TEN YEARS OF GEOTAIL AND ITS CURRENT STATUS: A BRIEF SUMMARY T. Mukai Institute of Space and Astronautical Science/JAXA, Yoshinodai, Sagamihara 229-8510, Japan
ABSTRACT Since its launch in 1992, GEOTAIL has extensively surveyed the magnetotail with a full set of plasma and field instruments over a wide range of distances from 9 Re to 220 Re away from the Earth by means of a sophisticated orbit strategy. In the first two years, the orbit was optimized to explore the distant tail, and thereafter was changed to study substorm processes in the near-Earth tail region. The near-tail orbit has also facilitated exploration of the dayside outer magnetosphere, the magnetopause, the magnetosheath, and the bow shock, as well as the upstream solar wind. GEOTAIL observations have revealed a number of new phenomena in these regions, and as of the end of 2003, about 600 papers have been published in refereed journals. Recent GEOTAIL studies have significantly advanced our understanding of the structure and formation of thin current sheets in the mid-tail plasma sheet during substorms, and have elucidated new kinetic aspects of magnetic reconnection. GEOTAIL has operated far beyond the designed mission life of three and half years. Most of the onboard instruments are still functioning well, and it is expected that GEOTAIL will continue to generate scientifically useful data.
GEOTAIL MISSION The GEOTAIL spacecraft was launched on 24 July 1992 by a Delta-II launch vehicle from Cape Canaveral, Florida, U. S. A. This is a joint program of the Institute of Space and Astronautical Science (ISAS) of Japan and the National Aeronautics and Space Administration (NASA) of U. S. A. ISAS developed the spacecraft and provided about two thirds of the science instruments, while NASA provided the launch and about one third of the science instruments. The spacecraft is operated from ISAS, but the data are acquired by both agencies. The prime objective is to investigate the structure and dynamics of the Earth's magnetotail with a comprehensive set of scientific instruments [Nishida, 1994]. As shown in Figure 1, in the first two years, the apogee was maintained on the nightside of the Earth by means of the lunar double swingby maneuvers, ranging from 80 to 220 Re in order to explore the distant tail. Later, from November 1994 the apogee was reduced first to 50 Re and then to 30 Re in order to study substorm processes in the near-Earth tail region. The perigee was set at 10 Re, and in June 1997, slightly reduced to about 9 Re. The inclination was set at -7° with respect to the ecliptic plane, so that the spacecraft has most frequently traversed the neutral sheet at the apogee near the December solstice. This orbit strategy has worked highly satisfactorily for the mission objectives, and also enabled exploration of the dayside outer magnetosphere, the magnetopause, the magnetosheath, the bow shock, and the upstream solar wind.
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Figure 2 shows a schematic view of the GEOTAIL spacecraft. The spacecraft weighed 1008kg initially, but now it is about 660 kg, since most of hydrazine propellant has been consumed for the orbit maneuvers. It has a cylindrical shape with diameter of 2.2 m and height of 1.6 m. Two masts which are 6-m long are deployed symmetrically to separate the magnetometers from the main body, and four 50-m antennas are deployed to measure the electric field from DC to 800 kHz. The spacecraft is spin-stabilized with the spin axis being nearly perpendicular to the solar ecliptic plane. The spin rate is 20 rpm.
Figure 2 A schematic view of the GEOTAIL spacecraft on orbit.
Figure 1 GEOTAIL orbit in the Geocentric Solar Magnetospheric (GSM) coordinates. The upper and lower panels show the distant-tail orbit (from September 1, 1992 to November 10, 1994) and the near-tail orbit (from November 10, 1994 to February 5, 1996), respectively. The present one is nearly the same as that shown in the lower panel.
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BRIEF SUMMARY OF RESULTS FROM GEOTAIL The initial GEOTAIL results were summarized in a special section of Geophysical Research Letters [December, 1994], and some highlights were compiled in Journal of Geomagnetism and Geoelectricity [Vol. 48, Nos. 5,6, 1996], a special section of Journal of Geophysical Research [March, 1998] and the AGU Geophysical Monograph 105 "New Perspectives on the Earth's Magnetotail" [1998]. Nishida [2000] reviewed the dynamic magnetotail based on the GEOTAIL observations. As of the end of 2003, about 600 papers have been published in refereed journals, as shown in Figure 3.
Figure 3
Number of GEOTAIL papers
Several new discoveries have resulted from the GEOTAIL observations, not only due to the orbit strategy optimized to the first extensive survey of the magnetotail over a wide range of distances, but also as a result of sophisticated measurement techniques. For example, the discovery of electrostatic solitary waves (ESW) in the plasma sheet boundary layer was due to the waveform capture (WFC) technique of the Plasma Wave Instrument [Matsumoto et al., 1994]. The WFC method has produced a number of new findings in wave properties in the magnetotail, the dayside magnetosphere, the magnetosheath, the bow shock and the upstream solar wind. Another important contribution to the success of GEOTAIL was the high time resolution measurement of 3-D distribution functions of ions and electrons with the Low Energy Particle (LEP) experiment [Mukai et al., 1994]. In particular, the LEP ion measurements with high sensitivity (large geometrical factors) have been quite effective for measurements of tenuous plasmas in the magnetotail; for example, the discovery of cold oxygen ion beams in the distant lobe/mantle regions [e.g., Seki et al., 1998]. The LEP ion measurements have also provided 24-hour continuous data of the ion fluid parameters (velocity moments; density, velocity and temperature), which have
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been very useful for statistical analysis of plasma properties. Combined with the magnetic field data, these plasma properties have been used to investigate the magnetotail structure and dynamics, which are well expressed in terms of the ideal MHD. Magnetic reconnection is one of the most important processes in space plasmas. Its in-depth understanding was the main target of GEOTAIL in terms of space plasma physics, and in the following we briefly summarize some important results on magnetic reconnection. Firstly, GEOTAIL has established that that the structure and dynamics of the magnetotail are basically determined by magnetic reconnection under both southward and northward IMF conditions, except possibly when IMF is almost due northward [Nishida et al., 1998]. A further important finding is that the neutral sheet in the distant tail is twisted under the influence of the By component of IMF, the twisting being more severe when the IMF is northward. GEOTAIL observations have clearly demonstrated the importance of magnetic reconnection for the magnetospheric substorms. Statistical results of plasma flow properties suggest that near-Earth reconnection starts before the expansion phase starts on the ground, and that this neutral line is initially formed at the distance of 22-30 Re in the midnight-premidnight region, coinciding with the local time range where the auroral and geomagnetic signatures of the expansion phase onset take place [Nagai, et al., 1998]. Studies on evolution of plasmoids have revealed that plasmoids, after initial formation in association with the near-Earth neutral line, expand toward the flank sides (dawnwad and duskward) during tailward propagation until the mid-tail region (-70 Re), and they have a full width of the magnetotail beyond this distance down the tail [Ieda et al., 1998]. The plasmoids are accelerated tailward until -100 Re, whereas the tailward speed is reduced beyond this distance, probably due to an interaction with pre-existing plasmas. In addition to the above MHD features, GEOTAIL observations of electron and ion velocity distribution functions have produced a new kinetic understanding of plasmoids and flux ropes [Mukai et al., 1998]. Characteristic signatures of the ion distribution functions are identified during the passage of plasmoids, and by comparison with computer simulations, they are qualitatively and quantitatively understood in terms of temporal/spatial evolution of the particle acceleration and heating by the magnetic reconnection [Hoshino et al., 1998]. Suprathermal electron acceleration in magnetic reconnection is also revealed by comparison between GEOTAIL observations and computer simulations [Hoshino et al., 2001]. Recent studies have significantly advanced our understanding of the evolution of thin current sheets and the associated occurrence of magnetic reconnection during the course of substorms [e.g., Mukai et al., 2000; Asano et al., 2003]. For these studies, the availability of reliable electron data has been crucial, and the structure of an extremely thin current sheet has been revealed with the current density estimated from the electron and ion velocity moments. One of the most important GEOTAIL discoveries is that the cross-tail current sheet in the late growth phase and the early expansion phase temporally forms a bifurcated structure in which the current density becomes the largest away from the neutral sheet [Asano et al., 2003]. The intense current sheet is thinner closer to the X-line, and the thickness becomes less than the ion inertial length. In the thin current sheet, the current is mainly carried by electrons, contrary to the ion-dominated diamagnetic current with the higher ion temperature than the electron one. The Hall electric field toward the neutral sheet causes the electron-dominated current sheet by enhancing the dawnward drift for both ions and electrons. GEOTAIL observations have also revealed the structure of the Hall current system and characteristic features of the associated electron distribution functions carrying the field-aligned current for magnetic reconnection [e.g., Nagai et al., 2003; Asano et al., 2004]. CURRENT STATUS Two data bases, namely the Key Parameters and the Science Data Base, have been produced from experiments onboard GEOTAIL. The Key Parameters are produced by the NASA Goddard Space Flight Center from the 24-hour continuous data dumped over the NASA/DSN, and are made available to the worldwide science community for event search and preliminary studies via the NASA/CDAWeb system. The calibrated plasma and
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magnetic field data with high time resolution are archived in the Science Data Base, and are also made available to the worldwide science community via DARTS (Data ARchive and Transmission System at the Center for Planning and Information Systems, ISAS/JAXA) at
. The plasma wave data and energetic particle data are also available via web-pages at the home institutions of the principal investigators,
and
, respectively. GEOTAIL has operated beyond the designed mission life of three and half years. The occurrence of eclipses with duration longer than the designed length (2 hours) is unavoidable, but the spacecraft could be operated without any serious problem for an eclipse lasting for 269 minutes in February 2000, the longest period in the past. Most of the onboard instruments are still functioning well, and it is expected that GEOTAIL will continue to generate scientifically useful data. ACKNOWLEDGMENTS The author thanks Professor Emeritus A. Nishida and Professor T. K. Uesugi of ISAS/JAXA, Dr. Mario Acuna of NASA Goddard Space Flight Center, and all Pi's and their teams for the success of GEOTAIL and production of excellent quality data. REFERENCES Asano, Y., T. Mukai, M. Hoshino, Y. Saito, H. Hayakawa, and T. Nagai; Evolution of the thin current sheet in a substorm observed by Geotail, J. Geophys. Res., 108, 1189, doi:l0.1029/2002JA009785, 2003. Asano, Y., T. Mukai, M. Hoshino, Y. Saito, H. Hayakawa, and T. Nagai; Current sheet structure around the nearEarth neutral line observed by Geotail, J. Geophys. Res., 109, A02212, doi:1029/2003JA010114, 2004. Hoshino, M., T. Mukai, T. Yamamoto, and S. Kokubun; Ion dynamics in magnetic reconnection: Comaprison between numerical simulation and Geotail observation, J. Geophys. Res., 103, 4509-4530, 1998. Hoshino, M., T. Mukai, T. Terasawa, and I. Shinohara; Suprathermal electron acceleration in magnetic reconnection,/. Geophys. Res., 106, 25,979-25,997, 2001. Ieda, A., S. Machida, T. Mukai, Y. Saito, T. Yamamoto, A. Nishida, T. Terasawa, and S. Kokubun; Statistical anaysis of the plasmoid evolution with Geotail observations, J. Geophys. Res., 103, 4453-4465, 1998. Matsumoto, H., H. Kojima, T. Miyatake, Y. Omura, M. Okada, I. Nagano, and M. Tsutsui; Electrostatic solitary waves (ESW) in the magnetotail: BEN wave forms observed by GEOTAIL, Geophys. Res. Letter, 21, 29152918, 1994. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The Low Energy Particle (LEP) experiment onboard the GEOTAIL satellite,./. Geomag. Geoelectr., 46, 669-692, 1994. Mukai, T., S. Machida, and T. Yamamoto; Dynamics and kinetic properties of plasmoids and flux ropes: GEOTAIL observations, "New Perspectives on the Earth 's Magnetotail" ed. A. Nishida, D. N. Baker and S. W. H. Cowley, Geophys. Mono., 105, 117-137, 1998. Mukai, T., T. Nagai, M. Hoshino, Y. Saito, I. Shinohara, T. Yamamoto, and S. Kokubun; GEOTAIL observations of magnetic reconnection in the near-Earth magnetotail, Adv. Space Res., 25 (No. 7/8), 1679-1684, 2000. Nagai, T., and S. Machida; Magnetic reconnection in the near-earth magnetotail, "New Perspectives on the Earth's Magnetotail" ed. A. Nishida, D. N. Baker and S. W. H. Cowley, Geophys. Mono., 105, 211-224, 1998. Nagai, T., I. Shinohara, M. Fujimoto, S. Machida, R. Nakamura, Y. Saito, and T. Mukai; The structure of the Hall current system in the vicinity of the magnetic reconnection site, J. Geophys. Res., 108, A10, 1357, doi: 10.1029/2003JA009900, 2003. Nishida, A.; The Geotail mission, Geophys. Res. Letters, 25, 2871-2873, 1994. Nishida, A., T. Mukai, T. Yamamoto, S. Kokubun, and K. Maezawa; A unified model of the magnetotail convection in geomagnetically quiet and active times, J. Geophys. Res., 103, 4409-4418, 1998. Nishida, A.; The Earth's dynamic magnetotail, Space Scie. Rev., 91, 507-577, 2000. Seki, K., M. Hirahara, T. Terasawa, T. Mukai, Y. Saito, S. Machida, T. Yamamoto, and S. Kokubun; Statistical properties and possible supply mechanisms of tailward cold O + beams in the lobe/mantle regions,./. Geophys. Res., 103, 4419-44Z9, 1998.
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SECTION 1: Magneto spheric Dynamics
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CLUSTER: NEW VIEW ON THE BOUNDARIES OF THE MAGNETOSPHERE C. P. Escoubet and M. Fehringer ESA/ESTEC, SCI-SH, Keplerlaan 1, 2200 AG Noordwijk, The Netherlands
ABSTRACT After 1.5 years of operations, the Cluster mission is fulfilling successfully his scientific objectives. The main goal of the Cluster mission is to study, in three dimensions, the small-scale plasma stuctures in the key plasma regions in the Earth environment: solar wind and bow shock, magnetopause, polar cusps, magnetotail, and auroral zone. The relative distance between the four spacecraft is varied, according to the scientific region, between 100 and 18000 km during the course of the mission. During the first phase of the mission, the four spacecraft crossed the exterior cusp (Feb. 2001) with an inter-spacecraft distance of 600 km, during the second phase they were in the magnetotail (Aug. 2001) with 2000 km inter-spacecraft distance. Since January 2002, the four spacecraft are again in the cusp/solar wind with the smallest inter-spacecraft distance of 100 km and in June 2002, in the tail, this distance will be increased to 4000 km.. A few results obtained during the first 1.5 years of operation are presented as well as the access to data through the Cluster Science data System and the future operations for the extended mission.
INTRODUCTION Cluster is one of the two missions - the other being the Solar and Heliospheric Observatory (SOHO) constituting the Solar Terrestrial Science Programme (STSP), the first 'Cornerstone' of ESA's Horizon 2000 Programme. The Cluster mission was first proposed in November 1982 in response to an ESA call for proposals for the next series of scientific missions. After the tragic accident of Ariane 5 on 4 June 1996 and the destruction of the four original Cluster spacecraft, the Cluster scientists convinced the ESA Science Programme Committee (SPC) that it was essential for the European scientific community to rebuild the spacecraft. This was agreed by the SPC in April 1997. In the meantime, SOHO, launched in December 1995, had begun to make some exciting discoveries about the Sun and its environment. Now, with the successful launch of the rebuilt Cluster satellites, the STSP Cornerstone is complete and it is possible to combine these two missions in order to study the full chain of processes from the Sun's interior to the Earth. When the first Soyuz blasted off from Baikonur Cosmodrome on 16 July 2000, we knew that Cluster was well on the way to recovery from the previous launch setback. However, it was not until the second launch on 9 August 2000 and the proper injection of the second pair of spacecraft into orbit that we knew that the Cluster mission was back on track. In fact, the experimenters said that they knew they had a mission only after switching on their last instruments on the fourth spacecraft. During the first 1.5 years of operations, Cluster visited the bow shock, the polar cusp, the magnetopause, the plasmasphere, the auroral region and the magnetotail.. In a first section, the Cluster instrumentation will be briefly described, then a few examples of Cluster observations will be presented and finally the data distribution through the Cluster Science Data System will be presented.
