are quite similar.

V. Icke: Aspherical two-wind configurations 30

1

i

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Fig. 1. (a) The evolution of the outer shock moving into a flattened RGE (Icke 1988). The left edge of the frame is the symmetry axis, the bottom edge is the equatorial plane. Note the formation of a cusp at about 0 = 45° at time r = 1.5; an equatorial cusp forms at T ~ 4.5. (b) Focusing by the inner shock. Top left panel shows the spherical case (no deflection); the other three panels summarize various configurations that occur in the numerical computations.

Below /3 ~ 2 the outer shock ranges from elliptical to peanut shaped. When /? > 3 the shock becomes bipolar but is never collimated well. The shock is a superposition of partial waves with different speeds. If /? J> 2 the fastest partial wave overtakes all others and the superposition forms a cusp. For most plausible RGE models this first occurs at intermediate latitudes, 0 « 40°. Later, a cusp also appears in the equatorial plane (Fig.la). The RGE gas immediately behind a cusp is compressed. The cusps show up as rings: in the first case, a pair of rings above the plane (e.g. in the Helix Nebula and the Owl); in the second case, a ring in the plane (as in the extreme butterflies). The shapes of bipolar nebulae are reproduced very well (Balick & Preston 1987, Balick et al. 1987, Mellema. 1993a, Frank 1993, 1994), even in extreme cases (Calabash Nebula: Morris & R.eipurth 1990, Icke & Preston 1989; note the evidence for an axial jet). The pressure in the hot bubble is almost uniform so that the shape of the outer shock communicates itself to the inner (standoff) shock: the inner shock has an elongation in the same direction as the outer one. This implies that the flow from the central star is focused along the axis (Eichler 1982, Icke 1988). The focusing patterns can become quite complex (Fig.lb) and are expected to influence the structure and evolution of the contact discontinuity.

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2 Numerical progress Early attempts at numerical modelling (Livio & Soker 1988, Soker & Livio 1989, Soker 1989) did not immediately provide correct answers because the density jump between the slow and the fast wind is of the order of 3000 and velocity jumps may exceed 200. This demands extraordinary caution in numerical work. Soker & Livio (1989) certainly understood the physical problem, but their numerical technique did not match that understanding. Residual numerical diffusion due to second-order errors produces unphysical results: the outward flow from the central star is smothered by the spurious inward momentum flux of the huge contact discontinuity between the fast wind and the RGE (Icke 1991). This is all the more important because such filling in does really occur under some circumstances (see Sect.3 below). A word to the wise about good numerical technique: if you are interested in making real models instead of ad libitum sketches on overhead transparencies, fine; but leave the design of hydrocodes to experts, just like you would the building of a telescope. Excellent numerical schemes exist (e.g. Woodward & Colella 1984, Boris & Book 1973, 1976, Boris et al. 1975, Zalesak 1979, Van Leer 1984, Roe 1986, Icke 1991, Eulderink 1993, Eulderink & Mellema 1994). Cadge a code from someone else or use their results for post-processing (see Sect.6). The interacting winds produce the usual triple discontinuity: an inner (reverse) shock surrounding the central source, an outer shock propagating into the RGE and a contact discontinuity separating the two. Both shocks confirm the analytic expectations within about 5% (Mellema et al. 1991, Icke et al. 1992a). The value a of the equator-to-pole density contrast mostly determines the timescale on which a strongly aspherical outflow pattern is established. The contact discontinuity and the flow behind the leading shock show two types of patterns, depending on (3. At (3 & 2 one finds elliptical to peanutshaped outer shocks, inside of which the aspherical structure is hardly apparent. The flow behind the outer shock on the symmetry axis has the appearance of an ascending vortex or mushroom cloud. Values f3 J> 3 describe funnel-like RGE's which produce elliptical inner shocks, with jet-like structures extending above them. These form quickly and flow almost parallel along the axis, finally protruding into a cap formed by the outer shock. The autocollimation arises as dense gas is entrained into the flow at the point where the outer shock makes a sharp inward cusp. This cool gas forms a chimney which confines the hot gas near the axis. At the head of the jet the gas forms a backflow which helps to stabilize it against the

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1. fast wind 2. inner shock 3. hot bubble 4. contact discontinuity (equatorial belt) 5. outer shock (equator) 6. contact discontinuity (chimney) 7. jet 8. X-shock 9. terminal shock and cap bubble 10. outer shock (cap) 11. streamline 12. backflow 13. slow wind (RGE) 14. cusp compression (ring above equator) Fig. 2. Numerical results of a shock propagating in an RGE with a = 0.8, 0 - 3, at time /. = 1.4 x 10 10 sec. RGE density p = 5 x 10~ 17 kg m~ 3 , speed v = 15 km/s, temperature T - 50 K. The fast wind has p = 10~ 20 kg m~ 3 , v = 2500 km/s, T - 105 K. The frame measures 6 by 2.1 x 10 15 m, 600 x 180 cells, and shows logarithmic greyscalejof the den-3 sity. Black is highest, at 10~ 15 , white is lowest at 10~ 20 kg m" "

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pressure in the jet (Fig.2). The jet is occasionally crossed by an X-shock of the type seen in cylindrical jets which emanate from a perfectly collimated orifice (Norman et al. 1983, Wiita 1991). In our case, however, the flow itself does the collimating through the combination of inertial confinement of the chimney and the pressure of the backflow gas. The mechanism is quite different from that introduced by Blandford & Rees (1974). If a chimney develops and the fast gas is collimated by inertial confinement, the flow in the resulting jet is usually extremely parallel. Deviations from perfect collimation are due to the small ripples in the contact discontinuity and due to X-shocks near the jet axis. Such shocks must occur when oblique shocks try to cross the axis (Courant & Friedrichs 1976, p.387). The configuration is like the letter X with a small transverse bar (a 'Mach disk') across the jet axis, forming an intermediate between the letters X and H. It is not clear how these shocks could be observed. In principle they could be detected by their characteristic velocity pattern. For applications to cooler jets, where line emission from X-shocks might in fact occur, I will sketch the basic procedure. The incoming flow with velocity v\ hits a shock which makes an angle

- 1

(7)

where M\ is the jet Mach number and 7 is the adiabatic index. From Eq.(7) and the Hugoniot relations one finds the jump conditions for v, the density p, the pressure P and M. When these are given in the incoming jet, one finds the conditions across the X by repeated application of these equations. A strong X-shock occurs at the end of the jet where the hot gas deflects off the expanding RGE cap. Beyond the Mach disk and inside the slip layer the speed is lower, and the cap gas buckles inward. This has been seen numerically in our own work (e.g. Mellema 1993a) and in jet calculations (Norman et al. 1983). The shapes are highly dynamical but the dimpled appearance recurs. The outflow around such recurved features is reminiscent of the 'fishtails' (NGC6396), but it remains to be seen if this identification is correct. The contact discontinuity is usually corrugated. This leads to the formation of a mixing zone between the hot bubble gas and the RGE material. The hot bubble inside is almost unobservable, but X-rays from the mixing zone might be detectable.

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3 Including more physics The luminosity of planetary nebulae does more than make them beautiful in the sky: radiation is important in their evolution, too. Therefore, the most urgent item to include is a radiative exchange term in the energy equation. Momentum deposition is of lesser concern, although it is likely to be important as line pressure in massive young stars and in active galactic nuclei, and pressure on dust in very young PNs and in young low-mass stars. It is easy to incorporate radiative transfer after a purely hydrodynamical computation run, because the radial-only approximation (Icke, Gatley & Israel 1980) can be used here, even in the presence of dust. This is particularly easy by means of the analytical line emissivity functions introduced by Balick et al. (1993). This post-processing does not affect the dynamics and can only be used as a first guess of what the PN models would look like. The central stars of young PNs evolve on a timescale comparable to the hydrodynamic evolution time of the two-wind shells. This co-evolution affects the global behaviour of the bubble-shell configuration. Breitschwerdt & Kahn (1988, 1990, Kahn & Breitschwerdt 1990a,b) noted in their study of spherical bubbles that the switch-on of the central wind should result in two regimes. If the gas density of the central wind is high ('momentum-driven' wind), the Mach number is small, the inner shock is weak and the bubble gas can cool promptly. In the converse case ('energy-driven' wind) the inner bubble remains at coronal temperatures. Furthermore, if the ionizing flux from the central star switches on rapidly enough, it can catch the two-wind nebula in the momentum-driven phase; in that case, Breitschwerdt and Kahn found that the contact discontinuity starts to evaporate. Because of the enormous density jump across the discontinuity the momentum flux in this evaporation pushes the gas inward, smothering the central wind. They predicted that this should lead to instabilities. This analysis is not applicable beyond the onset of instability. In higher dimensions the spatial exchange of fluid elements allows the formation of extremely complex structures. Therefore, the later behaviour needs explicit calculation. This was done in the numerical work of Mellema (1993a,b,c, 1994) and Frank (1993, 1994, Frank et al. 1992, 1993, Frank k Mellema 1994a,b). They found that the change-over between the two regimes did occur near a wind speed of 150 km/s, as expected. Radiation deposits energy in selected zones, due to atomic properties of lines and ionization edges. Thus, these quantummechanical jumps are expressed as spatial discontinuities by the radiative transfer.