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MISSION The scientific objectives of the Cluster mission are to study in three dimensions the plasma structures observed at the bow shock magnetopause polar cusp and magnetotail In addition the temporal variations of structures observed in the auroral zone mid-altitude polar cusp plasmasphere can be studied for the first time as the spacecraft are aligned as a "string of pearls" near perigee. To perform these objectives, the Cluster spacecraft have been placed on a 4x19.6 RE polar orbit (Figure 1). The spacecraft have a slightly different orbit to form a perfect tetrahedron in key regions of space like the Northern polar cusp, Southern magnetopause and plasmasheet (Figure 1)
Fig. 1. Regions of the magnetosphere crossed by the Cluster spacecraft. The left panel shows the orbit in February and the right panel in August. Table 1. Spacecraft separation distances Year 2001 2001 2002 2002 2003 2003 2004 2004 2005 2005
Phase Cusp Tail Cusp Tail Cusp Tail Cusp Tail Cusp Tail
Separation (km) 600 2000 100 4000 5000 100-700 100-700 10000 10000-20000 20000
Fig. 2. Separation distances during the course of the mission
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The separation distance is changed approximately every 6 months to study a particular physical process (Table 1 and Figure 2). It was decided to start with small distances (down to 100 km) and then to increase it toward the end of the mission (up to 20000 km). All measurements at small distances have to be done first since after 20000 km, the remaining fuel will not allow to substantially change the separation distance any more. The Cluster mission has been extended an additional 3 years from January 2003 to December 2005 to cover more separation distances and spend more time at a particular distance. After early 2003, the constellation manoeuvres will be done once a year that means that a tetrahedron will be formed in the tail and 6 months later in the Northern cusp without any manoeuvres in between. This was possible with the innovative manoeuvre method used by the Flight Dynamics Team at the European Space Operations Centre (ESOC). The small distances in the magnetotail in Aug. 2003 were not in the initial planning of the Cluster mission. But were recommended afterwards by the International Space Science Institute substorm working group. The small scales in the tail are necessary to investigate the processes that produce geomagnetic substorms. There are two competing models: magnetic reconnection and current disruption. The existence of a small " diffusion " region where the plasma is rapidly accelerated is expected in the first model, while a disruption of cross-tail current is expected in the second model. Both phenomena have a scale size of approximately 500 km, which will require a spacecraft separation distance of a few hundred kilometres to be studied. This separation distance will be achieved in the first year of extension. In addition, the mission data return has also been augmented by adding a second receiving ground-station. Due to the very large amount of data produced by Cluster, the baseline data return was limited to 50% of the orbit. After a few months of operations, it was however realised that many explosive events were missed due to their unpredictable behavior (sudden storm commencement, substorms, storms) and large scientific regions (like magnetotail and North and South cusp) could not be fully observed. In February 2002, the ESA SPC agreed to extend both the data coverage and the mission. The 100% coverage started in June 2002.
INSTRUMENTATION Each Cluster spacecraft contains a complete suite of instruments to measure magnetic fields, electric field, electromagnetic waves, and particles (Table 2). In addition a potential control instrument keep the spacecraft potential close to a few Volts positive in very tenuous plasma. More details on the payload can be found in Escoubetetal. (2001). Table 2: The 11 instruments on each of the four Cluster spacecraft. Instrument
Principal Investigator
ASPOC (Spacecraft potential control)
K. Torkar (IWF, A)
CIS (Ion composition, 0<E<40 keV)
H. Reme (CESR, F)
EDI (Plasma drift velocity)
G. Paschmann (MPE, D)
FGM (Magnetometer)
A. Balogh (IC, UK)
PEACE (Electrons, 0<E<30 keV)
A. Fazakerley (MSSL, UK)
RAPID (High energy electrons and ions, 20<Ee<400 keV, 10<Ei<1500 keV) DWP * (Wave processor)
P. Daly (MPAe, D)
EFW * (Electric field and waves)
M. Andre (IRFU, S)
STAFF * (Magnetic and electric
H. Alleyne (Sheffield, UK)
fluctuations)
N. Cornilleau (CETP, F)
WBD * (Electric field and wave forms)
D. Gurnett (IOWA, USA)
WHISPER * (Electron density and waves)
P. Decreau (LPCE, F)
* wave experiment consortium (WEC)
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CLUSTER RESULTS During the first months of operations, the spacecraft orbits allowed Cluster to visit key regions of the magnetosphere - the bow shock, the magnetopause and the plasmasphere (Figure 1 left panel). In addition, data obtained during a strong solar storm that occurred in November 2000 will be presented at the end of the section. These results are science highlights of the Cluster mission, for more details see C. Escoubet et al., 2001. The bow shock The bow shock is the surface that forms in front of the Earth's magnetosphere when the supersonic solar wind slams into it at a speed of about 400 km/s (around 1.5 million km per hour). This is similar to the shock wave (or sonic boom) when a plane flies faster than the speed of sound in the atmosphere. The bow shock slows down the solar wind and deflects it around the magnetosphere. In the process, the particles - electron and ions - are heated and the strength of the magnetic field is increased. Intense electromagnetic waves are also produced at the shock. Figure 3 shows the electric waves detected by the WHISPER instrument (Decreau et al., 2001) during two crossings of the bow shock by the four Cluster spacecraft. On the plot, which covers a period of 40 minutes, the wave frequency is plotted as a function of time. The power of the waves is plotted in false colours, red/brown for the most intense and blue for the less intense.
Fig. 3: Bow shock crossings by each of the four Cluster spacecraft on 22 December 2000. The left panels show the frequency-time spectrograms of the electric waves observed by the WHISPER instrument. Between about 08:25 and 08:35 UT the spacecraft were in the solar wind. The right diagrams show the configuration of Cluster during the bow shock crossing, in an overall view (top-right panel) and two enlarged views at 08:23 UT (middle-right panel) and 08:27 UT (bottom-right panel). The spacecraft and trajectories are color-coded: spacecraft 1 (RumbaSC1) in black, Salsa-SC2 in red, Samba-SC3 in green and Tango-SC4 in magenta (adapted from P. Decreau et al.,2001).
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The bow shock is characterised by an intense wave emission enhancement below 20 kHz that is observed around 08:25 UT and 08:35 UT. It should be noted that the crossings do not occur at the same time for each spacecraft. This is explained by the different positions of the spacecraft when they crossed the bow shock. The spacecraft were about 600 km from each other at that time. The right panel of Figure 3 shows the spacecraft configuration at the first crossing, when they were located on the right flank of the bow shock. A closer view of the spacecraft configuration is seen on the middle-right and bottom-right panels. Since the bow shock usually moves at a faster speed than the spacecraft, the spacecraft can be considered immobile and the bow shock is moving through the group. This motion is due to an increase in the pressure of the solar wind on the magnetosphere which pushes the bow shock closer to the Earth. It is clear from the figure that spacecraft 2 crossed the bow shock first and then the other spacecraft followed. A few minutes later, around 08:35 UT, the pressure decreased again and the bow shock crossed the spacecraft in the opposite direction. The second crossing was slightly different since the wave emission appears slightly broader and stronger. Emission of electric waves at the plasma frequency is seen on the spectrograms as a light blue line. From this emission the absolute electron density of the plasma can be deduced. This is shown on the lower panel. At first, the satellites were flying through the magnetosheath (Figure 3), which is characterised by a high particle density around 70 cm"3. After the crossing of the bow shock, the density in the solar wind decreased to around 20 cm'3. An interesting difference between spacecraft 2 and the other spacecraft is observed before the spacecraft again crossed the bow shock at around 08:35 UT. Spacecraft 1, 3 and 4 observed a gradual increase of density from 20 to 70 cm"3 which took about 3-4 minutes, while on spacecraft 2, this increase was very sharp, less than 1 minute. This is an interesting observation which shows that the bow shock can have different properties on quite short spatial scales, typically less than 600 km. Further studies, together with other data, will be conducted to understand the small scale structures of the bow shock.
The magnetopause Inside the bow shock is the magnetopause, the external boundary of the Earth's magnetic field. The magnetopause is characterised by a discontinuity of the magnetic field, a sudden change in particle distribution function and large emission of electromagnetic waves. In the following example, the magnetic field and the wave data will be used to study its geometry. The top panel in Figure 4 shows the north-south component (Bz) of the magnetic field (Balogh et al., 2001) and the total emission power of the magnetic waves (Cornilleau et al., 2003) generated around the magnetopause. The magnetopause is defined by the maximum power of the waves corresponding to the change of sign of Bz. The arrows at the bottom of the plot show the exact time of the magnetopause crossing. It is interesting to note that spacecraft 1 crossed the magnetopause first, then spacecraft 2 and 4 at the same time and finally spacecraft 3. Using a minimum variance analysis, the magnetopause plane can be defined for each spacecraft crossing. The result is shown on the right panel of Figure 4. The individual spacecraft positions as well as the magnetopause plane are shown. It is clear that spacecraft 1, 3 and 4 detected the magnetopause in approximately the same plane while spacecraft 2 detected the magnetopause almost in a perpendicular plane. This observation cannot be explained by a usual planar magnetopause surface but instead by a wave propagating along the magnetopause. The speed of this wave has been estimated at around 70 km/s. It is clear from this example that all four spacecraft are needed to measure the three-dimensional properties of the magnetopause. With two or three spacecraft we could have missed the wave. The plasmasphere The plasmasphere is the doughnut-shaped region located under 5 RE geocentric distance. It consists of very dense and low temperature plasma. The plasmasphere varies in size according to the magnetic activity (Laakso et al., 1997) and sometimes "detached" plasma is observed at its border (Chappell, 1974). On 17 May 2001, the Cluster spacecraft observed a structure at the border of the plasmasphere. Figure 5 shows the WHISPER data (Decreau et al., 2001) on that day. The light blue line marks the plasma frequency. This frequency increases as the local plasma density increases. The spacecraft 1 (top panel) entered the plasmasphere around 1930 UT as the plasma frequency line reaches 20 kHz (equivalent to a density around 4 cm'3).
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Fig. 4. Magnetopause crossing with the four spacecraft on 10 December 2000. The Bz component of the magnetic field is plotted as a function of time (high value means inside the magnetosphere and low value means outside). In addition the integrated wave power from the STAFF instrument is plotted as a function of time. The maximum in the wave power (marked at the bottom by an arrow) indicates the magnetopause crossing. Data courtesy of STAFF Principal Investigator N. Cornilleau-Wehrlin (CETP, France) and of FGM Principal Investigator, A. Balogh (Imperial College, UK).
Fig. 5: Plasmasphere observed by the Cluster WHISPER instruments on 17 May 2001 (same format as Figure 3). Data courtesy of WHISPER Principal Investigator, P. Decreau (LPCE, France). -14-
After 19:30 the plasma frequency continues to increase until 20:00 when it reaches the upper limit of the instrument (80 kHz, equivalent to a density above 80 cm"3). After a little more than 1 hour in the centre of the plasmasphere, the plasma frequency is detectable again (after 21:15 UT on SCI). The four spacecraft are following each other with about 30 min of time delay. Spacecraft 1 leads then come spacecraft 2, 4 and 3. The interesting observation is that the last spacecraft to leave the plasmasphere (number 3) detected a structure (shoulder in the plasma frequency) at around 2310 UT that was not observed by the other spacecraft a few hours before. Further analysis of this event with other instruments and other events with different spacecraft separation is needed to fully characterise the temporal behaviour of these structures. Solar storm on November 2000 With the Sun now at maximum activity in its 11-year cycle, numerous powerful solar storms are expected to occur. On 8 November 2000, the fourth biggest storm since 1976 was detected by SOHO. A huge cloud of plasma, in the form of a Coronal Mass Ejection (CME), was directed toward the Earth (figure 6). About 8 minutes later, the WHISPER instrument on Cluster detected the first consequence of the storm - an intense radio emission from 20 kHz up to above 80 kHz.
Figure 6: Solar storm on 8-10 November 2000. The SOHO LASCO image taken on 8 November 2000 at 23:26 UT is shown on the left panel. The right panels show Wind and Cluster data and the Magneto-Hydro-Dynamic model of the magnetosphere on 10 November 2000 at 05:30 UT (before arrival of the CME) and at 07:00 (at the arrival). The periods when Cluster was outside the magnetosphere are marked in red. Data courtesy of K. Ogilvie (NASA/GSFC), A. Balogh (IC, UK), the SOHO/LASCO team (ESA & NASA) and J. Berchem (UCLA/IGPP, USA).
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Then, about 20 minutes later, the first energetic protons (above 10 MeV) accelerated during the storm arrived at Earth. Their flux was 100,000 times higher than in quiet conditions. The effect of the CME on the magnetosphere took place about 1 day. It acted as a piston on the magnetosphere and decreased its size by half. The right middle and bottom panels of Figure 6 shows the Magneto-Hydro-Dynamics (MHD) model that simulates the magnetosphere's status during the storm. This model reproduces the global interaction of the solar wind with the magnetosphere. All key physical parameters - magnetic field, density and temperature - are calculated in three dimensions using solar wind data measured by the Wind spacecraft upstream of the bow shock. On the top panel, the magnetosphere is shown in dark blue, while to the left the magnetosheath is shown in green/yellow and further to the left the solar wind is in light blue. Due to the increased pressure coming from the solar wind at 07:00 UT, the colour of the above regions changed slightly, the solar wind becoming light yellow and the magnetosheath becoming red. Before the arrival of the cloud, at 05:30 UT, the magnetosphere was normal (right-middle panel) and Cluster was located inside (right top panel). After the CME arrival (right-bottom panel), the magnetosphere was compressed to about half of its normal size (the solar wind pressure reached a maximum of 40 nPa), and Cluster went outside, into the solar wind, for many hours (right top panel). This was about 10 days earlier than expected by the mission team.
Fig. 7: Magnetopause crossing during the solar storm on 10 November 2000. The left diagram shows the magnetic field Bx component as a function of time recorded on the four spacecraft (Rumba-SCl in black, Salsa-SC2 in red, Samba-SC3 in green and Tango-SC4 in blue). Bx is negative inside the magnetosphere and positive outside. The right diagrams show the configuration of the four spacecraft as they crossed the magnetopause in succession. Data courtesy of FGM Principal Investigator, A. Balogh (Imperial College, UK). Figure 7 shows the magnetometer data from one excursion outside the magnetosphere during the storm. Each spacecraft successively crossed the magnetopause in agreement with the model shown on the right hand. The order of the crossing, starting at 06:02 UT was first spacecraft 2 (Salsa-SC2), then Tango-SC4, Samba-SC3 and finally
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Rumba-SCl. A few minutes later, starting at 06:03:30, Rumba-SCl, Samba-SC3, Tango-SC4 and Salsa-SC2 reentered the magnetosphere in succession. The speed of the spacecraft was relatively slow compared to the motion of the magnetopause, so it should be seen as if the magnetopause was going back and forth through the spacecraft. The order of exit and entry is reversed, indicating that the magnetopause kept the same orientation during its motion. CLUSTER SCIENCE DATA SYSTEM The Cluster science data system has been set-up to distribute quicklook and processed data to all Cluster Principal and Co-Investigators, as well as to the scientific community. The Cluster community consists of 11 Principal Investigators and 241 Co-Investigators from 73 laboratories all over the word (Figure 8). To distribute the data efficiently to all users, a system of nine data centres located in Austria, China, France, Germany, Hungary, Netherlands, Sweden, United-Kingdom and United-States (see Figure 8) and interconnected with each other has been set-up. Each data centers store the full database of processed and validated data from all instruments. Data from February 2001 until end of December 2002 are available through the web at (http://sci2.estec.esa.nl/cluster/csds/csds.html). A quicklook plot is also available between a few hours to a few days after data acquisition and includes time series plots and spectrograms from most of the instruments. This is very useful for scientists to pick-up interesting events and for the Cluster project to monitor the progress of the mission. The access to data has started in February 2001 with a download rate of 700 Mbytes/month and since then has been increasing steadily at an average rate of about 150 Mbytes per month. The average download rate for the first three months of the year 2003 was above 6.5 Gbytes/month.
Fig. 8: Location of the Cluster community and data centres. CONCLUSION During the first few months of operations, the Cluster spacecraft have shown their full capability to make substantial advances in magnetospheric physics. For the first time structures in the magnetosphere have been studied in three dimensions which will bring new knowledge of processes taking place during the interaction between the Sun and Earth. The bow shock was captured between the four spacecraft, enabling its geometry and speed to be determined for the first time. Waves on the magnetopause were also directly observed for the first time and further studies of these should bring new insights into magnetic reconnection processes. Plasma structures were observed in the -17-
plasmasphere on a single spacecraft which suggest their temporal behavior. More data have been obtained in these regions and will allow scientists around the world to perform systematic studies of the physical processes involved. Another main target of the Cluster mission is the magnetotail, where magnetic reconnection, current disruptions and particle acceleration are taking place. During the next few months, Cluster will look for the first time at the spatial variation of these processes, casting new light on geomagnetic substorms that are responsible for the intense auroras on the nightside of the Earth. The latest news and the access to Cluster data through the Cluster Science Data System can be found at: http://sci.esa.int/cluster/ ACKNOWLEDGMENTS The authors thank all PI and their teams who provided the Cluster data, and the JSOC and ESOC team for their very efficient Cluster operations.
REFERENCES
Balogh A., C M . Carr, M.H. Acuna, M.W. Dunlop, T.J. Beek, P. Brown, K.-H. Fornacon, E. Georgescu, K.-H. Glassmeier, J. Harris, G. Musmann, T. Oddy, and K. Schwingenschuh, The CLUSTER magnetic field investigation: overview of in-flight performance and initial results, Ann. Gophys., 19, 2001. Chappell, C. R., Detached plasma regions in the magnetosphere, J. Geophys. Res., 79, 1861, 1974. Cornilleau-Wehrlin N., Chanteur G., Perraut S., Rezeau L., Robert P., Roux A., et al., First results obtained by the Cluster STAFF experiment, Ann. Gophys., 21,437, 2003. Decreau P. M. E., P.Fergeau, V.KrasnoselsTdkh, E.Le Guirriec, et al., Early results from the Whisper instrument on Cluster: an overview, Ann. Gophys., 19, 2001 Escoubet, C. P., M. fehringer, and M. Goldstein, The Cluster mission, Ann. Gophys., 19, 2001. Laakso, H, H. Opgenoorth, J. Wygant, P. Escoubet, J. Clemmons, M. Johnson, N. Maynard, F. Mozer, R. Pfaff and J. Scudder, Electron density distribution in the magnetosphere, ESA-SP-415, 1997.
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THE STRUCTURE OF THE PLASMA SHEET UNDER NORTHWARD IMF M. Pujimoto1, T. Mukai2, and S. Kokubun1 l
Dept. Earth Planet. ScL, Tokyo Inst. Tech., Meguro, Tokyo 152-8551, JAPAN 2 ISAS, Sagamihara, Kanagawa 229, JAPAN ABSTRACT
The structure of the plasma sheet under northward IMF is studied by data from the Geotail spacecraft. The plasma sheet is known to become cold and dense (T < 2 keV, n > 1 cm"3) during extended northward IMF periods. We show that such cold-dense ions (CDIs) appear at |Y"9sm| > 10 Re, that is, 10 Re off the tail axis to both flanks. CDIs on the dawnside have higher temperatures and some reach the upper-limit of 2 keV that we use to select CDIs. These CDIs having the highest temperature are distributed at the dawnside plasma sheet inner-edge (R < 15 Re) and is connected to the hot-dense ions (HDIs: T > 2 keV, n > 1 cm"3) in the further inner region. A survey shows that HDIs under nominal solar wind dynamic pressure appear mostly in the dawnside inner-magnetosphere during extended northward IMF intervals. Both results point to the idea that HDIs are the inner-magnetosphere extension of dawnside CDIs, while such a partner to the duskside CDIs cannot be identified. This structure of the plasma sheet suggests that there is significant dawn-dusk asymmetry in heating and transport in the magnetotail under northward IMF.
INTRODUCTION The plasma pressure in the plasma sheet is in balance with the magnetic pressure in the lobe and thus in balance with the dynamic pressure exterted on the flared magnetopause surface by the solar wind. It has been well known that the density and the kinetic energy of the solar wind control the density and the temperature, respectively, of the plasma sheet (e.g., Borovsky et al, 1997). The pressure balance requirement is naturally satisfied by this property. For typical cases, the plasma sheet density divided by the solar wind density is 0.05. Under the same solar wind condition, if the plasma sheet density is higher, its temperature has to be lower. If there is a factor that increases this density ratio, the plasma sheet would become denser and colder when the factor is in action. It has been known (Lennartsson and Shelley, 1986; Baumjohann et al., 1989) and has been confirmed very clearly by a recent statistical study using data from WIND and Geotail (Terasawa et al., 1997) that the plasma sheet status is controlled by the north-south component of IMF, cold and dense for extended northward IMF periods. By mapping data from low-altitude DMSP satellites, Wing and Newell (2002) has shown the spatial extent of the cold and dense ions under northward IMF. Understanding the formation mechanisms of the cold-dense plasma sheet (CDPS) that appears during extended northward IMF periods is challenging, for this is to unveil the unknown more efficient (note that the density ratio is increased) solar wind entry that results in a totally different state of the plasma sheet, and because this more efficient entry process is operative during the times when the magnetosphere is supposed to be most closed. The cold-dense status is at times realized by the presence of mostly unheated magnetosheath-like component that is totally different from the usual plasma sheet component (Fujimoto et al., 1998). As such, the entry processes under northward IMF should be able to accomplish transport across the boundary with little heating. The Geotail spacecraft has been collecting the data in the plasma sheet for more than seven years. Concurrent upstream solar wind/IMF data are available from WIND and ACE spacecrafts for these years. With these data, we are in a position to reveal the structure of the plasma sheet and its control by IMF. In this study, we show in what shape the plasma sheet is under extended northward IMF when the unknown -19-
Fig. 1. Variation of the plasma sheet status. All the plasma sheet data are binned and the detection probability of each bin is shown by the color code. The bins are logarithmically spaced, (a) For data taken close to the tail axis \Ygsm\ < 10 Re and (b) for data obtained close to the flanks \Ygsm\ > 10 Re. While the most probable density is < 0.5 cm~3 in both plots, the flank plasma sheet has more tendency to put on higher density/ lower temperature ( > lcm~ 3 , < 2 keV).
entry mechanism is operative. By compiling the plasma sheet data obtained in 1995-99, we show that, while the density is typically < 0.5 cm" 3 , it can be as dense as > 1 cm" 3 in the flanks. The pressure balance requires the temperature to be < 2 keV for these occasions. As such, these data points constitute Cold-Dense-Ions (CDIs: > 1 cm" 3 , < 2 keV) that are known to appear during northward IMF. CDIs are found to show dawn-dusk asymmetry when their temperatures are plotted versus Ygsm. Those samples having the highest temperature close to the 2 keV upper-limit of CDIs are distributed at the dawnside inner-edge of the plasma sheet R < 15 Re (R: the geocentric distance). We also find that Hot-Dense-Ions (HDIs: > 1 cm" 3 , > 2 keV) are distributed on the earthward side of this inner edge during northward IMF. Any counterpart on duskside to the HDIs cannot be identified. While composing this structure by compiling the five years of data, we use data obtained under steady northward IMF intervals from appropriate orbits to snapshot the key parts of this structure and strengthen the argument that such a structure is indeed spatial. The dawn-dusk asymmetry in the structure obtained above would reflect the dawn-dusk asymmetry in the plasma heating and transport processes under northward IMF. THE PROPERTIES OF COLD-DENSE PLASMA SHEET In this paper, we study the Geotail spacecraft data obtained in 1995-99. A data point is identified as a plasma sheet data point if it is obtained in the tail Xgsm < -15 Re and satisfies (1) \Vx\ < lOOkm/s, (2) density > 0.1 cm~3 and (3) temperature > 0.1 keV. Figure 1 shows how often various states of the plasma sheet appear. Shown is the color contour of the detection probability in log n - log T format. Figure la is for the part of the plasma sheet close to the tail axis |Yj5m| < 10 Re. One can see that the most probable plasma sheet density is < 0.5 cm~3. Figure lb shows the same thing for the region close to the flanks \YgSm\ > 10 Re. While the most probable density stays below 0.5 cm~3, now one can see that the density has non-negligible tendency to put on higher values in this region. When the density is above 1 cm~3, the temperature is below 2 keV, and thus satisfies the CDPS criteria. That is, CDPS seems to be detected more often closer to the flanks. In other words, when IMF turns northward, the part of the plasma sheet close to the flanks undergoes most dramatic change in its status. To study the plasma sheet structure under northward IMF, instead of sorting out the data obtained during northward IMF periods, we look for CDPS intervals, which are known to develop when the IMF condition persists. That is, we will be studying a well-matured state of the plasma sheet that would appear during -20-
Fig. 2. Spatial distribution of CDPS. One can clearly see that CDPS samples appear more often close to the flanks.
extended northward IMF intervals (The time required for its formation is not addressed in this study.). Among all the data obtained at Xgsm < 2.5i?e, the CDPS intervals are selected as follows: We first divide the Geotail dataset into 1-hour segments. Then every data point (12 sec. resolution) in each segment is inspected and if it satisfies the criteria (1) density n > 1 cm~3, (2) ion temperature T < 2 keV, and (3) the x component of the ion bulk flow Vx > -100 km/s, it is recorded as a CDPS data sample. The third condition is to reject data points from tailward flowing LLBL and the magnetosheath. We finally count up the number of CDPS samples and if they occupy more than 30 min. of their parent 1-hour segment, the 1-hour segment is identified as a CDPS interval. In this paper, in order to focus on prominent/long-duration events, we select only those CDPS intervals that last for more than 3 hours consecutively. Figure 2 shows the spatial distribution of the CDPS intervals selected in this way. Each point depicts the spacecraft position at the start time of each 1-hour CDPS segment. The region Xgsm < —30Re and R < WRe is not covered by the present dataset. It can be seen that the long-duration CDPS events are distributed widely over the Xgsm range of interest and are concentrated off the tail axis, that is, at | Ygsm | > 10 Re. This spatial distribution agrees well with what is shown in Figure 1 using all the data obtained within the plasma sheet. IMF Bz tended to be northward during and a few hours prior to these intervals. The points are distributed mostly equally on both flanks and we see no appreciable dawn-dusk asymmetry at this point. By inspecting the ion energy spectrum features, we have noticed that the CDPS events can be classsified into three categories. The three categories are as follows:(l) Category 1: CDPS composed of two-component ions (Fujimoto et al., 1998; Fuselier et al., 1999; Phan et al., 2000). This category is identified by the ion spectra feature showing a count rate peak locally below 0.75 keV and a significant flux at 5 keV F^eV > F* (F* corresponds to a differential energy flux at 5 keV for a (n, T)=(0.2 cm~3, 3 keV) Maxwellian). (2) Category 2: CDPS composed mostly solely of low energy ions. This would be the coldest state of the plasma sheet and is identified in the same way as Category 1 but without requiring F^eV > F*. (3) Category 3: CDPS without distinctly separated low energy ions (Phan et al., 1998). Cold (T < 2 keV) and dense (n > 1 cm~3) samples that do not have a local count rate peak at < 0.75 keV fall into this category. While Category 3 may contain two-component ions whose lower energy peak is above 0.75 keV, this category in reality is composed mostly of ions showing a single component feature whose peak is above 0.75 keV. By re-plotting Figure 1 with different categories depicted by different symbols, we found the Category 2 (cold component only) to appear on both dusk and dawnside, implying magnetosheath ions can access the plasma sheet with little heating from both the flanks. In contrast, Category 1 (two component) appears mostly only (90 %) on duskside whereas Category 3 (no cold component) are seen more (75 %) on dawnside (See -21-
Fujimoto et al. (2002) for details.). Here we note that the condition F^eV > F* is satisfied (although not required) in all of the Category 3 samples. Since the Categories 1 and 3 contain more energetic ions than the Category 2, they are likely to be obtained when the magnetospheric activity is higher than that for the Category 2 intervals. When such a slight activity exists, it results in heating of the sheath-like component on dawnside (Category 3), but the sheath-like population on duskside remains distinctly separated from the more energetic component that is presumably originated in the distant tail (Category 1). The symmetric distribution of Category 2 (cold ions only) implies that cold ions can access from dawnside as well, and CDIs that are categorized as Category 3 on dawnside would have had the spectrum shape that is identified as Category 1 close to the entry point presumably at the dawnside magnetopause. Quick heating of the cold component, however, seems to transform the spectrum to a one-component shape that is identified as Category 3. Note that this heating is only slight in the sense that it does not bring the temperature to above 2 keV but is enough to change the energy spectrum feature. The spatial distribution of Category 1 and 3 suggests that the heating is operative almost always on dawnside but infrequently on duskside in a slightly active magnetosphere. This dawn-dusk asymmetric nature of heating becomes more prominent when the temperature is plotted versus Ygsm. This plot for Category 3 (not shown) shows data samples having temperatures T > 1.5 keV to appear preferentially on dawnside. The same plot for Category 1 shows that most of the data points are distributed below 1 keV indicating that samples with > 1.5 keV are rare. Since the energetic component does exist in Category 1. this result suggests that the presence of this component is not enough but additional heating of the cold component is required to make T above 1.5 keV. That is, under northward IMF, dense samples (n > 1 era"3) with moderate temperature 1.5 keV < T (< 2 keV) belong mostly exclusively to Category 3 and appear only on dawnside where some heating of the cold component is present. By pursuing further into these higher temperature samples, we found that they are obtained at the dawnside inner-edge of the plasma sheet R < 15 Re. Since these distances are among the smallest, it is likely that these ions are heated more as they approach more to the Earth from dawnside. By approaching further closer to the Earth, the ions may become hot-dense ions (HDIs: n > lcm~3 and T> 2 keV) in the dawnside inner-magnetosphere. We will investigate this point in the next section.
Fig. 3. Variation of the plasma status in the tail region with geocentric distances 10 < R < 15 Re. (a) Duskside Ygsm > — 0-3X 9sm and (b) dawnside Ygsm < 0.3Xffsm. The presence of hot-dense ions (> 2 keV, > 1 cm'3) is evident in the plot for dawnside.
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Fig. 4. Locations of the long-duration HDI intervals are plotted by plus symbols. It is clear that they are all at the inner-edge of the plasma sheet R ~ 12 Re. Solid squares show long-duration HDI under nominal solar wind conditions, that is, a replotting of the data with HDIs under elevated solar wind density and/or dynamic pressure being excluded. All of the HDI under nominal solar wind conditions are on the dawnside inner-edge of the plasma sheet.