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Mellema and Frank also studied the evolution of the ionization front that produces a preheating wave as it travels through the incipient PN. The shock preceding this ionization front is followed later by a second shock which is caused by the impact of the energy-driven fast wind. The heat deposited by the ionization front in the R.GE causes the pressure there to increase so much that the expansion of the contact discontinuity is halted or even reversed; the central bubble fills up in about 2000 years unless the central momentum density is sufficiently high and timely. This interplay can lead to a large variety of cases, due to the fact that the various timescales are all of the same order of magnitude. I paraphrase Mellema (1993a, Chapter 4) on the behaviour of radiative spherical models: "The first outer shock can be identified with the edge of the surrounding envelope in PNs with attached envelopes. All PNs should at some stage have such an envelope. The second outer shock can be identified with the bright inner rim. Envelopes with sharp edges do not necessarily indicate a variable slow wind, but the correspondence with observations is better for a time dependent RGE. The outer envelope may expand faster than the rim. Since the envelope velocity stays roughly constant and the rim accelerates, faster envelopes are mainly expected in young nebulae. This is confirmed by observations (Chu 1989). Because the expansion of the rim stalls during the time of ionization of the slow wind, the dynamical age of nebulae in the models falls behind their evolutionary age: PNs are generally too small for their expansion age. Too low mass loss rates or velocities in the fast wind lead to the collapse of the nebula when ionization sets in. Fast wind mass loss rates lower than 10~9 M^yr" 1 at the moment of ionization are ruled out, except for centrally condensed nebulae like IC4593 and NGC6210."

The numerical findings in the aspherical case are interesting extensions of this. Again I paraphrase Mellema (Chapter 5): "The aspherical interacting winds model can explain the emission line images, the long slit spectra and the shapes: butterfly, rectangular, elliptical, barrel-shaped and round, as well as some of the substructures (eyes, major axis extrusions, bright edges). The models confirm the non-radiative ones, adding the new aspects of image substructures and kinematic structure. In radiative models the shock strength does not always increase; it may even become less than the value for non-radiative shocks. This implies that the swept-up shell is not always more compact than under non-radiative conditions! The inner shock causes turbulence in the hot bubble, disturbing the velocity field. This affects the shape of the inner shock but not that of the swept-up shell."

V. Icke: Aspherical two-wind configurations

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4 Puzzles

The enigmatic objects which Balick and co-workers have called FLIERS (fast low-ionization emission regions) are the main challenge to those who would understand PN hydrodynamics. The ansae (fast knots which show up in N + but hardly at all in neutral hydrogen) have been known for a long time. Cooling processes near the axis might produce them, but so far they have failed to show up (Mellema 1993a). They usually have an antisymmetric symmetry with respect to the central star and their spectroscopic properties show prima facie evidence for a vast nitrogen overabundance (cf. Perinotto in this volume). The antisymmetry leads naturally to the conjecture that they were formed by some sort of off-axis motion near the inner parts of the RGE disk. But why the bizarre composition? Theflowthrough flattened RGEs reproduces the PNs so well that I cannot believe that anything else might be the cause of the bipolar phenomenon. But why is the confining nebula flattened? The only mechanism that has been worked out in some detail is the common-envelope binary capture (Livio et al. 1979, Bodenheimer k Taam 1984, Taam k Bodenheimer 1989, Soker 1990; see also Soker's contribution in this volume). Paraphrasing Livio's contribution in this volume: we know that a star must lose lots of angular momentum when it forms. This remains stored in the matter of a disk, a planetary system or a companion star. Later, when the star begins to swell up, it may pick up some of that angular momentum again, which has been put into the deep freeze for billions of years. The result may be a flattened stellar envelope. It remains to be seen if it works; the statistics is tricky. 5 Getting ambitious Applications of the interacting-winds formalism to flow regimes other than those that occur in PNs include the wind-ring interaction in SN 1987A (cf. Blondin in this volume). To show that the PN formalism can include such exotic cases, the reader is encouraged to apply the technique of Sect.l above to the density distribution of Eq.(2) , in which the angular function A describes a disk: A = (1/6) (a + (9 - TT/2) 2 ). Note especially the formation of cusps and compare this with Blondin's diagrams. Another interesting example where biconical flow might occur is the Homunculus of rj Carinae (Icke 1981). The 'inertia! confinement' mechanism for the formation of gaseous jets should apply to all situations that are geometrically similar to the PN case, ranging from young stars and Herbig-Haro objects to active galactic nuclei

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(Icke et al. 1992b, Frank & Noriega-Crespo 1994). In the relativistic case the collimation gets even better (Eulderink 1993, Eulderink & Mellema 1994). The application of the two-wind mechanism to young stellar objects requires a modification of the equations of motion, because radiation pressure on dust becomes important. Also, energy losses due to cooling and (photo)dissociation need to be considered. Optical depth effects make life extremely difficult, because intrinsic nonlinearities (ionization, dust opacity, line opacity) cause the photon momentum to couple into the motion in curious and non-intuitive ways. 6 Use and abuse of models There are two enormous obstacles to the comparison between theory and observations. First, there is a pronounced tendency among observers to publish 'connect-the-dots' drawings rather than the spectra themselves, Please don't! We need to compare our theories with the real stuff. Second, hydrocodes produce a prodigious amount of output, and making this generally available is no simple thing. We are trying to improve on this by making our results available on Internet via 'anonymous ftp'. Please contact me, Frank or Mellema if you are interested. Too frequently, even in refereed papers, do people publish back-of-theenvelope scribbles called 'models'. I will not embarrass the individual authors, and the stack would be too high anyway, but stop arbitrarily sketching shapes and guessing velocities and densities! To encourage everyone to make actual first-order models, rather than resort to the magic-marker hydrodynamics which some people extemporize with felt tip pens on overhead transparencies, I have presented in Sect.l a capsule summary of the way in which simple but meaningful hydro can be done analytically. Acknowledgements It is an honour and a privilege to thank the friends-and-colleagues who have been involved in this project for almost a decade: Bruce Balick, Garrelt Mellema, Adam Frank and Frits Eulderink. I owe a debt of gratitude to Mario Livio and Noam Soker for spirited discussions, and to Hugo Schwarz for keeping everybody's nose to the ground. Most of the research summarized above was supported by the Netherlands Council for Supercomputer Facilities (NCF); by the Netherlands Foundation for Research in Astronomy (ASTRON); by the North Atlantic Treaty Organization (NATO) under grant 0898/87; and through the U.S. National Science Foundation (NSF) un-

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der grants AST89-13639 and INT87-14186. I apologize to the participants of the Conference, and to its organizers, for being so scandalously late with my contribution.