HOT-DENSE IONS IN THE DAWNSIDE INNER-MAGNETOSPHERE In this section hot-dense ions (HDIs) in the magnetotail are investigated. First we show the state of the tail inner-magnetosphere by plotting the detection probability in the log n - log T plane. Here the data obtained in the tail at the geocentric distances of 10 < R < 15 Re are considered. Note that because of the Geotail orbit, the region R < 10 Re is not investigated. Magnetosheath/LLBL are excluded as before. Figure 3 shows the results with dawnside (Ygsm < Q.3Xgsm) and duskside (Ygsm > — 0.3Xgsm) plotted separately. One can see that the typical density in this region is 0.5 cm" 3 on both sides. With smaller probabilities, higher density > 1 cm" 3 samples appear preferentially on danwside. By noting that some of them have temperatures above 2 keV, we can tell that HDIs are present at R < 15 Re on dawnside. Let us now focus on the HDIs. The procedure for selecting HDI intervals is the same as CDPS in the previous section except that the criteria for the ion temperature is now T > 2 keV (instead of < 2 keV). We plot in Figure 4 the spatial distribution of the HDI intervals by plus symbols. We plot only those detected consecutively for more than 3 hours. Excluding one exceptional event, all the long duration HDI intervals are on dawnside at the radial distances of ~ 12Re. Since the region closer than 10i?e to the Earth is not surveyed by Geotail, the inner limit of HDI cannot be determined by the present dataset.
Fig. 5. IMF Bz for HDIs under nominal solar wind condition. T=0 are the start times of the events. Data for T=—6 to +6 hours are shown. Tendency for northward IMF from 3 hours prior to and during the events can be seen.
Inspection of the concurrent solar wind condition (WIND KP data, courtesy, A. J. Lazarus) shows that some of the long-duration HDI events plotted in Figure 4 are associated with elevated solar wind density -23-
and/or dynamic pressure. Since these cases are more or less trivial, we exclude the HDI events that are associated with solar wind density exceeding 20 cm~3 or dynamic pressure exceeding 5 nPa. We re-plot by solid squares in Figure 4 the long duration HDI events detected under normal solar wind condition. It can be seen that the exceptional case on duskside is due to the elevated solar wind density and long duration HDIs under normal solar wind condition are distributed only on dawnside. Figure 5 shows the IMF condition (WIND KP data, courtesy, R. Lepping) for normal solar wind density HDI events. T=0 in this plot corresponds to the start time of the HDI events and the IMF Bz component from T=—6 to +6 hours interval are shown. Time lags are calculated by Xgsmtw/Vsw {Xgsm^y. GSM X of WIND, Vsw: concurrent solar wind speed) and are taken into account in plotting Figure 5. The plot suggests a tendency for these long duration HDIs to appear during extended northward IMF. It should be noted that CDPS also tends to appear during northward IMF (e.g., Terasawa et al., 1997). Indeed, this study on HDI is initiated by the idea that a category of CDIs in the plasma sheet might transform to HDIs in the inner magnetosphere while they experience more heating in the way. The fact that both CDPS and HDIs favor northward IMF implies that both appear under northward IMF with HDI being the inner-magnetosphere extension of CDPS.
"SNAPSHOTS" TAKEN UNDER STEADY NORTHWARD IMF The above statistical analyses reveal the structure of the plasma sheet under northward IMF as follows: CDIs appear on both flanks. CDIs on the dawnside are heated as they approach the inner-magnetosphere. The temperature of CDIs on duskside stays below 1 keV, suggesting no heating and thus no approach to the inner-magnetosphere is in action on this side of the plasma sheet. This asymmetric transport leads to the asymmetric distribution of the HDIs in the inner-magnetosphere. More than 60 extended CDI and 9 extended HDI intervals obtained in 1995-99 set the basis for this conclusion.
Fig. 6. A schematic figure to show the appropriate orbits from which snapshots of the plasma sheet structure under northward IMF can be taken, (a) Skimming the plasma sheet inner-edge, mostly symmetric about midnight, (b) From the dawnside plasma sheet inner-edge to deeper into the dawnside flank. If data are taken under steady nominal solar wind and northward IMF conditons, they should snapshot (a) dawn-dusk asymmetry of the HDIs and (b) CDI-HDI connection on dawnside, respectively.
The above picture will be clearer if we inspect data from appropriate orbits under stable solar wind/IMF condition that snapshot the spatial structures described above. In the followings, we show two examples -24-
that indeed show (1) the dawn-dusk asymmetry of HDIs and (2) the connection of CDIs to HDIs at the plasma sheet inner-edge on dawnside. The two orbits are schematically shown in Figure 6. To show the dawn-dusk asymmetry of HDIs, we look for HDIs obtained along the orbits that skim the inner-edge of the plasma sheet nearly symmetric about midnight. We then check if there is any temporal variation in the solar wind/IMF condition during the traversals of the magnetotail. If not, that would imply spatial dawn-dusk asymmetry. Figure 7a shows the data from a traversal selected in this manner. Shown are data obtained in 24 hours during which the spacecraft moved all the way from duskside LLBL slightly tailward of the terminator (0000 UT), across the midnight meridian at ~ 10 Re in the tail (1200 UT), and finally to the dawnside terminator magnetopause (2300 UT). The data in 0400-0700 UT show the duskside low-latitude (small \BX\, 3rd panel) plasma sheet at the radial distances of 13 Re and MLT=19-21. The density (1st panel) remains small (< 1 cm~3) and the temperature (2nd, in keV unit) stays high (> 4 keV) in this region. As such, this duskside plasma sheet inner-edge shows the typical status. The density is elevated to above 1 cm" 3 at 1600 UT when the spacecraft is on dawnside (MLT=2.5). The temperature keeps > 2 keV and thus the criteria for HDIs are satisfied. The interval that strictly satisfies the HDI criteria described before spans 1600 - 2300 UT. It is evident that the Bx (3rd) and By (4th) components have only moderate magnitudes indicating that the HDIs are situated at low-latitude part of the plasma sheet inner-edge on dawnside. As already described, it is the low-latitude part of the plasma sheet inner-edge on both sides that are surveyed by the spacecraft. The striking contrast between dusk and dawn seen in Figure 7a is not due to difference in latitudes. To confirm the connection between CDIs and HDIs on dawnside, we look for HDIs on the orbits that is at ~ 10 Re from the Earth on the midnight meridian and then move tailward into the dawnside plasma sheet. In Figure 7b, which shows the data obtained on such an orbit under stable solar wind/IMF condition. HDIs are identified in 0300-0600 UT when the spacecraft is on the dawnside plasma sheet inner-edge. Then, as the spacecraft moves deeper tailward into the dawnside plasma sheet, CDIs are detected at 0600-1000 UT. This finding of HDIs and CDIs being situated adjacent to each other strongly supports the idea that HDIs are the inner extension of the CDIs.
Fig. 7. Data from the orbits described in Figure 6. (a) HDIs on dawnside (1600-2300 UT) but no such a counterpart on duskside. (b) CDIs sitting adjacent to HDIs on dawnside.
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DISCUSSION In this paper, we have investigated the structure of the plasma sheet under northward IMF. CDPS (n > 1 cm~3 and T < 2 keV) is well known to appear under northward IMF. We have shown that they appear mostly in the region more than 10 Re off the tail axis to the flanks. By inspecting the ion energy spectrum features, we have found that the CDIs can be classified into three categories. One of the categories that preferentially appears on dawnside is indicative of a heating process on low energy (< 1 keV) ions. A part of these showing temperatures > 1.5 keV are located at the dawnside inner-edge of the plasma sheet. On the earthward side of this region, present under northward IMF are HDIs (n > 1 cm~3 and T > 2 keV). CDIs in the plasma sheet are transformed to HDIs as they are heated more in the course of earthward convection. The fact that there are no HDIs on duskside implies that such supply of plasma from the duskside plasma sheet is not present under northward IMF. Indeed, duskside CDIs are characterized by a spectrum feature that indicates absence of heating on the cold component. The dawn CDPS-HDI connection can be explained by the test particle transport calculations by Spence and Kivelson (1994) and Wang et al. (2001). Assuming two sources for the plasma sheet ions, one in the distant tail that produces the normal energetic component and the other at the dawnside flank, these calculations follow the adiabatic motions of test ions in model magnetic and electric fields. An ion's thermal energy is calculated from the adiabatic relation. For quite intervals, the effects of the convection electric fields do not dominate and the cold component drifting duskward from the dawn-flank can substantially contribute in populating the plasma sheet. This produces CDPS close to the dawn-flanks and, since they are heated adiabatically as they move inward and duskward, produces HDI in the dawnside inner-magnetosphere as well. In the model, a cold ion source at the dusk flank is not situated. Even if it is situated, the dawn-todusk directed motion will make it not effective in supplying ions over a wide area in the plasma sheet. Thus the duskside plasma sheet inner-edge is dominated by the ions from the distant tail source and shows the normal hot and tenuous status. While the agreement between the observations and the model calculations may seem simple, it points to the two items that do not seem to be widely accepted in the magnetospheric community: (1) The dawn flank is a substantial source for the plasma sheet during quiet times. (2) The effects of non-E x B drifts can be significant in shaping the spatial distribution of (not only the energetic but also) the thermal population during quiet times. These two points have been theorized by Spence and Kivelson (1994) but are now supported by a data-based study. An inferred two-dimensional structure of the plasma sheet by mapping DMSP data seems to point to the same conclusion (Wing and Newell, 1998). One question that is not answered by these models are, of course, the formation mechanism of duskside CDPS. The above agreement with the adiabatic model calculations makes us expect gradual change from CDIs to HDIs. The example in Fig. 7b, however, is puzzling in this aspect. One can see that the transition at 0600 UT is rather sharp, not in accordance with the slow adiabatic picture. Furthermore, the magnitude of the magnetic field is not larger on the HDI side of the boundary. There seems more to be considered on the heating mechanism. Let us note on the fate of the HDIs. If they drift duskward, we should be detecting HDIs on duskside as well. Since we do not, this is not the case. It could be that, in the dawn-tail inner magnetosphere, the weak electric field convection under northward IMF that tries to move the particles dawn-sunward and the weak grad-B/curvature drift of lower energy ions that is directed dusk-tailward are mostly balanced. Then the lower energy ions that carry most of the high density may be stagnant in the dawn-tail part. If not stagnant, because they are absent on duskside, they should be moving dawn-sunward. In this case, we would expect the dawn-dusk asymmetry to be seen in the dayside as well. ACKNOWLEDGMENTS M. F. acknowledges fruitful discussion with T. Phan, H. Hasegawa, and M.N. Nishino. The key parameter data from WIND spacecraft were provided by the NASA/GSFC data processing team. REFERENCES Baumjohann, W., G. Paschmann, and C. A. Cattell, Average plasma properties in the central plasma sheet, J. Geophys. Res., 94, 6597, 1989.
Borovsky, J. E., M. F. Thomsen, and D. J. McComas, The superdense plasma sheet: Plasmasperic origin, solar wind origin, or ionospheric origin?, J. Geophys. Res., 102, 22,089, 1997. -26-
Fujimoto, M., et al, Plasma entry from the flanks of the near-Earth magnetotail: Geotail observations J. Geophys. Res., 103, 4391, 1998. Fujimoto, M., T. Mukai, and S. Kokubun, The cold-dense plasma sheet and the hot-dense ions in the inner magnetosphere, Adv. Space Res., 30 (10), 2279, 2002. Fuselier, S. A., et al, Composition measurements in the dusk flank magnetosphere, J. Geophys. Res., 104, 4515, 1999. Lennartsson, W. and E. G. Shelley, Survey of 0.1 to 16 keV/e plasma sheet ion composition, J. Geophys. Res., 91, 3061, 1986. Phan, T. D., et al., WIND observations of the halo/cold plasma sheet, Substorms-4, Eds. S. Kokubun and Y. Kamide, Terra Sci. Pub., Tokyo, 1998. Phan, T. D., et al., WIND observations of mixed magnetosheath-plasma sheet ions deep inside the magnetosphere, J. Geophys. Res., 105, 5497, 2000. Spence, H. E., and M. G. Kivelson, Contributions of the low-latitude boundary layer to the finite width magnetotail convection model, J. Geophys. Res., 98, 15,487, 1993. Terasawa, T., et al., Solar wind control of density and temperature in the near-Earth plasma sheet: WINDGEOTAIL collaboration, Geophys. Res. Lett, 24, 935, 1997. Wang, C. P., et al., Modeling the quiet time inner plasma sheet protons, J. Geophys. Res., 106, 6161, 2001. Wing, S. and P. T. Newell, Central plasma sheet ion properties as inferred from ionospheric observations, J. Geophys. Res., 103, 6785, 1998. Wing, S. and P. T. Newell, 2D plasma sheet ion density and temperature profiles for northward and southward IMF, Goephys. Res. Lett, 29, 10.1029/2001GL013950, 2002.
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GEOTAIL OBSERVATIONS OF THE COLD PLASMA SHEET ON THE DUSKSIDE MAGNETOTAIL M. N. Nishino1, T. Terasawa1, and M. Hoshino1 1
University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo, 113-0033, JAPAN ABSTRACT
We have studied signatures of the cold plasma sheet in the duskside near-earth magnetotail with the Geotail data. In the duskside plasma sheet, cold plasmas are found 3-4 hours after the northward turning of the interplanetary magnetic field. The cold plasma sheet with two-temperature ions is often found on the duskside, while cold plasma sheet with one-temperature are also found on the duskside when northward interplanetary magnetic field continues for a very long interval (more than several hours). We discuss the development and evolution of the cold plasma sheet on the duskside.