References Balick, B., 1987, Astron. J. 94, 671. Balick, B., 1988, Astron. J. 97, 476. Balick, B., Mellema, G., k Frank, A., 1993, Astron. Astrophys. 275, 588. Balick, B., k Preston, H.L., 1987, Astron. J. 94, 958. Balick, B., Preston, H.L., k Icke, V., 1987, Astron. J. 94, 1641. Blandford, R.D., k Rees, M.J., 1974, Monthly Notices Roy. Astron. Soc. 169, 395. Bodenheimer, P., k Taam, R.E., 1984, Astrophys. J. 280, 771. Boris, J.P., k Book, D.L., 1973, J.Computat.Phys. 11, 38. Boris, J.P., k Book, D.L., 1976, J.Computat.Phys. 20, 397. Boris, J.P., k Book, D.L., Hain, K., 1975, J.Computat.Phys. 18, 248. Breitschwerdt, D., k Kahn, F.D., 1988, Monthly Notices Roy. Astron. Soc. 235, 1011. Breitschwerdt, D., k Kahn, F.D., 1990, Monthly Notices Roy. Astron. Soc. 244, 521. Castor, J., McCray, R., k Weaver, J., 1975, Astrophys. J. (Letters) 200, L107. Chu, Y., 1989, In: Planetary Nebulae, IAU Symposium 131, S. Torres-Peimbert (Ed.). Courant, R., k Friedrichs, K.O., 1976, Supersonic Flow and Shock Waves (Springer, Berlin). Courant, R., k Hilbert, D., 1968, Methoden der mathematischen Physik (Springer, Berlin). Eichler, D., 1982, Astrophys. J. 263, 571. Eulderink, F., 1993, Numerical relativistic hydrodynamics, Ph.D. Thesis, Leiden. Eulderink, F., & Mellema, G., 1994, Astron. Astrophys., in the press. Frank, A., 1993, In: Planetary Nebulae, IAU Symposium 155, R. Weinberger k A. Acker (Eds.), p.311. Frank, A., 1994, Astron. J. 107, 254. Frank, A., Balick, B., Icke, V., k Mellema, G., 1993, Astrophys. J. 404, 125. Frank, A., k Mellema, G., 1994a, Astron. Astrophys., in the press. Frank, A., & Mellema, G., 1994b, Astrophys. J., in the press. Frank, A., k Noriega-Crespo, A., 1994, Astron. Astrophys., in the press. Frank, A., Noriega-Crespo, A., k Balick, B., 1992, Astron. J. 104, 1120. Icke, V., 1981, Astrophys. J. 247, 152. Icke, V., 1988, Astron. Astrophys. 202, 177. Icke, V., 1991, Astron. Astrophys. 251, 369. Icke, V., Balick, B., k Frank, A., 1992a, Astron. Astrophys. 253, 224. Icke, V., Gatley, I.A., k Israel, F.P., 1980, Astrophys. J. 236, 808. Icke, V., Mellema, G., Balick, B., Eulderink, F., k Frank, A., 1992b, Nature 355, 524. Icke, V., k Preston, H.L., 1989, Astron. Astrophys. 211, 409. Icke, V., Preston, H.L., k Balick, B., 1989, Astron. J. 97, 462. Kahn, F.D., 1983, In: Planetary Nebulae, I.A.U. Symposium 103, edited by D.R. Flower (Reidel, Dordrecht), p. 305. Kahn, F.D., k Breitschwerdt, D., 1990a, Monthly Notices Roy. Astron. Soc. 242, 209. Kahn, F.D., k Breitschwerdt, D., 1990b, Monthly Notices Roy. Astron. Soc. 242, 505. Kahn, F.D., k West, K.A., 1985, Monthly Notices Roy. Astron. Soc. 212, 837. Kompaneets, A.S., 1960, Doklady Akad. Nauk SSSR 130, 1001. Kwok, S., 1982, Astrophys. J. 258, 280. Lamers, H.J.G.L.M., 1983, In: Diffuse matter in galaxies, J. Audouze (ed.), Reidel, Dordrecht, p.35 Landau, L.D., k Lifshitz, E.M., 1987, Fluid mechanics (Pergamon, Oxford).

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Livio, M., Salzman, J., &; Shaviv, G., 1979, Monthly Notices Roy. Astron. Soc. 188, 1. Livio, M., & Soker, N., 1988, Astrophys. J. 329, 764. Mellema, G., 1993a, Numerical models for the formation of aspherical planetary nebulae, Ph.D. Thesis, Leiden. Mellema, G., 1993b, In: Mass Joss on the AGB and beyond, H.E. Schwarz (Ed.), ESO, Garching, p. 13. Mellema, G., 1993c, In: Planetary Nebulae, IAU Symposium 155, R. Weinberger b. A. Acker (Eds.), p.369. Mellema, G., 1994, Astron. Astrophys., in the press. Mellema, G., Eulderink, F., &; Icke, V., 1991, Astron. Astrophys. 252, 718. Morris M., & Reipurth B., 1990, Publ. Astron. Soc. Pacific 102, 446. Norman, M.L., Winkler, K.H.A., k. Smarr, L.L., 1983, In: Astrophyskal Jets, A. Ferrari & A.G. Pacholczyk (Eds.), Astrophys. Space Set. Lib. 103, Reidel, Dordrecht, p.227. Roe, P., 1986, Ann. Rev. Fluid Mech. 18, 337. Soker, N., 1989, Astrophys. J. 340, 927. Soker, N., 1990, Astron. J. 99, 1869. Soker, N., & Livio, M., 1989, Astrophys. J. 339, 268. Taam, R.E., & Bodenheimer, P., 1989, Astrophys. J. 337, 849. Van Leer, B., 1984, SJAM J.Scient.Statist.Computat. 5, 1. Weaver, J., McCray, R., Castor, J., Shapiro, P., & Moore, R., 1977, Astrophys. J. 218, 377. Wiita, P.J., 1991, In: Beams and Jets in Astrophysics, P.A. Hughes (Ed.), Cambridge University Press, Chapter 8. Woodward, P., & Colella, P., 1984, J. Computat. Phys. 54, 115. Zalesak, S.T., 1979, J.Computat.Phys. 31, 335.

Novae and related stars as tracers of mass loss Michael Bode School of Chemical and Physical Sciences, Liverpool John Moores University

Abstract In this review we consider the ways in which novae and related objects can be used to give insights into mass loss from the evolved stellar components of these interacting binaries, or from previous phases of binary evolution. We do not concern ourselves with the processes of mass loss at outburst per se. We pay specific attention to symbiotic stars, recurrent novae and classical novae.

1 Introduction Of the three sub-types of interacting binary considered in this review, there is no doubt that classical novae (CN) are the best defined at present in terms of our knowledge of the composition of the central binary (white dwarf plus late-type main sequence star) and cause of outburst (thermonuclear runaway - see e.g. Bode and Evans 1989 and references therein). Recurrent novae (RN) on the other hand form a small, and surprisingly heterogeneous group of nine known members, with either red giant or main sequence massdonating stars and either white dwarf or main-sequence accretors (see e.g. Bode 1987, Webbink et al. 1987). The much larger class of symbiotic stars (SS) is equally heterogeneous, and as with RN, the cause of outburst is less clear than for classical novae, though it seems that the evidence in favour of most of these systems containing a white dwarf is increasing (e.g. Miirset et al. 1991). What has been clear since the class was first defined is that their cool components are evolved. We now proceed to discuss each of these types of object in turn.

2 Symbiotic Stars Symbiotics are divided into two sub-types: 5 with near IR colours those of a stellar photosphere, and D where this wavelength range is dominated by emission from hot dust (Webster & Allen 1975). The former contain red giant cool components, whereas the latter contain Mira variables. In both cases, the radiation field of the hot, accreting star ionises and excites the wind of the cool component giving rise to the emission lines we observe in the optical and UV, and the free-free continuum emission seen at longer wavelengths. In some cases, such emission may also be associated with 321

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material ejected in a slow nova eruption, or result from interaction of the slow wind from the red giant and a fast wind from the hot component. The circumstellar environment of SS has been explored observationally in several ways. For example, Nussbaumer et al. (1988), using emission lines, showed that the ionised winds of the late-type components of SS have similar CNO abundances to those of isolated red giants. However, one of the most fruitful lines of investigation in this regard has been to observe free-free emission at radio wavelengths. In a series of papers, E.R. Seaquist (Toronto), A.R. Taylor (Calgary) and their collaborators have shown that SS are among the most prolific stellar radio sources. In their latest contribution (Seaquist et al. 1993) 99 symbiotics were observed with the VLA at 3.6cm. Over half were detected at the O.lmjy level and above. They found, as one might expect, a good correlation between radio and hot dust emission, but a surprisingly poor one with cool dust. Mass-loss rates were found to increase with later spectral type, but these were generally higher in SS than in local field giants. Whitelock & Munari (1992) ha.ve suggested that the red giant components of SS are more akin to those of the Galactic bulge population, which might go some way to explaining this discrepancy. In general, D type SS are found to suffer higher mass loss rates and are more radio luminous than 5 types. Ivison et al. (1991) had previously concluded from a combination of optical spectroscopy and radio observations that D types have wider binary separations than S types. This makes sense in terms of the relative physical sizes of red giant and Mira variable stars. The work of Ivison et al. (1991) led in turn to their extensive observations of SS at mm/sub-mm wavelengths with the JCMT (Ivison et al. 1992). Figure 1 illustrates the importance of these observations in determining the free-free spectral turnover, the position of which can be used to test models of the radio-emitting region such as that proposed by Seaquist et al. (1984). 3 Recurrent Novae Recurrent novae are sometimes (perhaps mistakenly) included in lists of SS. In fact, they generally only show high excitation emission lines at outburst, and not all contain red giants. One of the best studied members of this class is RS Oph which has had recorded outbursts in 1898, 1933, 1958, 1967 and 1985. The 1985 outburst was studied across the electromagnetic spectrum from the radio to the X-ray region (Bode 1987 and references therein). Prior to the latest eruption, there was strong evidence that at outburst high velocity ejecta interact with a slow wind from the M2III star. Such

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M.F. Bode: Novae as tracers of mass loss

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evidence includes a pronounced narrowing of emission lines with time after outburst, signifying deceleration of ejecta (Dufay et al. 1964), and the emergence of very strong optical coronal lines (up to [FeXIV], Joy 1961). Indeed, it was in the 1933 outburst that coronal lines were first identified in any object other than the Sun. When we were notified about the latest outburst in 1985, a major multifrequency observational campaign was set in motion. Among the most important new observations that were secured were those in the X-ray (using the EXOSAT satellite, Mason et al. 1987 - see figure 2) and the radio (Padin et al. 1985, (Hjellming et al. 1986), the latter culminating in European VLBI Network (EVN) observations (Taylor et al, 1989) which gave a stunning insight into the morphology of the ejection (figure 3). To a first approximation, the evolution of the remnant of RS Oph is ex-

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pected to be a low-energy analogue of that of a type II supernova (i.e. as described elsewhere in these proceedings, high velocity ejecta encountering a low velocity wind with a 1/r2 density distribution). Bode fc Kahn (1985) derived a simple, spherically-symmetric model (constrained by the free-free optical depth in the radio) and compared this to the observed X-ray luminosity and shock temperature. It was concluded that at the time of the first X-ray observations (55 days from outburst) the remnant was in transition between phase II (adiabatic) and phase III (isothermal) stages of development. This simple model was refined by O'Brien k Kahn (1987) who included

M.F. Bode: Novae as tracers of mass loss

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Fig. 3. EVN map at 1.7 GHz, 77 days after the 1985 outburst of R.S Oph (from Taylor et al. 1989). a first order treatment of radiative cooling to investigate its effects on remnant dynamics, and O'Brien et al. (1992) who used this model to fit the X-ray spectrum on day 55 then followed the evolution to later times. These latter authors concluded that the outburst energy was 1.1 X 1043ergs, ejected mass ~ 1O" 6 M 0 , and the mass loss rate in the pre-outburst wind ~ 2 x lO-'MQyr" 1 (for a 20 km s" 1 wind). In the context of this model, the rapid fall-off in X-rays beyond 60-70 days after outburst is attributable to the finite extent of the pre-shock circumstellar medium provided by the wind of the red giant. In the previous (1967) outburst, the wind is swept clear and only re-established over the 18 year period up to 1985. This implies an outer wind radius of ~ 1015cm. In the 1985 outburst, the shock wave reaches the edge of the wind in about 70 days and the subsequent, rapid, adiabatic expansion of the shocked gas results in the observed reduction in X-ray emission. However, although the model is generally very successful in fitting the early spectral data, exact agreement with later observations was still not possible and it is likely that the cause of this may be found in the assumption of spherical symmetry. In order to produce the bipolar structure seen in figure 3, Lloyd et al. (1993) considered (a) isotropic ejection into an anisotropic wind, and (b)

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anisotropic ejection into an isotropic wind. As detailed elsewhere in these proceedings, wind anisotropies are now recognised as being the key to understanding the shaping of planetary nebulae. However, it was found that the match with radio observations was poor. This is due to the finite extent of the wind (see above) which ensures that the ejecta cannot be collimated to sufficiently large distances, even for highly anisotropic winds. Anisotropic ejection of two high velocity blobs of material did however give a much better fit. Of particular note is the formation of a. dense equatorial ring where an oblique shock wave interaction is taking place. A natural consequence of this model is that when the ejection does not take place in the plane of the sky, optical depth effects mean that the central emission peak is displaced to lie closer to the brighter of the outer two, as observed (see Figure 4). However, if such a model is to be believed, then the question must be raised as to the physics of an outburst which would give rise to such an anisotropic ejection.

4 Classical Novae 4-1 GK Persei - a nova super-remnant As described in the introduction, classical novae are generally accepted to have unevolved mass-donating components in their central binaries. Hence, they are not expected to contain much in the way of a dense circumstellar medium prior to outburst. GK Per (Nova Persei 1901) was however detected by Reynolds & Chevalier (1984) as a non-thermal radio remnant using the VLA. More detailed studies by Seaquist et al. (1989) have shown this object again to be a low-energy analogue of a supernova remnant in which around 1% of the outburst energy has been converted into the energies of relativistic electrons and enhanced magnetic field (see Figure 5). A radio survey of 27 other CN failed to detect any similar objects (Bode et al. 1987a). The question is, why is GK Per apparently so unusual? This particular CN does however turn out to be exceptional in several other ways. For example, it has by far the longest known orbital period among the CN (~ 2 days); the cool component of the binary is evolved (K2IV) and the white dwarf is magnetised (B ~ 5xlO 5 G), and is of relatively high mass (0.72 < M-yyjy ^ 1-3M©); it undergoes dwarf nova-like outbursts of a few magnitudes, and finally it was one of only two CN around which light echoes were observed following outburst (see Seaquist et al. 1989, and references therein). As part of the investigation of GK Per and its environment, Seaquist et al. used IRAS survey data to explore the region of the light echoes. It

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M.F. Bode: Novae as tracers of mass loss Unshocked Wind

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Fig. 4. Schematic (top) showing the main features of the dynamics resulting from bipolar ejection into a finite, spherically-symmetric wind. Absorption in the unshocked wind from the front part of the ring results in the emission peak being displaced from the centre of symmetry. Below is the model radio map at l.TGHz of RS Oph, 77 days from the 1985 outburst, derived from bipolar ejection models as calculated in Lloyd et al. (1993). The radio map is constructed assuming synchrotron emission from a power law distribution of electron energies and an equipartition magnetic field.

was found that the nova lies on the saddle point joining two lobes of far infrared emission (again, see figure 5) whose colour temperature was effectively constant (T = 22 ± 1/i') and which contained a total mass of dust of 0.06 — 0.08M(TV This structure is also coincident with HI emission, where M J J J > O.6M0 . Bode et al. (1987b) suggested that the IRAS emission is in the form of a. toroid around GK Per, and could be associated with the fossil planetary nebula, of material ejected following the common envelope phase of the central binary. GK Per's 1901 outburst might then be its first, and the high velocity ejecta. were encountering material lying in cones along the axis of the toroid. Despite the large amount of circumstantial evidence that this is indeed the case, there are some problems with this interpretation, in-

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03h30min

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Fig. 5. IRAS 100//m map of the region around GK Per with superimposed a sketch of the disposition of the superlight nebulosity in 1901 and (inset) a bGHz radio map of non-thermal radio emission of the central remnant in August. 1901 (see Bode et al. 1987, and Seaquist et al. 1989 for further details). eluding the poorer correlation of CO than III emission with the IRAS nebula (Hessman 1989) and the current low-luminosity of the white dwarf (unless MWD really is very high). However, there is no doubt that GK Per is an object worthy of further detailed attention.