INTRODUCTION The origin of the plasma particles in the near-earth plasma sheet has not been fully understood. For southward interplanetary magnetic field (IMF) condition, it is widely accepted that magnetic reconnection plays an important role in plasma entry from the magnetosheath into the plasma sheet and plasma heating there. On the other hand, two mechanisms have been proposed to explain plasma transport during the northward IMF periods; that is, magnetic reconnection at high-latitude magnetopause and diffusive process at the low-latitude boundary layer (LLBL). However, what is really going on at the magnetopause and in the plasma sheet under the northward IMF conditions is still an open question. Several authors (Fairfield et al., 1981; Lennartsson and Shelly, 1986; Baumjohann et al., 1989; Lennartsson, 1992) reported that the near-earth plasma sheet becomes cold and dense during geomagnetically quiet periods. In recent years several studies have suggested that plasma particles which are of magnetosheath origin enter directly into the plasma sheet via near-earth tail flanks (e.g. Fujimoto et al., 1997; Terasawa et al., 1997; Borovsky et al., 1997). Nishino et al. (2002) reported that the total plasma content in the near-earth plasma sheet increases during northward IMF periods, and they concluded that the effect of the plasma transport during the northward IMF intervals is not negligible. These studies have shown that the cold plasma sheet (Cold-PS) is formed in the tail flanks during the northward IMF periods. Fujimoto et al. (1997) concluded that cold ions of magnetosheath origin directly enter into the flank plasma sheet via the magnetopause to form the Cold-PS with a two-temperature velocity distribution function. This two-temperature state shows that hot ions of magnetospheric origin and cold ions of magnetosheath origin co-exist in the flank plasma sheet. Recently, Takashima (2002) reported that the Cold-PS with twotemperature ions exists only on the duskside flank, and that not only IMF-Bz but also IMF-By plays an important role in formation of a Cold-PS with two-temperature ions. In this report we focus on the development and evolution of the Cold-PS on the duskside near-earth magnetotail flank and its dependence on the northward IMF. INSTRUMENTATION For the plasma sheet data, we use 12 s averaged magnetic field data from the magnetic field experiment (MGF) (Kokubun, et al., 1994) and 12 s averaged plasma moment data (density, velocity, and temperature) from the low energy particle experiment (LEP) (Mukai et al., 1994) on board Geotail. For the solar wind -28-
Table 1. Event list of cold plasma sheet in the dusk flank magnetotail. From left, date, T\Ag (see text for its definition), three components of the IMF, solar wind density (Nsw), GEOTAIL position, and ion velocity distribution (one-/two-temperature) in Cold-PS are shown. No. (1) (2) (3) (4) (5) (6) (7) (8) (9)
date yyyy/ m m /dd 1995/02/24 1995/03/07 1995/03/24 1995/04/04 1995/05/01 1996/02/22 1996/04/01 1997/02/12 1997/02/24
Tlag
IMF Bx
hour
nT
4.3 5.8 8 30 7.8 7.5 2.6 9.3 13
-0.8 ±1.3 -0.2 ±1.6 1.4 ±1.5 -2.8 ±3.9 -2.2 ±1.3 3.9 ±0.7
3.1 ±0.7 -1.0±0.9 2.2 ±3.4
IMF By nT 2.2 ±1.9 2.8 ±2.4 -2.6 ±2.6 2.9 ±2.8 2.4 ±1.2 0.9 ±4.3 0.1 ±1.2 0.7 ±1.2 1.2 ±2.5
IMFBz nT
Nsw /cc
1.9±1.0 19.0 ±0.8 4.6 ±1.5 4.8 ±0.4 9.2 ±1.4 16.6 ±3.0 6.8 ±2.2 4.5 ±1.0 2.2 ±1.3 6.0 ± 0.6 2.1 ±1.9 5.1 ±2.0 2.4 ±0.6 20.8 ±2.2 2.5 ±0.8 7.2 ±1.6 2.5 ±2.0 5.1 ±1.2
GEOTAIL
Ions
(X,Y,Z)RE (-26.1,20.1,7.4) (-18.0,18.2,8.2) (-16.1,16.3,1.8) (-10.8,19.7,9.9) (-7.7,17.4,4.5) (-18.8,23.3,4.0) (-16.5,18.5,3.2) (-17.3,20.3,8.2) (-21.7,20.5,4.7)
2-Temp. 2-Temp. 2-Temp. 1-Temp. 2-Temp. 1-Temp. 2-Temp. 1-Temp. 1-Temp.
data, 92 s averaged data from WIND/SWE and MFI provided via CDAWeb are used. We use the geocentric solar magnetosphere (GSM) coordinate throughout this report. OBSERVATION We have surveyed the Geotail data during January 1994 - December 1997 to find Cold-PS in the nearearth magnetotail (0 > X > —50 RE)- The criteria for Cold-PS are Tj(ion temperature)< 1 keV and /3(ratio of magnetic and plasma pressure) > 1. We do not impose a condition on the ion density (N\), because the ion density in the plasma sheet depends on the solar wind density as well as the distance from the earth. To avoid contaminations of the low-latitude boundary layer (LLBL) and the magnetosheath, the events with tailward flowing ions (Vx < —100 km/s) in the flanks are excluded. In our data set we have selected 9 typical Cold-PS events near the dusk magnetopause (Table 1). To study the effect of the length of northward IMF duration on formation of Cold-PS, we introduce T]ag which represents the time delay between the northward turning of the IMF and the time when Cold-PS was first seen. In defining the beginning of a northward IMF interval we require 30 minutes of southward field and then we select the following northward field interval that is closest to the Cold-PS observation. We have neglected short (interval of < 30 min) and small (time integration of < —30 nT-min) southward excursions of the IMF during prolonged northward IMF intervals. We have checked the ion velocity distribution function of these Cold-PS events and divided them into two groups; that is, two-temperature Cold-PS (Figure 1 (a)) and one-temperature Cold-PS (Figure 1 (b)). The IMF-Bz for each event is shown in Figure 1 (convection delay included), where the intervals of T]ag are shown as hatched regions. It seems that the duskside plasma sheet is filled with two-temperature ions if the northward IMF continues for 3-4 hours, and that the ions there become one-temperature if the northward IMF holds for more than several hours. The dependence of these Cold-PS emergence on IMF-By j if any, can not be discussed with our data set, because IMF-By was positive in all cases except one. Here we focus on an event of Cold-PS with one-temperature ions. Figure 2 shows the Wind and Geotail observations on April 2-4, 1995 (Event No. 4 in Table 1). The left panels (a) show the Wind data between 0 UT of April 2 and 8 UT of April 4. From the top, magnitude of magnetic field (\B\) and its Z-component (Bz), the latitudinal angle of magnetic field ( 30°), weakly northward (green, 0° < 6\MF < 30°), and southward (red, #IMF < 0°) cases. In the right panels (b), the Geotail data between 0 UT and 8 UT of April 4 are shown (from top, |£?|, ion temperature (T{), ion density (N{), Vx, A and energy-time diagram of omnidirectional electrons and ions). After a northward turning of the IMF at ~ 19 UT of April 2, the IMF kept its northward direction for more than 30 hours (T]ag for this event is ~ 30 hours). The solar wind velocity and density were normal (300-350 km/s and 4-8 /cc, respectively) during this interval. Geotail came into
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Fig. 1. \MF-Bz of each Cold-PS event (convection delay included), Left (right) panels are for one-(two-) temperature events. The hatched regions correspond to intervals of T] ag
the Cold-PS at 0030 UT of April 4 (at the last of this prolonged northward IMF interval), when it was located at (—11.0,19.5,9.6) RE- In this Cold-PS, the ion distribution was one-temperature, and the ion temperature was extremely low (~0.3-0.4 keV) in comparison with the usual Cold-PS (~0.5-l keV). After 0330 UT, on the other hand, the hotter and tenuous plasmas (Ts ~l-2 keV, Nt ~0.2-0.4 /cc) are observed in the plasma sheet, while the IMF were still (weakly) northward. At the transition from Cold-PS to hotter plasma sheet, plasma flows in the plasma sheet were quite stagnant (< 100 km/s) and no substorm features were observed. Geotail was located at (—12.3, 18A,6.9)RE at 0330 UT. This means that the cold plasma could not enter into this region during this prolonged northward IMF interval, and that the cold plasma was restricted to near the magnetopause. (Our statistical survey (not shown here) of cold plasma sheet shows that there seems to be a region where cold plasma is never observed even after a prolonged northward IMF interval. This might be related to transport pass of the cold plasma, which is beyond the scope of this report.) Next we show an event in which the cold ions changed from two-temperature to one-temperature after a prolonged northward IMF interval. Figure 3 shows (a) the solar wind, (b) the magnetosheath, and (c) the magnetosphere observations by Wind and Geotail on February 16-17, 1995. At the Wind location (X = 198 RE), the IMF was weakly northward (~ 1 nT) after 1730 UT of February 16 (Hereafter, convection delay (~ 50 minutes) from Wind to Geotail location is included). However, Geotail in the duskside magnetosheath observed a few southward excursions of the IMF; between 2115 UT and 2250 UT (95 minutes), between 2311 UT and 2322 UT (11 minutes), and between 2356 UT and 2400 UT (5 minutes). Therefore we define 2250 UT as the beginning of the northward IMF interval at the Geotail location. After the northward turning of the IMF around 2250 UT, Geotail observed the magnetosheath and the LLBL until 0410 UT of -30-
Fig. 2. An example of Cold-PS with one-temperature ions on the duskside. (a) WIND observation (56-hour plot) during 0 UT of April 2, 1995 and 8 UT of April 4 (convection delay included) and (b) GEOTAIL observation during 0-8 UT of April 4, 1995 are shown. The IMF polarity is color-coded with 3 cases; that is, strongly northward (0IMF > 30°), weakly northward (0° < ftMF < 30°), and southward intervals.
February 17, 1995. Geotail came into the cold plasma sheet with two-temperature ions around 0410 UT. when Geotail was located at (—34.4,23.9,4.6) RE- This two-temperature state continued until around 07 UT. During this two-temperature ion interval, the ion temperature was 0.6 - 0.9 keV, the density was 0.5 0.9 /cc, and the intensity of the magnetic field was 3-12 nT. The ion counts had two peaks around 7 keV and 0.7 keV, and this two-temperature state continued until around 07 UT. Between 0540 UT and 0600 UT Geotail repeated multiple crossings of the neutral sheet with a period of ~2-3 minutes, which were due to the flapping motion of the plasma sheet related to substorm activities that is thought to be triggered by the southward IMF between 0520 UT and 0620 UT. An intensification of the westward electrojet was observed by the CANOPUS magnetometers around 06 UT, although no fast flow was observed at the Geotail location. Around 07 UT, the higher part of ions began to decrease, and the one-temperature Cold-PS appeared. This one-temperature state continued until 09 UT. During this one-temperature interval, Geotail stayed near the central plasma sheet except for a brief excursion to the tail lobe around 0710 UT. In this one-temperature Cold-PS, the ion temperature and density were 0.5-0.8 keV and 0.4-0.7 /cc, respectively, which were slightly colder than that of the preceding two-temperature Cold-PS. The peak in the counts of this one-temperature ions was around 0.7-0.8 keV that is as low as that of the lower-energy part of the previous two-temperature ions. The plasma flow was stagnant (< 100 km/s) throughout the interval of Cold-PS observation between 0410 UT and 0900 UT. In particular, no fast flow was observed when the transition from two-temperature to one-temperature occurred. The IMF turned to be southward around 0810 UT, and thus hot and tenuous plasma sheet emerged and fast tailward flows were observed after 09 UT. In this event, the duration from -31-
Fig. 3. Wind and Geotail observations on February 16-17, 1995. (a) The solar wind observations (including a convection delay ~ 50 minutes), (b) Bz in the magnetosheath observed by Geotail, and (c) Geotail observations between 04 UT and 10 UT of February 17.
the northward turning of the IMF to the emergence of one-temperature Cold-PS was ~ 8.2 hours, which is longer than those of 2-temperature events (Table 1).
DISCUSSION There is a dawn-dusk asymmetry in the distribution function of cold ions; that is, two-temperature ions are often observed in the dusk flank plasma sheet and never in the dawnside (Takashima, 2002). He mentioned that both lower- and higher-energy parts of two-temperature ions are of magnetosheath origin, and that the higher-energy part is created by magnetic reconnection during the northward IMF intervals. However, our result shows that the higher-energy part is of magnetospheric origin, since the ion distribution function in the duskside Cold-PS changes from two-temperature to one-temperature under the prolonged northward IMF condition. If higher energy part of ions were of magnetosheath origin (or was created by magnetic reconnection during northward IMF intervals), the ion distribution should continue to be two-temperature after a prolonged northward IMF intervals. Therefore, the dawn-dusk asymmetry of the higher energy part may result from ion drift motion in the near-earth plasma sheet during southward IMF periods. We propose a story of the evolution of the duskside near-earth plasma sheet during northward IMF periods: With substorm activities the hot-tenuous plasma (HTP) dominates the plasma sheet. After the northward turning of the IMF the cold ions enter into the duskside near-earth plasma sheet so that Cold-PS with two-temperature ions is established. If the northward IMF continues for more than several hours, higher energy ions disappear and Cold-PS with one-temperature ions is formed. On the dawnside, on the other hand, only HTP and Cold-PS with one-temperature ions are formed.
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Fujimoto et al. (2002) noted that the dawnside Cold-PS does not seem to require steady northward IMF, and that the presence of the cold ions along the plasma sheet dusk flank is less apparent during the southward IMF. This dawn-dusk asymmetry remains to be solved in future works. In this report we have studied the emergence of Cold-PS from the viewpoint of the length of northward IMF durations. However, there is a possibility that not only positive IMF-_Bz but also the density in the solar wind (Afew) also affects the formation of Cold-PS. In the event on April 1, 1996 (Event No.7 in Table 1) when iVsw was larger than 20 /cc, Cold-PS with two-temperature ions was found only 2.6 hours after the northward turning of the IMF. The effect of solar wind density on formation of Cold-PS also need to be investigated. ACKNOWLEDGEMENTS We thank T. Mukai for use of GEOTAIL LEP data. We are grateful to S. Kokubun and T. Nagai for use of GEOTAIL MGF data. We also thank K. W. Ogilvie and R. P. Lepping for providing us with key parameter data from the SWE and MFI instruments on WIND, respectively. Canadian Space Agency is acknowledged for use of CANOPUS magnetometer data. This research is partially supported by ACTJST (Research and Development for Applying Advanced Computational Science and Technology of Japan Science and Technology Corporation). REFERENCES Borovsky, J. E., M. F. Thomsen, and D. J. McComas, The superdense plasma sheet: Plasmasheric origin, solar wind origin, or ionospheric origin?, J. Geophys. Res., 102, 22,089-22,097, 1997. Baumjohann, W., G. Paschmann, and C. A. Cattell, Average plasma properties in the central plasma sheet, J. Geophys. Res., 94, 6597-6606, 1989. Fairfield, D. H., R. P. Lepping, E. W. Hones, Jr., S. J. Bame, and J. R. Asbridge, Simultaneous measurements of magnetotail dynamics by IMP spacecraft, J. Geophys. Res., , 86, 1396-1414, 1981. Fujimoto, M., T. Terasawa, T. Mukai, Y. Saito, T. Yamamoto, and S. Kokubun, Plasma entry from the flanks of the near-Earth magnetotail: Geotail observations, J. Geophys. Res., 103, 4391-4408, 1998. Fujimoto, M., T. Tonooka, and T. Mukai, Vortex-like fluctuations in the magnetotail flanks and their possible roles in plasma transport, in Geophys. Monogr., 133, edited by P. T. Newell, and T. Onsager, AGU. Washington, D. C , pp. 241-251, 2002. Hill, T. W., Origin of the plasma sheet, Rev. Geophys. Space Set, 12, 379-388, 1974. Lennartsson, W., and E. G. Shelley, Survey of 0.1- to 16- keV/e plasma sheet ion composition, J. Geophys. Res., 91, 3061-3076, 1986. Lennartsson, W., A scenario for solar wind penetration of earth's magnetotail based on ion composition data from the ISEE 1 spacecraft, J. Geophys. Res., 97, 19,221-19,238, 1992. Nishino, M. N., T. Terasawa, and M. Hoshino, Increase of the tail plasma content during the northward interplanetary magnetic field intervals: Case studies, J. Geophys. Res., doi:10.1029/2002JA009268, 2002. Takashima, S., Two component of plasmas in the low-latitude boundary layer, M. S. Thesis, Department of Earth and Planetary Sciences, Tokyo Institute of Technology, 2002. Terasawa, T., M. Fujimoto, T. Mukai, I. Shinohara, Y. Saito, T. Yamamoto, S. Machida, S. Kokubun, A. J. Lazarus, J. T. Steinberg, and R. P. Lepping, Solar wind control of density and temperature in the near-Earth plasma sheet: Wind/Geotail collaboration, Geophys. Res. Lett, 24, 935-938, 1997.
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THE DAWN-DUSK ASYMMETRY IN MAGNETOSHEATH AND THE LEAKAGE OF ENERGETIC ELECTRONS: THE GEOTAIL OBSERVATION S. Imada1. M. Hoshino1, and T. Mukai2 1
Department of Earth and Planetary Science, University of Tokyo, Hongo 7-3-1 Bunkyo-ku Tokyo, Japan ^Institute of Space and Astronautical Science, Yoshinodai 3-1-1 Sagamihara Kanagawa, Japan ABSTRACT
We statistically and systematically studied the dawn-dusk magnetosheath asymmetry by using both thermal/middle energy electrons (< 40 keV: LEP) and energetic electrons (> 38 keV: EPIC) experiments onboard the Geotail spacecraft. We found the dawn-dusk magnetosheath asymmetry of energetic electrons distribution (> 38 keV), while the clear asymmetry could not be observed, for the lower energy electrons. It was also found that energetic electrons flux did not strongly change across the magnetopause as compared with that of the thermal/middle energy electrons in the dawn side magnetopause, and that the intensity of the energetic electrons increases with going towards the dawn side magnetopause from the bow shock. Those results suggest that energetic electrons leak out of magnetosphere into magnetosheath.