4.2 Recent CN Prior to the ROSATdetection of Nova Her 1991 5 days after outburst (Lloyd et al. 1992), X-ray detections of CN had only been made months after the event and had been ascribed to remnant burning on the surface of the white dwarf (Ogelman et al. 1987). In the case of Nova Her, Lloyd et al. have shown that the ROSAT PSFC data is best fitted by emission from a hot gas with kT ~ lOkeV and is inconsistent with the black body emission expected from residual burning, or pseudophotospheric emission. It is however consistent with shock temperatures that would result either from interaction of the ejecta with a pre-existing dense ambient medium, or possibly the interaction between material ejected at different velocities during outburst (see O'Brien h Lloyd 1994). Recent observations of early X-ray emission in Nova Cygni

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1992 (Krautter et al. 1993) also bear further scrutiny in this regard. This interpretation, in which X-ray emission indicates shock interactions within the ejecta, may be linked with the predictions of the common envelope phase in the outbursts of CN. Here, slower mass outflow in the orbital plane, followed by higher velocity ejection shapes the remnant of a CN in a similar manner to the shaping of planetary nebulae.

5 Future directions Our work will follow two main lines over the next few years: observations with a new generation of facilities (MERLIN II, SCUBA on the JCMT, and later the ISO satellite), and further development of our hydrodynamic models to explore the development of X-ray emission in non-spherically symmetric remnants, with particular emphasis on understanding the most recent observations of classical novae.

Acknowledgements I am very grateful to my colleagues Huw Lloyd and Tim O'Brien for commenting on a first draft of this manuscript.

References Bode, M.F., 1987, editor, RS Ophiuchi (1985) and the recurrent, nova phenomenon, VNU Sci. Press, Utrecht. Bode, M.F., & Kahn, F.D., 1985, MNRAS, 217, 205. Bode, M.F., Seaquist, E.R., k. Evans, A., 1987a, MNRAS, 228, 217. Bode, M.F., Seaquist, E.R., Frail, D.A., Roberts, J.R., Evans, A., & Albinson, J.S., 1987b, Nature, 329, 519. Bode, M.F., fc Evans, A., 1989, Classical novae, Wiley, Chichester. Dufay, J., Bloch, M., Bertaud, C , &: Diifay, M., 1964, Ann. Astrophys., 27, 555. Hessman, F.V., 1989, MNRAS, 239, 759. Hjellming, R.M., van Gorkom, J.H., Taylor, A.R., Seaquist, E.R., Padin, S., Davis, R.J., & Bode, M.F., 1986, ApJL, 305, L71. Ivison, R.J., 1992, Ph.D. Thesis, University of Central Lancashire. Ivison, R.J., Bode, M.F., Roberts, J.A., Meaburn, J., Davis, R.J., Nelson, R.F., & Spencer, R.E., 1991, MNRAS, 249, 374. Ivison, R.J., Hughes, D.H., & Bode, M.F., 1992, MNRAS, 257, 47. Joy, A.H., 1961, ApJ, 133, 493. Kenyon, S.J., 1986, The symbiotic stars, CUP, Cambridge. Kenyon, S.J., Fernandez-Castro, T., & Stencel, R.E., 1988, AJ, 95, 1817. Krautter, J., Ogelman, H., Starrfield, S., Triimper, J. fc Wichmann, R., 1993, Annals Israeli Phys. Soc, 10, 28. Lloyd, H.M., Bode, M.F., Predehl, P., Schmitt, J.H.M.M., Triimper, J., Watson, M.G., & Pounds, K.A., 1992, Nature, 356, 222. Lloyd, H.M., Bode, M.F., O'Brien, T.J., & Kahn, F.D., 1993, MNRAS, 265, 457. Lorenzetti, D., Saraceno, P., & Strafella, F., 1985, ApJ, 298, 350.

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Mason, K.O., Cordova, F.A., Bode, M.F., k Barr, P., 1987, in RS Ophiuchi (1985) and the recurrent nova phenomenon, VNU Sci. Press, Utrecht, pl67. Murset, U., Nussbaumer, H., Schmid, H.M., k Vogel, M., 1991, A&A, 248, 458. Nussbaumer, H., Schild, H., Schmidt, H.M., k Vogel, M., 1988, A&A, 198, 179. O'Brien, T.J. k Kahn, F.D., 1987, MNRAS, 228, 277. O'Brien, T.J., k Lloyd, H.M., 1994, in Kinematics and dynamics of diffuse astrophysical media, ed. J.E. Dyson, APSS, (in press). O'Brien, T.J., Bode, M.F., k Kahn, F.D., 1992, MNRAS, 255, 683. Ogelman, H., Beuermann, K., k Krautter, J., 1987, A&A, 177, 110. Padin, S., Davis, R.J., k Bode, M.F., 1985, Nature, 315, 306. Purton, C.R., Feldman, P.A., Marsh, K.A., Allen, D.A., k Wright, A.E., 1982, MNRAS, 198, 321. Reynolds, S.P., k Chevalier, R.A., 1984, ApJL, 281, L33. Seaquist, E.R., Taylor, A.R., k Button, S, 1984, ApJ, 284, 202. Seaquist, E.R., Bode, M.F., Frail, D.A., Roberts, J.A., Evans, A., k Albinson, J.S., 1989, ApJ, 344, 805. Seaquist, E.R., Krogulec, M., k Taylor, A.R., 1993, ApJ, 410, 260. Taylor, A.R., Davis, R.J., Porcas, R.W., k Bode, M.F, 1989, MNRAS, 237, 81. Webbink, R.F., Livio, M, k Truran, J., 1987, ApJ, 314, 653. Webster, B.L, k Allen, D.A., 1975, MNRAS, 171, 171. Whitelock, P.A, k Munari, U., 1992, A&A, 255, 171.

Circumstellar scattering processes in symbiotic systems H.M. Schmid and H. Schild Institute of Astronomy, ETH Zentrurn, CH-8092 Zurich, Switzerland

1 Circumstellar matter in symbiotic stars By definition, symbiotic stars exhibit simultaneously the absorption features of a cool giant and high excitation emission lines, e.g. Hell, [OIII], of an ionised nebula. Thus the presence of circumstellar material is a necessary classification criterion of these objects. It is now generally accepted that symbiotic stars are binary systems consisting of a red giant and a hot radiation source, in most cases a hot white dwarf. The strong emission lines originate in a dense nebula which is thought to be wind material lost by the cool giant and ionised by the hot companion. Besides the compact nebula other components of circumstellar material are observed. Enhanced IR emission due to circumstellar dust is found in D-type (D for dust) symbiotic systems. Very extended ionised regions have been mapped with radio interferometers or optical imaging techniques (e.g. Taylor 1988, Solf 1988, Corradi & Schwarz 1993, Schwarz 1993). Some features of the extended structures can be associated with bipolar outflow (velocities ~ 100 km/s) from the central, unresolved binary system. In this paper we discuss how the geometric structure of the circumstellar environment of symbiotic systems can be clarified from an analysis of light scattering processes. 2 Scattering processes in the circumstellar environment of symbiotic stars Polarisation measurements are a well known tool for studying scattering processes. Polarimetric observations of symbiotic stars in broad and narrow band filters have shown that these objects are often intrinsically polarised (e.g. Piirola 1983, Schulte-Ladbeck 1985, Schulte-Ladbeck et al. 1990). The obtained polarisation pattern is often very complex and variable in time suggesting that more than one polarising mechanism is in operation. Therefore the polarimetric data are difficult to interpret because from measurements in a few filters only a very limited amount of the polarisation information can be extracted. But the polarisation studies showed that the polarisation is due to scattering on asymmetrically distributed (or illuminated) circum331

332

H.M. Schmid & H. Schild: Light scattering in symbiotic stars

stellar material and that dust grains are a dominant agent for producing polarisation. Rayleigh scattering by atomic hydrogen is observed in eclipsing symbiotic systems (Isliker et al. 1989). When the hot component approaches occultation the UV radiation below 1400A is strongly attenuated due to Rayleigh scattering in the extended atmosphere of the cool giant. The Rayleigh cross section depends strongly on wavelength, and increases sharply towards the Lycn A1215 resonance transition, with a corresponding heavy effect on the continuum shape. From detailed analysis of the observed spectral changes, the density structure of the stellar wind from the cool giant and the extent of the neutral gas in the innermost region of symbiotic binaries can be derived (Gonzalez-Riestra et al. 1990, Vogel 1991). Raman scattered emission lines at A6825 and A7082 are special features which are only observed in symbiotic systems. These lines are produced by a Raman scattering process of the OVI AA1032,1038 resonance lines by neutral hydrogen (Schmid 1989). In this process the incident OVI photon with frequency i>,- excites atomic hydrogen from its ground state ls 2 S to an intermediate virtual state from where the Raman scattered photon vj is emitted, leaving hydrogen in the excited state 2s 2 S. The frequency of the scattered photon is given by i/j = U{ — f;/, where Vij is the frequency corresponding to the energy difference between initial state (ls 2 S) and final state (2s 2S) of hydrogen. In symbiotic systems strong OVI line radiation is produced near the hot component and Raman scattered in the neutral region near the red giant star. The asymmetric geometry is expected to produce a net polarisation in the Raman scattered lines, which is indeed detected (Schmid & Schild 1990). 3 Polarisation in the Raman scattered lines In order to investigate the scattering geometry for the Raman lines we made a spectropolarimetric survey of 15 symbiotic stars (Schild & Schmid 1992, Schmid & Schild 1994). This survey shows that the Raman scattered OVI lines AA6825,7082 are always, and sometimes very strongly polarised. Typically the polarisation integrated over the Raman lines is around 5%. In addition we find structural variations in the polarisation degree and angle across the line profile. Maximum polarisation in parts of the line can be as high as 15%. The polarisation structure is always very similar in the A6825 and A7082 components (see Fig. 1). We can distinguish between three types of profiles with increasing com-

H.M. Schmid & H. Schild: Light scattering in symbiotic stars

6800

6820

6840

7070

7090

333

7110

Fig. 1. Type III polarisation profile of the Raman lines A6825 (left) and A7082 (right) in RRTel. The total flux (top panel), percentage polarisation (center panel) and polarisation angle (bottom panel) is shown. plexity: (I) Profiles with constant polarisation through the line. (II) Profiles with decreasing polarisation towards the red line wing and constant polarisation angle. (Ill) Profiles with two components and a rotating or flipping polarisation angle. There are about equal numbers of symbiotic stars in each class. An interesting property is that type III profiles are only seen in systems containing a Mira variable or a late M giant (>M5). A simple model can explain the observed pattern in the polarisation profiles (see Schmid & Schild 1994 for details). Consider first OVI photons from the hot component traveling along the binary axis (~ parallel) towards the cool giant. They will be Raman scattered between the two stars by approaching H°-atoms in the red giant's wind. This produces a blue shifted Raman photon and a polarisation direction perpendicular to the binary axis. OVI photons with directions roughly perpendicular to the binary axis will be scattered by receding atoms in the outer wind of the red giant. Thereby, they get a. red shift and a polarisation parallel to the binary axis. Thus a

334

H.M. Schmid & H. Schild: Light scattering in symbiotic stars

type III polarisation profile, with a distinct red component and an angle flip, is only expected if the scattering probability, or the H°-density, in the outer regions of the cool giant's wind is high. This explains why type III profiles are only observed in systems with very late type giants undergoing heavy mass loss. Another result of this simple model is that the polarisation direction in the blue line wing has to be perpendicular to the binary axis. 4 Outlook The polarisation found in the A6825 and A7082 features now opens a new avenue of investigation. Repeated observations will show the periodic variations of the polarisation vector due to the binary motion. This will permit to establish accurate orbital parameters such as period, inclination and orientation of the orbital plane. This geometric information is also interesting in relation to other observations of the structure of symbiotic nebulae from radio interferometers or optical imaging techniques. In particular the polarisation direction of the Raman lines should clarify, whether mass outflow from symbiotic binaries occurs in polar, meridional or equatorial directions. We expect that with elaborate theoretical calculations much more information about symbiotic systems can be extracted (see Schmid 1992). The Raman lines contain velocity and polarisation information which can be obtained for various orbital phases. Time series of polarisation measurements will thus permit to map the density and velocity structure of the scattering H°-atoms in symbiotic systems. This is particularly interesting for understanding the interaction processes taking place in binary stars. References Corradi R.L.M., Schwarz H.E., 1993, Astr. Astrophys., 268, 714 Gonzalez-Riestra R., Cassatella A., & Fernandez-Castro T., 1990, Astr. Astrophys., 237, 385 Isliker H., Nussbaumer H., k. Vogel M., 1989, Astr. Astrophys., 219, 271 Piirola V., 1983, in: Cataclysimic Variables and Related Objects, M. Livio & G. Shaviv (eds.), IAU Coll. 72, 2H Schild H., & Schmid H.M., 1992, Gemini - Newsletter RGO, 37, 4 Schmid H.M., 1989, Astr. Astrophys., 211, L31 Schmid H.M., 1992, Astr. Astrophys., 254, 224 Schmid H.M., & Schild H., 1990, Astr. Astrophys., 236, L13 Schmid H.M., & Schild H., 1994, Astr. Astrophys., 281, 145 Schulte-Ladbeck R.E., 1985, Astr. Astrophys., 142, 333 Schulte-Ladbeck R.E., Aspin C , Magalhaes A.M., &; Schwarz H.E., 1990, Astr. Astrophys. Suppl, 86, 227 Schwarz H.E., 1993, this volume Solf J., 1988, in: The Symbiotic Phenomenon, J. Mikolajewska et al. (eds.), IAU Coll. 103, 85 Taylor A.R., 1988, in: The Symbiotic Phenomenon, J. Mikolajewska et al. (eds.), IAU Coll. 103, 77 Vogel M., 1991, Astr. Astrophys., 249, 173

Poster Papers Presented

Stellar Evolution and Wind Theory From Red Giants to White Dwarfs: what can be said about mass loss? H. E. Kley & C. A. Tout Radiation pressure on dust as a mechanism of acceleration of circumstellar envelopes N. Netzer Planets, binary stars and axisymmetric mass-loss N. Soker

Wolf-Rayet stars and Luminous Blue Variables Direct imaging and velocity-resolved long-slit spectra of the nebula around P-Cygni M. J. Barlow, J. E. Drew, J. Meaburn & R. Massey A theoretical study of the kinematic properties of VVR. ring nebulae A. Burkert & N. Langer Coronographic imaging of the P-Cygni nebula M. Clampin, M. Robberto, A. Nota, F. Paresce & J. Staude Thermal and non-thermal radio emission from VVR. 147 R. J. Davis & P. E. Pavelin The radio variability of P Cygni K. Exter, M. Barlow, M. Bode, R. Davis, I. Uoxtmrt.li & C Skinner Nebulae around LBVs D. Hutsemekers & E. van Drom Wind collision in WR+O binaries: effects of energy losses and thermal conductivity A. V. Myasnikov & S. A. Zhekov The nature of extended shells around Wolf-Rayet stars with ring nebulae J. Nichols-Bohlin & R. A. Fesen 335

336

Poster Papers

Consistent hydrodynamic atmosphere models for hot luminous stars: photospheric and wind diagnostics D. Schaerer & W. Schmutz Wind asymmetries in AG Car detected with spectropolarimetry T. J. Harries, R. E. Schulte-Ladbeck, G. C. Clayton, D. J. Hillier & I. D. Howarth Wind inhomogeneities in WR40, the central star of the WR ring nebula RCW58 R. E. Schulte-Ladbeck, D. J. Hillier & G. C. Clayton Properties and variability of the stellar wind from P Cygni 5. Scuderi, G. Bonanno, D. Spadaro, N. Panagia, H. J. G. L. M. Lamers & A. de Koter Colliding stellar winds: near IR spectroscopy and new hydro calculations /. R. Stevens, A. M. T. Pollock & I. D. Howmih Spectroscopy of LBV nebulae in local group spirals J. V. Vilchez Dust condensation around AG Carinae and its evolutionary stage