INTRODUCTION In the magnetospheric physics, several satellite observations of energetic particles have been carried out over the several decades [e.g., Sarris et al., 1976, Sarafopoulos et al., 2001]. It is considered that these energetic particles in the plasma sheet can be accelerated by the magnetic reconnection which converts the magnetic field energy to the kinetic energy. Many observations support the idea that energetic electrons bursts can be related to the magnetic reconnection and the formation of a neutral line [e.g. Terasawa and Nishida, 1976; Mobius et al., 1983]. The spatial distribution of energetic particles in the plasma sheet was also discussed by many peoples [e.g. Sarris et al., 1976; Sarafopoulos et al., 2001]. The intensity of energetic electrons (ions) flux in the dawn side (dusk side) is higher than that in the dusk side (dawn side). They also discussed a correlation between the geomagnetic activity and the energetic particles burst events. Sarafopoulos et al. [2001] studied energetic particle spectrum (from 25 to 850 keV) using the Interball tail probe. They concluded that the dawn-dusk asymmetry is mainly caused by the betatron acceleration mechanism due to the global convection motion in magnetotail. Energetic particles in magnetosheath were also discussed by many people [e.g. Paschalidis et al., 1994: Sarafopoulos et al., 2000]. Paschalidis et al. [1994] observed a dawn-dusk local time asymmetry in the magnetosheath ion distribution at energies > 50 keV, and they also found the magnetosheath ion spectra were harder in the dusk side than in the dawn side. Sarafopoulos et al. [2000] discussed the leakage of energetic particles from magnetosphere to magnetosheath, and they reported energetic electrons leak out from the equatorial dawn LLBL and protons leak out from the dusk. In this work, we statistically and systematically study the dawn-dusk magnetosheath asymmetry and discuss the energy dependence of the asymmetry. The IMF Bz dependence is also discussed from the magnetopause crossing event. These results may give important information concerning on the leakage and -34-
Fig. 1. The dawn-dusk magnetosheath asymmetry of electrons > 38 keV (top left), 43 keV ions (top right), 9.3 keV electrons (bottom left), 3.2 keV electrons (bottom right). The vertical axis is YQSM (RE) and the horizontal axis is XQSM (RE)- The color contour is averaged electrons/ions flux.
mixing of the magnetospheric and solar wind plasmas through magnetopause. DATA ANALYSIS Data Set We study the dawn-dusk magnetosheath asymmetry using the Geotail data from January 1995 to June 1997. We used electrons differential flux (the energy: 3.2 keV and 9.3 keV), ion differential flux (the energy: 43 keV) and ion moments (the ion density TV,, temperature T,, and the ion bulk velocity Vi) obtained from the Low Energy Particle (LEP) instrument with 1-min resolution. The coordinate is in GSM system hereinafter. For higher energy electron, we used the integrated electron flux (the energy: > 38 keV and > 110 keV) measured by the Energetic Particles and Ion Composition (EPIC) instrument with 1-min resolution. The particles flux was integrated over pitch angles by assuming that the high energy electrons are isotropic. We also used 1-min averages of the vector magnetic field data (Bx,By.Bz). The magnetosheath/solar wind data were selected by TV, > 3 /cc, Tt < 0.2 keV and Vx < -250 km/sec. The Dawn-Dusk Magnetosheath Asymmetry Figure 1 shows the statistical results of the dawn-dusk magnetosheath asymmetry of energetic electrons (> 38 keV) flux (top left), 43 keV ions flux (top right), 9.3 keV electrons differential flux (bottom left) and 3.2 keV electrons differential flux (bottom right). In this result, the IMF dependence is not taken into account. We divided this rectangular region into 11 x 11 bins each having a 5 RE X 4 RR. The color intensity shows the average electron and ion flux (top-left: /cm2/sec/str, others: /cm2/sec/str/keV). We find that energetic electrons flux of the dawn side magnetosheath is higher than the dusk side one, and that the intensity of energetic electrons flux near the magnetopause is larger than that around the bow shock. For the lower energy electron, however, such an asymmetry is not clear. The intensity of 43 keV ions is higher in the dusk side than in the dawn side, which shows the opposite asymmetry of the energetic -35-
Fig. 2. The dawn side magnetopause crossing event on November 29th, 1995. From the top, electrons flux 3.2(keV), 9.3(keV), >38(keV) and >110(keV), the magnetic field Bz in the unit of nT, the ion density Nt in cm~3, the ion temperature T; in keV and the ion bulk velocity Vx in km/sec.
electrons case. These results suggest that the presence of the energetic particles in the magnetosheath adjacent to magnetopause is strongly correlated with that in the plasma sheet, because in the plasma sheet the energetic electrons and ions are rich in the dawn side and in the dusk side, respectively [Imada et al.. 2002]. Dawn Side Observations on November 29, 1995 To understand the leakage process of the energetic electrons in detail, we studied the dawn side magnetopause crossing event. Figure 2 shows the event observed at (X, Y, Z) = (-21.7, -20.4, -0.2) B.R. From the top panel the thermal/middle energy and the energetic electrons flux, the z-component of magnetic field, the ion density, the ion temperature and the ion x-component bulk velocity are depicted. The hot and tenuous plasma was observed from 15:00 to 16:00 UT, which indicates that the Geotail was situated in the plasma sheet. From 16:30 to 17:15 UT, the spacecraft crossed magnetopause several times. Comparing the flux of energetic electrons in the plasma sheet with one in magnetosheath, one can find that the flux was almost constant, while lower energy ones dramatically changed across magnetopause. From this event study, we think that the lower energy electrons are well confined in magnetosphere, but the energetic electrons are likely to leak out across magnetopause. DISCUSSION and CONCLUSIONS We conducted the statistical analysis of the dawn-dusk magnetosheath asymmetry of electrons distribution that depends on the particle's energy. This result may be interpreted as the leakage of the plasma sheet electrons into the adjacent magnetosheath. We also studied the dawn side magnetopause crossing event to understand the leakage processes of energetic electrons in detail. We found that the flux of the energetic electrons across the dawn side magnetopause did not strongly change compared with the lower energy electrons. We think energetic electrons effectively leak through magnetopause. To explain the behavior of these energetic electrons in magnetosheath, we propose two mechanisms: One is the acceleration at the bow shock, and the other is the leakage from magnetosphere. Although one may think that the Earth's bow shock produces the dawn-dusk asymmetry by the different acceleration process in the parallel or the perpendicular shock, we found from Figure 1 that electron flux intensity is not strong near the bow shock but near the magnetopause. We conclude that the energy dependence of the dawndusk asymmetry is controlled by a local magnetopause process. When we think about the leakage from magnetosphere, two routes can be proposed: One is from the flank of the magnetotail, and the other is from -36-
Fig. 3. The relationship between energetic electrons > 38 keV integral flux and Bz of the magnetosheath on November 29th, 1995. The correlation can be observed between flux and Bz.
the dayside subsolar region. If these energetic electrons flew from subsolar region, the dawn-dusk asymmetry cannot appear. The dawn-dusk magnetosheath asymmetry reflects the plasma sheet conditions. It is proper to say that these electrons leak from the magnetotail flank. Figure 3 shows the relationship between energetic electrons flux of > 38 keV and local magnetic field Bz of the magnetosheath on November 29, 1995. We can see a good correlation between the flux and Bz. and it suggests that the northward Bz is in favor of the effective leakage of energetic electrons to magnetosheath thorough magnetopause. Hence the local IMF Bz seems to control the leakage process of energetic electrons through magnetopause. We are planning to reveal the relationship between the local IMF condition and the mechanism of the energetic particles leakage process through magnetopause. ACKNOWLEDGEMENTS We are grateful to all members of the Geotail team, especially to T. Terasawa, K. Maezawa, Y. Saito. The authors also thank R. W. McEntire and T. Hori for providing us the energetic electron data of EPIC/ICS. REFERENCES Imada. S.. M. Hoshino, and T. Mukai, The Dawn-Dusk Asymmetry of Energetic and Thermal Electrons: The Geotail Observation, in Proc. Sixth International Conference on Substroms, 388-393, 2002. Mobius, E., M. Scholer, D. Hovestadt, and G. Paschmann, Energetic particles in the vicinity of a possible neutral line in the plasma sheet, J. Geophys. Res., 88, 7742-7752, 1983. Paschalidis, N. P., E. T. Sarris, S. M. Krimigis, R. W. McEntire, M. D. Levine, I. A. Daglis and G.C. Anagnostopoulos, Energetic ion distributions on both sides of the Earth's magnetopause, J. Geophys. Res., 99, 8687-8703, 2000. Sarafopoulos, D. V., M. A. Athanasiu, D. G. Sibeck, R. W. McEntire, E. T. Sarris, and S. Kokubun, Energetic proton and electron dispersion signatures in the nightside magnetosheath supporting their leakage out of the magnetopause, J. Geophys. Res., 105, 15,729-15739, 2000. Sarafopoulos, D. V., N. F. Sidiropoulos, E. T. Sarris, V. Lutsenko, and K. Kudela, The dawn-dusk plasma sheet asymmetry of energetic particles: An Interball perspective, J. Geophys. Res., 106, 13,053-13,065, 2001. Sarris, E. T., S. M. Krimigis, and T. P. Armstrong, Observation of magnetospheric bursts of high-energy protons and electrons at 35 RE with IMP7, J. Geophys. Res., 81, 2342-2355, 1976. Terasawa, T. and A. Nishida, Simultaneous observations of relativistic electron bursts and neutral-line signature in the magnetotail, Planet. Space Sci., 24, 855-866, 1976. -37-
ELECTRON INERTIA EFFECTS IN AN MHD-SCALE KELVIN-HELMHOLTZ VORTEX T. K. M. Nakamura1 and M. Fujimoto1 l
Dept. Earth Planet. Set., Tokyo Inst. Tech., Meguro, Tokyo 152-8551, JAPAN ABSTRACT
We study the electron inertial effects on an MHD-scale Kelvin-Helmholtz (K-H) vortex. An LLBL like situation, that is, an MHD/ion scale velocity shear layer collocated with a current layer, is set up and the evolution of an MHD/ion-scale K-H mode is followed. The magnetic field is assumed to be perpendicular to the flow and the simulation plane. For a duskside-like situation, it is shown that smaller vortices appear within the MHD-scale parent K-H vortex. For the reasonable initial condition we adopt here, these appear only when the electron inertia effects are turned on. The smaller vortices grow as they are entrained into the center of the larger vortex and eventually destroy the parent's vortex pattern. For a dawnside situation, in contrast, such a decay process is not observed. We will present detailed analyses revealing the nature of these more or less surprising results.
INTRODUCTION The Low-Latitude Boundary Layer (LLBL) is known for its mixed plasma feature (e.g., Sckopke et al., 1981). The LLBL ions are composed of the solar wind-like cold component and the more energetic magnetospheric component. The LLBL seems to thicken up with distance down the tail during extended northward IMF (Mitchell et al., 1987; Fujimoto et al., 1998). It is also reported that there is a dawn-dusk asymmetry in the tail-LLBL ion spectrum feature, which may reflect the dawn-dusk asymmetry of the mixing process (Fujimoto et al., 2002). The Kelvin-Helmholtz (KH) vortex has been considered as an important element for understanding the dynamics of the tail-flanks (e.g., Miura, 1984, 1987). Since the tail-flanks are the region where the ion mixing becomes most prominent, one is inclined to think that the KH vortex is helping to produce the plasma mixing. Because the KH vortex size is of MHD-scale, one may think that the behavior of the vortex can be well described by the MHD equations. The plasma mixing, however, does not take place by definition in ideal-MHD. This has led Fujimoto et al. (1994) to think of the finite ion Larmor radius effects, and Otto and Fairfield (2000) to think of reconnection embedded within the vortex caused by a non-ideal effect, for KH vortices to become viable as a plasma mixing machine. Wilber and Winglee (1995) studied a KH vortex by a full particle code, however, the size of the vortex was not fully up to the MHD-scale. In this paper, we will study the effects of another non-MHD term on an MHD-scale KH vortex. We will consider a velocity shear layer of a transverse geometry, collocated with a current layer. We use the equation that takes the finite mass of electrons into account. While the ion inertia (the Hall term) does not change the result from an MHD calculation, the electron inertial effects make an MHD-scale KH vortex to decay. This totally different result, which is quite interesting from the plasma mixing point of view, is available even when the thickness of the velocity shear layer is of MHD-scale, and only in duskside-like situations. Simulation model The basic equations we use are the same as the Hall-MHD equations except that the electron finite inertia is taken into account in the magnetic field induction equation (Biskamp, 2000). In the normalized form, -38-
Fig. 1. Density contours to show the decay of an MHD scale vortex when the electron inertia is taken into account for a duskside-like case. Smaller vortices appear and grow within the parent MHD scale to destroy the well-ordered pattern. The decay process is not available for a dawnside-like case.
In this paper, we show the results with the ion-to-electron mass ratio of Mi/Me=25. The electron temperature is assumed to be zero for simplicity, however, this will not affect the results significantly. Normalizations are made by quantities on the magnetospheric side. The ion inertial length and the inverse of the ion gyrofrequency are the unit for the spatial scale and time, respectively. The magnetic field is assumed to be perpendicular to the flow and the simulation plane. The stream-wise size of the simulation box is determined based on the wavelength of fundamental KH mode 15D, where D is the initial half-thickness of the ion velocity shear layer. While the results for D =2 are shown, the vortex decay process is found to proceed similarly when D is varied (Hayashi et al., 2003). The magnitude of ion velocity jump across the magnetopause is equal to the Alfven speed. The density and the magnetic field magnitude in the magnetosheath is 4 and 0.5, respectively. The magnetic field does not change sign across the magnetopause but does change in intensity to keep the pressure balance. The ion beta in the magnetosphere is 0.25 and the ion teperature is uniform. The fundamental KH mode is activated by initially adding small amplitude perturbation in a velocity component. To induce electron inertial effects, a small amplitude random magnetic perturbation is added.
Results The fundamental MHD-scale KH vortex is nicely rolled-up by T=80D and this state stays rather stable when the simulation is done by the MHD equations. The vortex behaves similarly when the ion inertial effects are included (Hall MHD). When the electron dynamics is taken into account, however, we observe the decay of the MHD-scale vortex for the duskside-like situation. Figure 1 shows the density contours from four cases, with/without electron inertia, and dusk/dawnside situations. Panels for the duskside case show that smaller vortices appear and are entrained into the fundamental MHD-scale vortex only when the electron inertia effects are included. Hereafter, they will be called A-vortices. As the A-vortices are entrained into the center of the parent vortex, they grow quickly in time and expand outward to destroy the well-ordered parent vortex pattern. In contrast, in the dawnside case, even with the electron finite mass effects, A-vortices don't show up and no decay process is observed. Here we note that further smaller vortices (hereafter called B-Vortices) appear at the edges of the highdensity surge intruding into the magnetospheric side for both cases only when the electron inertia effects are included. By varying the initial half-thickness D and the ion-electron mass ratio, we have found that the size of the A-vortex is proportional to D. That is, the decay process follows a similarity law (Hayashi et al., 2003). On the other hand, the size of the B-vortex is found to be controlled by the mass ratio as jjf-4. -39-
Fig. 2. Two quantities to feature the current sheet kink instability at the hyperbolic point. The instability due to the finite electron inertia is activated at the hyperbolic point where the current density is enhanced. The generated magnetic perturbations are conveyed by the electron flows in the directions of the red arrows.