R. Viotti, C. Rossi, V. F. Polcaro, A. Cassatella & M. Barylak Episodic condensation of dust in WR stellar winds P. M. Williams, K. A. van der Hucht, M. R. Kidger, P. Bouchet & P. A. Whitelock

Supernovae Circumstellar interactions: Kepler's supernova remnant K. J. Borkowski, J. M. Blondin & C. L. Sarazin The progenitor star of SN 1987A and its circumstellar nebula A. P. S. Crotts, W. E. Kunkel, S. R. Heathcote & E. Bonar Circumstellar emission from Supernova 1992ad R. J. Cumming & W. P. S. Meikle High-resolution spectroscopy of the SN 1987A circumstellar medium R. J. Cumming, W. P. S. Meikle, J. Spyromilio. D. A. Allen & P. Lundqvist

Poster Papers

337

Planetary Nebulae Detection of C I emission at 609/im from the Ring Nebula in Lyra R. Bachiller, P. J. Huggins, P. Cox & T. Forveille HST-WFPC observations of the planetary nebula K648 in M15 L. Bianchi et al. Structure and kinematics of planetary nebulae and their faint giant haloes M. Bryce Morphological segregation of planetary nebulae: which progenitors? R. L. M. Corradi & H. E. Schwarz High spatial resolution HCO + observations of the PN NGC 7027 P. Cox, S. Guilloteau, R. Bachiller, P. J. Huggins, A. Omont & T. Forveille Radiation gas dynamics of planetary nebulae: from diversity to unity A. Frank Binary evolution and the formation of bipolar planetary nebula Zhanwen Han, Ph. Podsiadlowski & P. Eggle.ton The OH/IR - PN connection: kinematics and masses W. J. Maciel & R. Ortiz The formation of surrounding shells around PNe G. Mellema Correlations between PN morphology and central star evolution L. Stanghellini, R. Corradi & H. Schwarz The structure of the knots in the Helix nebula (NGC 7293) J. R. Walsh & J. Meaburn X-ray emission from wind-blown bubbles: ROSAT observations of NGC 6888 M. Wrigge & H. J. Wendker The planetary nebula IRAS 08355-4027 R. D. Wolstencroft, S. M. Scarrott, M. A. Read & Q. A. Parker X-ray emission from planetary nebulae during their evolution 5. Zhekov & M. Perinotto

338

Poster Papers

Asymptotic Giant Branch Stars and their Envelopes Molecular absorption in the circumstellar environment of the Post-AGB Star HD 56126 E. J. Bakker The evolution from visible to infrared variable AGB stars J. A. D. L. Blommaert, A. G. A. Brown, H. J. Habing, Y. K. Ng & W. E. C. J. van tier Veen The discrimination of 0- and C-rich circumstellar envelopes from molecular observations V. Bujarrabal, A. Fuente, A. Omont & J. Alcoleu Induced nucleation of carbon dust in red giant stars B. J. Cadwell, Hai Wang, E. D. Feigelson & M. Frenklach Hydrogen lines and recombination mechanisms in Long-Period Variables P. de Laverny & C. A. Magnan Detection of far infrared line emission from Alpha Orionis A. E. Glassgold & M. R. Haas The dust shells around infrared carbon stars M. A. T. Groenewegen New detections of neutral K and Na scattering around red giants C. Guilain & N. Mauron Modelling the continuum emission from the envelopes of post-AGB stars J. Giirtler, Th. Henning & C. Kompe A study of IRAS oxygen-rich AGB stars 0. Hashimoto Impact of shock-waves on SiO masers around Long-Period Variables A. Heske, J. Hron & H. Olofsson Physics of shock-waves in envelopes of evolved stars E. Huguet & J-P. Lafon Dust and gas mass-loss rates from AGB Stars K. Justtanont, C. J. Skinner & A. G. G. M. Tielens The onset of the superwind in OH/IR stars B. M. Lewis

Poster Papers

339

Radiative transfer in axisymmetric circumstellar shells J. Lefevre, B. Lopez & D. Mekarnia A new type of OH Mira emission A. M. Le Squeren, E. Gerard & S. Etoka SiC in circumstellar shells around C stars J. Lorenz-Martins & J. Lefevre Infrared spectra of post-AGB Stars A. Manchado, P. Garcia-Lario, K. Sahu & S. R. Pottasch The 150 AU structure of the radio continuum and the ammonia bipolar outflow in CRL 618 J. Martin-Pintado, R. Gaume, R. Bachiller & K. Johnston Carbon stars with detached CO envelopes H. Olofsson, P. Bergman, U. Carlstrom, K. Eriksson & B. Gustafsson Near-infrared hydrogen emission lines in IRC+10-120 R. Oudmaijer, T. Geballe, R. Waters & K. Sahu Detection of Herbig-Haro condensations in the bipolar outflow of M 1-92 J. Solf High-resolution imaging of evolved stars P. G. Tuthill, J. E. Baldwin, C. A. Haniff & R. W. Wilson The radial distribution of dust around evolved stars W. E. J. C. Van der Veen The circumstellar environment of R CrB stars H. J. Walker Images of IRC+10216 in the lines of cyanopolyyene molecules H. A. Wootten, R. Sahai, Nguyen-Q-Rieu & Truong-Bach 13

CO interferometric observations of CRL 618, the 'Egg nebula' /. Yamamura, K. Shibata & S. Deguchi

Kinematic modelling of H2O maser shells J. A. Yates

340

Poster Papers

Novae and Symbiotic Stars Wind mass loss during nova outburst and theoretical light curves M. Kato & I. Hachisu Evidence signalling the start of counterjet activity in R Aquarii A. G. Michalitsianos, M. Perez & M. Kafatos Merlin radio observations of Nova Cygni 1992 P. E. Pavelin, R. J. Davis, L. V. Morrison, M. F. Bode & R. J. Ivison IUE observations of the symbiotic binary CH Cygni during its current outburst A. Skopal & M. F. Bode Resolved nebulosities around symbiotic stars U. Munari & F. Patat Periodic dust obscurations in symbiotic Miras U. Munari & B. Yudin

Miscellaneous A modified radiative wind model for B[e] stars F. X. de Araujo & D. Petrini Evolved metal rich low mass stars as contributors of the UV emission in elliptical galaxies F. Fagotto, C. Chiosi, A. Bresson & G. Bertelli

Author index

Bachiller, R., 304 Balick, B., 291 Blondin, J. M., 139 Bode, M. F., 321 Borkowski, K. J., 300 Centurion, M., 227 Cherchneff, I., 232 Chevalier, R, 308 Chu, Y-H. 73 Chugai, N., 148 Clampin, M., 89 Clegg, R. E. S., 300 Cox, P., 304 Cumming, R. J., 192 Davidson, K., 95 Dopita, M. A., 73 Drissen, L., 89 Drew, J. E., 27 Dufour, R. E., 78 Dyson, J. E., 52 Ebbets, D., 95 Fleischer, A. J., 262 Forveille, T., 304 Fransson, C , 120, 198 Garcia-Segura, G., 85 Garner, H., 95 Gauger, A., 262 Geballe, T. R.,, 296 Green, D. A., 203 Guelin, M., 266 Hajian, A., 291. Harrington, J. P., 300. Hartquist, T. W., 52 Hrivnak, B., 296 Hsu, J. J. L., 179, 187 Huggins, P. J., 304 Icke, V., 309

Joss, P., 179, 187 King, D. L., 227 Kwok, S., 296 Langer, N., 1 Leibundgut, B., 100 Leitherer, C , 89 Lipman, K., 227 Livio, M., 35 Lozinskaya, T. A., 73 Lucas, R., 266 Lundqvist, P., 192, 213 Maccioni. A., 291 Mac Low, M-M 85 Malumuth, E., 95 Meikle, W. P. S., 192 Neri, R., 266 Nomoto, K., 221 Nota, A., 89 Olofsson, H., 246 Panagia, N., 112, 207 Paresce, F., 89 Pascual, P. R., 198 Pasquali, A., 89 Perinotto, M., 291 Podsiadlowski, Ph., 179, 187 Pooley, G. G., 203 Robberto, M., 89 Robert, C , 89 Ross, R. R., 179, 187 Rupen, M. P., 207 Schild. H., 331 Schmid, H. M., 331 Schwarz, H. E., 274 Sedlmayr, E., 262 Shigeyama, T., 221

341

342

Author index

Smith, L.J., 64 Sonneborn, G., 198 Sramke, R. A., 112, 207 Suzuki, T., 221 Tenorio-Tagle, G., 166 Terlevich, R. J., 153 Terzian, Y., 291 Tielens, A. G. G. M., 232 Tsvetsanov, Z., 300 Unger, S., 227 Van Dyk, S., 112, 207 Vladilo, G., 227 Walborn, N., 95 Walton, N., 192, 227 Wamsteker, W., 198 Weiler, K., 112, 207 White, R., 95 Winters, J. M., 262 Wood, P., 15

Object index

30 Doradus, 75 41.9+58, 160, 163 Abell 30, 300-301, 303 Abell 35, 48 Abell 41, 48 Abell 46, 48 Abell 58, 300 Abell 63, 48 Abell 65, 48 Abell 78, 300-303 AG Carinae, 65-71, 81, 93-94 AG Pegasi, 290 AS201, 290 B0954 + 658, 204 BD +30°3639, 251, 297 BD -22°3467, 48 BE 294, 90-91 BE 381, 90-91 BE UMa, 48 Beta Canis Majoris, 33 BI Crucis, 35, 286, 290 Calabash Nebula, 311 Cassiopeiae A, 135-136, 159, 162-163, 172, 177 CH Cygni, 286, 290 Chi Cygni, 249 CPD - 5 6 8032, 297 CRL 2688, 250, 256-257, 272 CRL 618, 250, 256, 270-272 CW Leonis, 266 DS 1, 48 Dumbbell, 305 EGB 5, 49 Epsilon Canis Majoris, 33 Eta Carinae, 65-70, 81, 95-98, 317 Gamma Velorum, 10 GK Persei, 326-328 Hl-36, 290

HD 101584,251, 257 HD 112313,48 HD 160529, 66-67 HD 192163, 79 He 2-104, 35, 286, 290 He 2-111, 281, 283-284 He 2-113, 297 He 2-114, 305 He 2-147, 290 He 2-36, 35 He 2-442, 290 He 2-77, 304 He 3-519, 65-67, 71 Helix nebula, 258, 311 Hen 1357, 35 Hen 1383, 290 HFG 1, 48 HM Sagittae, 290 Homunculus, 95-97, 99, 317 HR Carinae, 66, 69 Ht Tr 4, 48 IC 2149, 291, 294 IC 4182, 207 IC 4406, 35, 252, 257 IC 4593, 316 IC 4634, 35, 46 IN Comae Berenices, 48 IRAS 04296+3429, 297-298 IRAS 05113+1347,298 IRAS 05341+0852, 297-298 IRAS 07134+1005, 297-298 IRAS 15154-5258,300 IRAS 16342-3814, 257 IRAS 17253-2824, 24 IRAS 17514-1555,297 IRAS 18333-2357, 277, 280, 300 IRAS 18520+0533, 24 IRAS 19500-1709, 250 IRAS 20000+3239, 298 IRAS 21282+5050, 251, 257, 297-298 IRAS 22223+4327, 250, 298 IRAS 22272+5435, 297-298

343

344

Object index

IRAS 2257+6609, 298 IRAS 23304+6147, 250, 298 IRC+10011, 243-244 IRC+10216, 244, 252-253, 255-256, 266-272 IRS16, 60 IRS7, 52-53, 60-61 J 320, 35 K 3-72, 35 Kl-2, 35, 37, 48 K4-45, 304 Kepler SNR, 172 KV Velorum, 48 L1641.44 LH 100, 75 LMC N201, 277 LMC N66, 277, 279 LoTr 1, 48 Lo Tr 5, 48 LSS 2018, 48 LW Hya, 48 Ml-67, 66, 70, 81 Ml-13, 305 Ml-16, 280-283, 285 Ml-17, 251 Ml-78, 304 Ml-92, 257 Ml 5, 277, 280 M2-51, 258 M2-9, 281 M22, 277, 280, 300 M4-18, 297 M4-9, 257 Mrk 297A, 113, 159-160 MT Serpentis, 48 My Cn 18, 35 Mz 1, 305 NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC NGC

1566, 168-169, 176 2346, 35, 37, 48, 252, 257-258 2359, 66, 78, 81-84 2976, 230 3031, 198, 203, 207 3077, 230 3132, 252, 257 3242, 291, 294 5455, 103 5457, 103 5548, 171 6072, 257 6210, 316 6302, 284, 286 6396, 314 6537, 81 6543, 57, 291, 293-294 6563, 257 6720, 257, 307 6772, 257-258

NGC 6781, 252, 257-258 NGC 6826, 291, 293-294 NGC 6829, 280 NGC 6853, 305 NGC 6888, 9, 78-81, 84, 88 NGC 7009, 35, 291, 293-294 NGC 7027, 251-252, 257, 272, 280-282, 297-298, 306-307 NGC 7293, 52-53, 59, 61. 258 NGC 7662, 291, 294 Nova Cygni 1992, 329 Nova Herculis 1991, 328 OH 17.7-2.0,35 OH 231.8+4.2, 250, 256 OH 26.5, 242 Owl Nebula, 311 P Cygni, 10, 65-69 PHL 932, 49 Pi1 Giuis, 249 PS1, 277, 280 Piippis A, 172 R Aquarii, 35. 290 R Sculptoris. 254 R127, 65, 70, 93-94 R71, 65, 93 RCW58, 52, 54-57, 61-62, 66, 68, 70 Red Rectangle, 280, 282 Ring Nebula, 257, 307 RR Telescopii. 333 RS Ophiuchi, 290, 322-325, 327 RX Puppis, 290 R Cassiopeiae. 249 S 119,89-93 S 308, 66, 81 S Doradus, 65, 93 S Scuti, 254 ScWe 2. 35 ScWe 3, 35 Sh2-71, 48 Sk -69°202, 139, 184-185 SMC N2, 277 SMC N5, 277, 279 SN 1937C, 203, 207 SN 1950B, 103 SN 1957D, 101-103, 106, 109-110, 113, 133 SN 1961V, 101-102 SN 1969L, 148. 150 SN 1970G, 101-103, 106, 109, 112, 116-117, 148, 169 SN 1978K, 101-102, 113, 133, 136, 150-151, 160 SN 1979C, 100-104, 109-110, 112-114, 116-118, 120, 130-131, 133, 148, 153, 160 SN 1980K, 100-106, 109, 113, 116-118, 130-131, 133, 148, 160 SN 1981K, 101-102, 113, 116-117 SN 1982E, 131 SN 1983N, 101, 103, 113, 116-117 SN 1984E, 106, 131

Object index SN 1984L, 101, 113, 116-117 SN 1986J, 101-102, 106-109, 113, 115-119, 133, 135-136, 150-151, 153, 159-160, 162-163 SN 1987A, 2, 4, 7, 12, 100-102, 107-108, 113, 121, 126, 131, 139-143, 145-146, 148-149, 172, 179-180, 184-185, 187, 190, 197, 203, 213-214, 221, 226-227, 308, 317 SN 1987F, 148-149, 153, 160, 163 SN 19881, 160 SN 1988Z, 101-102, 107-109, 113-114, 116-118, 133, 135, 149-151, 153, 160-163 SN 1989R, 101, 107-109 SN 1990B, 101, 113-114, 116-117 SN 1993J, 7, 100-102, 113, 116, 119-120, 127-130, 136, 179-180, 185, 187-188, 192, 196-201, 203-205, 207-211, 213, 216-218, 224, 226-230 SN~ 1955, 159 Sp 1, 48 SS96, 290 Tau Scorpii, 30 TT Cygni, 254 TychoSNR, 172 U Antilae, 254 U Orionis, 253, 256 UU Sagittae, 48 V Hydrae, 49, 249 V1016 Cygni, 290, 323 V417 Centauri, 286, 290 V477 Lyrae, 48 V605 Aquilae, 300 V651 Monocerotis, 48 V664 Carinae, 48 Vo 1, 297 VV47, 257-258 VW Pyxidis, 48 W Hydrae, 252 W43A, 257 WRA 751, 66, 69

345

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