We first show that B-vortex is due to the finite electron inertia, as can be suggested from its size scaling on the mass ratio. Figure 2 shows the intensity of the term which characterizes the electron inertial effects and the color contour for the current density. At the hyperbolic point where the flow is converging from top and bottom, the current sheet is pinched. The intense current density is subject to the current sheet kink instability (Suzuki et al., 2002). The instability is due to the electron inertia and produces magnetic perturbations whose wavelength is scaled to the mass ratio by j ^ - 4 , which is in agreement with the present results. We observe the magnetic perturbations to be carried by the electron flow, directed opposite to the electric current, in the direction of the red arrows in the Figure. For the duskside-like situation, these magnetic perturbations are carried to the velocity shear layer situated at the outer-edge of the parent vortex. Figure 3 shows the structure of this velocity shear region. The line plots show the Vx component along the black broken lines in the contour plots for density. The red (blue) trace shows the profile at T=0 (65D). For both dusk and dawnside cases, a secondary shear layer is generated at the outer edge of the non-linearly rolled-up vortex. The thickness of the secondary shear layer is about 1/4 of the initial shear layer width (Variation in the plasma beta changes the amount of the velocity difference across the secondary layer but its thickness is always 1/4 of the initial). When magnetic perturbations are convected into this region, it would become the seed for a secondary KH instability. From the ratio of the thickness of the shear layers, the secondary KH vortex size is expected to be 1/4 of the parent vortex. Indeed, for the duskside-like case, B-vortices are carried by the electron flow in this region to give onset to the secondary KH instability that produces A-vortices. This naturally explains why the size of the A-vortices
Fig. 3. The structure of the secondary shear layer located at the outer edge of the parent vortex. Its width is about 1/4 of the initial.
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Fig. 4. Summary of the decay process. For a duskside case, magnetic perturbations conveyed to the secondary shear layer seed the secondary KH instability. This produces the smaller vortices that destroy the parent KH vortex. For a dawnside case, there is no seeding to the secondary shear layer.
is always about 1/4 of the parent's when the initial thickness D is varied. In contrast to the duskside case, the B-vortices generated at the hyperbolic point is carried to the opposite direction in the dawnside case. Since the perturbations do not seed the potential instability in the secondary shear layer, A-vortices do not appear. One may think that the A-vortices are of MHD nature and the electron inertial effects work only to trigger their growth. Then the decay process should be obtained as long as the seed to the secondary KH instability is somehow added. We have confirmed this to be true by numerical experiments where perturbations are artificially added to the secondary shear layer at the right time. The growth of the Avortices and subsequent decay of the parent vortex are observed in an MHD simulation as well. The electron inertia effects, in contrast, naturally give rise to the seeding in the right place at the right time.
Summary and Discussion Figure 4 summarizes how the electron inertial effects lead to the self-similar decay of an MHD scale K-H vortex. The current sheet kink instability that produces perturbations of the hybrid-scale is activated at the hyperbolic point. For a duskside-like case, the perturbations are conveyed by the electron flow to the secondary shear layer at the outer edge of the parent vortex. The secondary KH vortices (A-Vortices) grow as they are entrained into the center of the parent vortex to destroy the rolling-up pattern. For a dawnsidelike case, on the other hand, the perturbations do not propagate to the secondary shear layer and thus the decay does not take place. The similarity of the decay process observed when the initial thickness D and/or the mass ratio are varied is because the secondary shear layer thickness and so the size of the A-vortices are always 1/4 of the initial and the parent's, respectively. This explains the puzzling nature of the decay process that requires finite electron inertia but follows the hydrodynamic scaling at the same time. Decay of a large vortex that has entrained plasma widely from both sides of the shear layer is an attracting way of accomplishing plasma mixing across the magnetopause. The decay process shows dawn-dusk asymmetry, which may be related to the dawn-dusk asymmetry of the mixed-ion feature observed in the tail-flanks (Fujimoto et al., 2002). In order to apply the present results to the magnetotail situation, however, there are items that must be cleared. The first among many is to ask if plasma mixing is trully realized by the flows after the decay, whose main components are at MHD scales. The next would be to see how inthe-plane magnetic component changes the results. Since the decay is due to the secondary KH instability, one can imagine that the decay process in its present form will be shut down when the component is strong enough to stabilize the secondary KH. Will there be new effects due to the electron inertia in this regime? Will reconnection of the severely stretched field lines within the parent vortex (Otto and Fairfield, 2000) be triggered by the electron inertia? Studies on these issues are being performed (Nakamura and Fujimoto, in preparation). -41-
Acknowledgements T. N. thanks I. Shinohara, H. Matsumoto, and M. Hoshino for their stimulating discussion. M.F. is supported by the Grant-in-Aid #13640446 from the Ministry of Education, Science, and Technology. M. F. is a member of the ACT-JST project 12D-1. REFERENCES Biskamp, D., Magnetic Reconnection in Plasmas, Cambridge Univ. Press, 2000. Fujimoto, M., and T. Terasawa, Anomalous ion mixing within an MHD scale K-H vortex, J. Geophys. Res., 99, 8601, 1994. Fujimoto, M., et al, Plasma entry from the flanks of the near-Earth magnetotail: Geotail observations J. Geophys. Res., 103, 4391, 1998. Fujimoto, M., T. Mukai, and S. Kokubun, Cold-dense plasma sheet and the hot-dense ions in the magnetotail, Adv. Space Res., 30 (10), 2279, 2002. Hayashi, D., et al., Decay of MHD-scale vortices by parasitic electron dynamics, submitted to Phys. Rev. Lett, 2003. Mitchell, D. G., et al., An extended study of the low-latitude boundary layer on the dawn and dusk flanks of the magnetosphere, J. Geophys. Res., 92, 7394, 1987. Miura, A., Anomalous transport by magnetohydrodynamic Kelvin-Helmholtz instabilities in the solar wind magnetosphere interactions, J. Geophys. Res., 89, 801, 1984. Miura, A., Simulation of KH instability at the magnetospheric boundary, J. Geophys. Res., #£3195, 1987. Otto, A., and D. H. Fairfield, Kelvin-Helmholtz instability at the magnetotail boundary: MHD simulation and comparison with Geotail observations, J. Geophys. Res., 105, 21,175, 2000. Sckopke, N. G., et al., Structure of the low-latitude boundary layer, J. Geophys. Res., 86, 2099, 1981. Suzuki, H., M. Fujimoto, and I. Shinohara, Current sheet kink instability at ion-electron hybrid scale, Adv. Space Res., 30 (12), 2663, 2002. Wilber,M., and R.M. Winglee, Dawn-dusk asymmetries in the low-latitude boundary layer arising from Kelvin-Helmholtz instability: A particle simulation, J. Geophys. Res., 100, 1883, 1995.
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ON ION PROPERTIES WITHIN THE SUBSOLAR MAGNETOPAUSE CURRENT LAYER UNDER THE NORTHWARD AND SOUTHWARD IMF M. Nowada1'2, T. Sakurai2, and T. Mukai1 'Institute of Space and Astronautical Science, Sagamihara, Japan Department of Aeronautics and Astronautics, Tokai University, Hiratsuka, Japan
2
ABSTRACT This paper describes the ion properties within the subsolar magnetopause current layer (MPCL) under the northward and southward IMF. This study was based on the magnetic field and plasma data measured by the GEOTAIL spacecraft on November 10, 1995 and January 3, 1996. An interesting signature of a "collimated" ion distribution was found under the southward IMF. This can be interpreted as evidence that the ions originating from the magnetosheath were distributed along the open field lines formed by the magnetic reconnection on the MPCL. Under the northward IMF the low energy ions were found to be coexisting (or mixing) with energetic magnetospheric ions on the outer edge of the magnetosphere. This result indicates that the LLBL existed under the northward IMF.
INTRODUCTION The structure and dynamics of the Earth's dayside magnetopause current layer (MPCL) are of great interest because they are an important factor in determining the transport of the mass, momentum, and energy of plasma from the magnetosheath into the magnetosphere. Using particle data measured by the ISEE satellite, Song et ah (1993) found that the MPCL had multiboundary structures. The plasma within each layer consisted of three different ion components, that is, the magnetosheath and magnetospheric ions, and the intermediate energy ions between the magnetosheath and the magnetosphere. No associated significant heating (or cooling) of the plasma due to the magnetic reconnection was observed within the layers of the MPCL. However, these plasma characteristics within the MPCL revealed by Song et al. (1993) were observed under the extremely northward IMF with a magnitude of approximately 40.0 nT. In the present study, ion properties within the dayside MPCL under the northward and southward IMF are examined using the ion energy - time spectrograms (E-t diagram) and two-dimensional ion distribution functions. The magnitude of the IMF Bz was often approximately 2.0 nT in both cases. OBSERVATIONS This study is based on the magnetic field (Kokubun et al, 1994) and low energy plasma moment (LEP) (Mukai et ah, 1994) data obtained when the GEOTAIL spacecraft was crossing the dayside MPCL on November 10, 1995 (hereafter, referred to as Case 1) and January 3, 1996 (hereafter, referred to as Case 2). The time resolutions of the magnetic field and the plasma moment data are 3.0 seconds and 12.0 seconds, respectively. In particular, the LEP instrument onboard GEOTAIL covers the energy range from 32.0 eV/Q to 39.0 keV/Q, which is a broader range than the energy coverage between 75.0 eV/Q and 20.0 keV/Q of the LEP instrument onboard previous satellites such as ISEE. The MPCL crossings are identified based on the plasma moment data measured with sufficient
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counting statistics by means of an LEP-EA (Energy-per-charge Analyzer) instrument. In this study, the MPCL is defined by a clear jump of the north-south magnetic field component (i.e. Bz component in GSM coordinate system) due to Chapman - Ferraro current. The MPCL crossing was observed between 084433 UT and 084442 UT on November 10, 1995 under the northward IMF. The MPCL crossing on January 3, 1996 was observed between 021330 UT and 021345 UT under the southward IMF. MPCL CROSSING UNDER THE NORTHWARD IMF ON NOVEMBER 10,1995: CASE 1 Case 1 is the inbound MPCL crossing which occurred under the northward IMF. Local time during this MPCL crossing was 121200 LT. Figure la shows the E-t diagrams for four directions (dawnward, duskward, sunward, and anti-sunward (tailward)). The horizontal and vertical axes provide the universal time (UT) and energy of the ion (keV), respectively. The MPCL is bracketed by broken lines. The ion population is also shown by the lower color scale with the range between 5.0 counts/sample and 500 counts/sample. In the magnetosheath, the ions in the sunward and duskward directions are distributed up to 7.0 keV while those in the dawnward and anti-sunward directions are distributed up to 4.0 keV. Within the MPCL, the dawnward and anti-sunward ion populations are also less than those in the sunward and duskward directions. The upper limit of the ion energy in the sunward and duskward directions is 3.0 keV higher than that in the dawnward and anti-sunward directions (1.0 keV). Although in the magnetosphere, the ions with energy higher than 10.0 keV are distributed in all directions, the lower energy ions between 0.1 keV and 0.7 keV are seen in the sunward and duskward directions, adjacent to the MPCL and coexisting (or mixing) with the high energy magnetospheric ions. In order to clarify the ion distribution property within the MPCL, we examine the ion distribution function and compare it with that in the magnetosheath. Figures lb and c show two-dimensional ion distribution functions projected onto an equatorial plane in the magnetosheath observed at 084401 UT and within the MPCL at 084438 UT, shown with the black and white dotted lines in Figure 1 a, respectively. The upper, lower, left and right sides correspond to the dawnward, duskward, sunward and antisunward directions, respectively. The color bar shows the logarithms of the phase space density with the unit of m" V 3 . The direction of the magnetic field is superimposed on the distribution functions with a thick arrow. The range of the distribution function is appropriately fixed so that the distribution of the low energy ions is particularly clear. The cold ions with the energy range between 200 eV and 400 eV shown by the dotted circles are distributed along the magnetic field lines within both the magnetosheath and the MPCL. No essential difference between the distributions can be seen.
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Fig. 1 (a) Ion energy - time spectrograms for four directions (the dawnward, sunward, duskward and anti-sunward directions) of the MPCL crossing during the interval from 084300 UT to 084600 UT on November 10, 1995. The population of the ion from 5.0 counts/sample to 500 counts/sample is shown with the lower color bar. The MPCL is bracketed by thick broken lines. (b),(c) two-dimensional ion distribution functions projected onto an equatorial plane in the magnetosheath (b) at 084401 UT and within the MPCL (c) at 084438 UT, shown with the black and white dotted lines in Fig. l(a). The upper, lower, right and left sides correspond to the dawnward, duskward, tail (anti-sun) ward and sunward directions, respectively. The color bar shows the logarithms of the phase space density with the unit of m"V3. MPCL CROSSING UNDER THE SOUTHWARD IMF ON JANUARY 3,1996: CASE 2 Case 2 is the inbound MPCL crossing which occurred under the southward IMF. Local time during the MPCL crossing was 111 800 LT. Figure 2a shows the E-t diagrams obtained during the MPCL crossing and the format is the same as that in Figure la. Although the ions in the sunward, duskward and anti-sunward directions in the magnetosheath are distributed up to 7.0 keV, those in the dawnward direction are only distributed up to 3.0 keV. Within the MPCL, the upper limit of the ion energy in the sunward, duskward and anti-sunward directions is 5.0 keV higher than that in the dawnward direction (2.0 keV). The ion populations in these three directions are almost the same, and are much higher than the ion population in the dawnward direction. In the magnetosphere, the energy of the dominant ions is higher than 10.0 keV. The low energy ions seen adjacent to the MPCL under the northward IMF are absent. Figures 2b and c show the two-dimensional distribution functions in the magnetosheath observed at 021258 UT, and within the MPCL at 021335 UT, shown with the black and white dotted lines in Figure 2a, respectively. The format is the same as that in Figures lb and c. The distribution of the cold ions as shown by the dotted circle in Figure 2b is shifted duskward and is along the magnetic field lines. However, within the MPCL, these ions have a "collimated" distribution along the magnetic field lines as shown with a dotted oval in Figure 2c.
Fig.2 The MPCL crossing data under the southward IMF observed on January 3, 1996. The formats of all figures are the same as those in Fig.l.
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SUMMARY AND DISCUSSION In this study, we examined the ion properties within the MPCL under the northward and southward IMF. Under the northward IMF, the ion distribution function is similar to that in the magnetosheath. This result suggests that the magnetic field lines in the magnetosheath were piled up on the MPCL, and the satellite observed the magnetosheath ions that were distributed along the magnetic field lines.
Fig.3 Plot of the intensity of the magnetic field in Case 1. The MPCL crossing is bracketed by solid lines. Figure 3 shows the plot of the magnetic field intensity in Case 1. The universal time is shown in the bottom of the figure. The MPCL crossing is bracketed by solid lines. From 084330 UT, the magnetic field intensity in the magnetosheath tends to increase closer to the MPCL. During the MPCL crossing, the intensity also increases abruptly up from 21.0 nT to 54.0 nT. The low energy ions were seen on the outer edge of the magnetosphere and coexisted (or mixed) with the energetic magnetospheric ions. This is considered as direct evidence that the LLBL existed in the outermost region of the magnetosphere under the northward IMF. This result supports the previous findings of Mitchell et al. (1987). However, under the southward IMF, the ions within the MPCL have a "collimated" distribution along the magnetic field lines. Finding of this characteristic ion distribution suggests that the ions originate from the magnetosheath and are distributed along the open field lines formed by the magnetic reconnection (Gosling et al., 1990, Fuselier <=* a/., 1991). The low energy ions on the outer edge of the magnetosphere as observed in Case 1 were not present. This result indicates that the LLBL was absent under the southward IMF. ACKNOWLEDGMENTS We thank the GEOTAIL mission team for providing both the magnetic and plasma moment data. REFERENCES Fuselier, S. A., D. M. Klumpar, and E. G. Shelly, Ion reflection and transmission during reconnection at the earth's subsolar magnetopause, Geophys. Res. Lett., 18 139 - 142, 1991. Gosling, J. T., M. F. Thomsen, S. J. Bame, T. G. Onsager, and C. T. Russell, The electron edge of the low latitude boundary layer during accelerated flow events, Geophys. Res. Lett., 17, 1833 - 1836, 1990. Kokubun, S., T. Yamamoto, M. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The Geotail magnetic field experiment, J. Geomag. Geoelect., 46, 7 - 2 1 , 1994. Mitchell, D. G., F. Kutchko, D. J. Williams, T. E. Eastman, L. A. Frank, and C. T. Russell, An extended study of the low-latitude boundary layer on the dawn and dusk flanks of the magnetosphere, J. Geophys. Res., 92, 7394 -7404, 1987.
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Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The Low Energy Particle (LEP) experiment onboard the GEOTAIL satellite,/ Geomag. Geoelect., 46, 669 - 692, 1994. Song, P., C. T. Russell, R. J. Fitzenreiter, J. T. Gosling, M. F. Thomsen, D. G. Mitchell, S. A. Fuselier, G. K. Parks, R. R. Anderson, and D. Hubert, Structure and properties of the subsolar magnetopause for northward interplanetary magnetic field: Multi-instrument particle observations, J. Geophys. Res., 98, 11,319 - 11,337, 1993. E-mail address of Motoharu Nowada: nowada(a)stp.isas.ac.jp
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INFLUENCE OF SOLAR WIND ON SOURCE OF RING CURRENT PLASMA M. Nose1, R. W. McEntire2, and S. P. Christon3 1
Data Analysis Center for Geomagnetism and Space Magnetism, Graduate School of Science, Kyoto University, Oiwake-cho, Kitashirakawa, Sakyo-ku, Kyoto 606-8502, Japan 2 Applied Physics Laboratory, Johns Hopkins University, 11100 Johns Hopkins Road, Laurel, MD 20723-6099, USA 3 Focused Analysis and Research, Columbia, MD 21044, USA
ABSTRACT In the present study we aim to derive empirical equations relating source plasma of the ring current to the solar wind. We used the energy spectra at energies of 9-135 keV obtained by the suprathermal ion composition spectrometer (STICS) sensor of the energetic particle and ion composition (EPIC) instrument on the Geotail spacecraft. The plasma parameters (i.e., number density and temperature) of H+, O+, and He+ were estimated by fitting the K-distribution function to the energy spectra in the region of a geocentric distance of 8.5-10.5 RE and magnetic local time of 2200-0200 hour. The results showed that the H+ number density in the plasma sheet correlated with the solar wind density, while the O+ and He+ number density had no correlation with the solar wind parameters. Thus the origin of H+ ions in the plasma sheet is thought to be the solar wind. O+ and He+ ions in the plasma sheet are expected to have different origin from the solar wind. It was also found that the temperature of H+, O+, and He+ has a good correlation with the solar wind velocity and that gradients of the derived empirical equations can be ordered by ion mass. This implies that ions are accelerated in a mass-dependent way.
INTRODUCTION From a viewpoint of the space weather forecast, a numerical simulation of the ring current has been conducted by a large number of researchers. In such a numerical simulation, one of the important elements is the source term, that is, the outer boundary condition of the plasma on the night side. The boundary condition of the plasma is thought to have a great influence on results of a numerical simulation. Kozyra et al. (1998) examined an effect of the plasma density by conducting numerical simulations with two different outer boundary conditions. In one simulation, the plasma density of the inner plasma sheet was fixed at its prestorm value, while in the second simulation, the density was changed as indicated by the observations at geosynchronous orbit. The results showed that the Dst index calculated in the second simulation was much closer to the observed Dst than that in the first simulation. Thomsen et al. (1998) insisted that the intensity of storm-time ring current (the minimum value of Dst) is determined by a combination of the plasma sheet density and the interplanetary electric field. Therefore we need to estimate reasonably the plasma parameters at the outer boundary (i.e., the plasma parameters in the near-Earth plasma sheet) for simulation studies of the ring current. It is expected that the solar wind controls the plasma parameters in the near-Earth plasma sheet (Terasawa et al., 1997; Borovsky et al., 1997, 1998). Thus, in the present study, we aim to derive empirical equations relating the plasma parameters in the near-Earth plasma sheet to the solar wind parameters, contributing to the space weather forecast. DATA SET We used the H+, O+, and He+ ion flux data obtained by the suprathermal ion composition spectrometer (STICS) sensor of the energetic particle and ion composition (EPIC) instrument on board the Geotail satellite
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(Williams et al., 1994). The EPIC/STICS instrument measures energetic (9-210 keV) ion fluxes in eight energy steps, but we focused on the first seven energy steps (i.e., 9-135 keV), because the ion flux at 210 keV was sometimes below one count level. We made use of 10-min averages of the omnidirectional differential flux. The data period used here is from January 1995 through December 1998. We restricted a surveying area by a geocentric distance (r) of 8.5-10.5 RE and a magnetic local time of 2200-0200 hour, because we are interested in the near-Earth region. The plasma sheet was identified by the criteria of the magnetic field, that is, Bx^(Bx2+By2)V2<20 nT and ^s=tan"'(5z/5^y)>15°, which are similar to the criteria used by Baumjohann et al. (1990). The magnetic field data were obtained by the magnetic field (MGF) experiment (Kokubun et al., 1994). In respect of the solar wind data, we used 1-hour averages that are available from the OMNIWeb database (http://nssdc.gsfc.nasa.gov/omniweb/). DERIVATION OF PLASMA PARAMETER We estimated plasma parameters such as density and temperature in the near-Earth plasma sheet by fitting the K-distribution function to the observed energy spectrum. The K-distribution function is the generalized Maxwellian distribution function and describes well the energy spectrum at high energies. For this distribution function, the differential flux J at energy E is given by the following form:
WJ
{2XKEJ
[r(K-\/2)){
KEJ
where n, m, Eo, F, K are, respectively, the number density, the ion mass, the characteristic energy, the Gamma function, and the spectral index. At the limit of K^-CO, the distribution becomes the Maxwellian distribution. The plasma temperature 7" is expressed by T^E,^—. K-1.5
In order to evaluate how well the fitting is done, we used a £ function given by
X2 = ^ ( l o g 1 0 Job.{E,) - log10 J^AE,))2,
(3)
where •/„/,., and Jm0M are the observed and model differential fluxes at the ith energy step E,, respectively. We need to find the best set of plasma parameters (», T, K) that give the minimum value of £ Ci^min)- This task was accomplished by using the gradient-search method (Bevington and Robinson, 1992). When the value of ^min is less than or equal to 0.05, we consider that the fitting has been done adequately and used the fitting parameters in the subsequent analysis. Figure 1 shows examples of fitting to the energy spectra observed at 0838:48-0848:32 UT on June 26, 1998 (DAY 98177). The observed energy spectra of H+, O+, and He+ ions are presented by diamonds, triangles, and squares, respectively. The results of fitting, which are shown by dotted lines, were found to be fairly reasonable. The minimum value of £ was much smaller than 0.05 for all ion species, that is, 0.0063 for H+, 0.0024 for O+, and 0.0063 for He+. The derived plasma parameters are («, T, K)=(0.81 cm"3, 10.5 keV, 5.3) for H+, 3
+
3
(0.24 cm" , 12.5 keV, 7.1) for O , and (0.0071 cm" , -49-
Fig. 1. Examples of fitting to the energy spectra observed at 0838:48-0848:32 UT on June 26, 1998 (DAY 98177). Dotted lines are results of fitting.
(2)
10.1 keV, 4.2) for He+. Fitting the K-distribution function to the energy spectra observed in the near-Earth plasma sheet, we obtained 283 sets of plasma parameters (n, T, K) for H+, 100 sets for O+, and 234 sets for He+. The distributions of ^min for these plasma parameters are displayed in Figure 2. Occurrence probability has a peak at 22min=0.010-0.015 for H+ and / min =0.005-0.010 for O+ and He+. More than 70% of events have / m i n ^ 0.025 for all ion species.
Fig. 2. Distributions of the minimum value of x2 for the fitting plasma parameters, (a) H*, (b) 0 + , and (c) He*.
It should be noted that energy spectra in the energy range of 9-135 keV were used in the present study. The plasma parameters were derived with the assumption that ion distribution below 9 keV follows the K-distribution. However, recent studies reported satellite observations of enhancement of cold component (< 1 keV) in ion distributions in the central part of the near-Earth plasma sheet (Fujimoto et al., 2002; Seki et al., 2003). Fujimoto et al. (2002) found a few events containing cold plasma from the 3-year Geotail data, while Seki et al. (2003) showed that the probability of cold ion detection in the near-Earth plasma sheet was 30-40 % when Geotail entered the shadow of the Earth. There is a big discrepancy in occurrence probability of the cold plasma between the above two studies, but in case of appearance of such cold plasma, the plasma density (temperature) we derived here could be underestimated (overestimated). ANALYSIS Relation Between Plasma Number Density and Solar Wind Parameters We calculated correlation coefficients between the plasma number density in the near-Earth plasma sheet («PS) and the solar wind parameters. One hour averages of solar wind parameters were compared with the plasma sheet number density that is included in the 1-hour interval; that is, we considered nearly simultaneous response of the plasma sheet to the solar wind. The solar wind parameters were chosen as the number density («sw), the velocity (Ksw), the z component of the IMF (Bz), and the>" component of the electric field (isr=-Fsw52). The results are shown in the top three lines of Table 1. We found a large correlation coefficient between the H+ number density in the near-Earth plasma sheet and the solar wind number density (correlation coefficient (c.c.) of 0.67). There were no correlations between the number density of O+ and He+ and any solar wind parameters. (The number density of
Plasma Sheet
Table 1. Correlation coefficients between plasma parameters in the near-Earth plasma sheet and the solar wind parameters
«psfH+l «PsfO+l «psfHe + l
7kfH+l 7>s[O + l
FpsfHel
«sw 0.67 0.19 0.22 -0.16 -0.29 -0.25
Solar Wind Bz Ksw -0.20 0.27 0.34 0.085 -0.031 0.078 -0.25 0.49 0.64 -0.11 0.58 -0.071
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Ev—Fsw5z -0.30 -0.30 -0.20 0.30 0.27 0.21
0 + ions was in the range of 0.018-0.236 cm"3 with the mean of 0.060 cm"3. The number density of He+ ions was in the range of 0.0015-0.0132 cm"3 with the mean of 0.0060 cm"3.) In Figure 3a we plotted data points of «PS[H+] and « s w to derive an empirical equation between them. The empirical equation was found to be «PS[H+](cm"3)=0.0276-«sw(cm"3)+0.422
(4)
by the least square method, which is indicated by a solid line.
Fig. 3. (a) Relation between the H + number density in the plasma sheet and the solar wind number density, (b) Relation between the H + temperature in the plasma sheet and the solar wind speed, (c) Same as Figure 3b, except for the O + ion. (d) Same as Figure 3b, except for the He + ions.
Relation Between Plasma Temperature and Solar Wind Parameters Correlation coefficients between the plasma temperature in the near-Earth plasma sheet and solar wind parameters are presented in the bottom three lines of Table 1. Here we also considered nearly simultaneous response of the plasma sheet temperature. There were good correlations between the plasma temperature and the solar wind velocity for all ion species (c.c.=0.49-0.64). We found no good correlations between the temperature and other solar wind parameters. Figures 3b-3d shows relation between the plasma temperature and the solar wind velocity. Solid lines are regression lines that were determined by the least square method. These regression lines are expressed by rPS[H+](keV)=0.0162FSw(km/s)-0.00193, rPS[O+](keV)=0.0368-FSw(km/s)-6.142, rPstHe+](keV)=0.0228Fsw(km/s)-0.881.
(5a) (5b) (5c)
Relation Between Spectral Index and Solar Wind Parameters We performed correlation study between K and the solar wind parameters, but there was no good correlation between them. Thus we examined statistical characteristics of K[H + ], K[O+], and /c[He+]. Results are shown in -51-
Figure 4. We found that almost all of K values are less than 10 for all ion species. Their mean values are 8.26 for H+, 6.58 for O+, and 4.81 for He+.
Fig. 4.
Distributions of the spectral index (K) for (a) H+, (b) O+, and (c) He*.
DISCUSSION Number Density in Plasma Sheet It was found that the H+ number density is correlated well with the solar wind density. This might be caused by (1) permeation of the solar wind plasma to the plasma sheet or (2) compression of the magnetosphere. If mechanism 2 is plausible, we should observe good correlations for all ion species, which is in contradiction to the observational results. Therefore we consider that the solar wind plasma is transported to the plasma sheet and contributes to «ps[H+]. No correlation between the number density of O+ and He+ and the solar wind density implies that origin of O+ and He+ ions is not the solar wind. In future study we will examine what controls the O+ and He+ density. Temperature in Plasma Sheet The temperature in the plasma sheet was correlated well with the solar wind velocity for all ion species. We consider that the kinetic energy of the solar wind is transferred to the plasma sheet plasma. We found that gradients of the empirical equations can be ordered by ion mass, that is, the largest gradient (0.0368) for O+, the next large gradient (0.0228) for He+, and the smallest gradient (0.0162) for H+, implying that ion heating or acceleration mechanism is mass-dependent. Spectral Index in Plasma Sheet There was no good correlation between the spectral index and the solar wind parameters. As can be seen in Figure 4, more than 65 % of spectral index are distributed within a narrow range around the mean value (6-10 for H+, 4-7 for O+, and 3-6 for H e ) . These results indicate that spectral shape at high-energy tail is rather steady. The spectral shape is expected to be far from that of the Maxwellian distribution since the values of the spectral index are small (<10). We noticed that the mean values of K are ordered as K[He+]
their results we can expect that a time lag at r=8.5-10.5 RE, where we surveyed, will range from 0 hour to 3-4 hours. From a statistical survey of Geotail observations at X—15 to -50 RE, Terasawa et al. (1997) reported that there is a good correlation (c.c.=0.66) between the 1-hour plasma sheet density and the 1-hour solar wind density (see their Figure 2b); both quantities are for the same intervals. They also intimated the similar good correlation for the solar wind density averaged over 9 hours prior to the plasma sheet observation (see their Figure 3b). Therefore our results seem to be reasonable in comparison with the previous studies. However, to examine the solar wind-magnetosphere coupling in more detail, we need to consider time lag between the solar wind and the plasma sheet. ACKNOWLEDGMENTS We thank D. J. Williams and S. R. Nylund for their help in analysis of the Geotail/EPIC/STICS data. The Geotail/MGF data were supplied by T. Nagai. We thank J. H. King and N. Papitashvili for providing the OMNI data. Thanks are due to S. Ohtani, K. Takahashi, and A. T. Y. Lui for their helpful comments. This work was partly supported by the Sasagawa Scientific Research Grant from The Japan Science Society. REFERENCES Baumjohann, W., G. Paschmann, and H. Ltihr, Characteristics of high-speed ion flows in the plasma sheet, J. Geophys. Res., 95, 3801-3809, 1990. Bevington, P. R., and D. K. Robinson, Data Reduction and Error Analysis for the Physical Sciences, second edition, pp. 153-156, McGraw-Hill, Boston, 1992. Borovsky, J. E., M. F. Thomsen, and D. J. McComas, The superdense plasma sheet: Plasmaspheric origin, solar wind origin, or ionospheric origin?, J. Geophys. Res., 102, 22,089-22,097, 1997. Borovsky, J. E., M. F. Thomsen, and R. C. Elphic, The driving of the plasma sheet by the solar wind, J. Geophys. Res., 103, 17,617-17,639, 1998. Fujimoto, M., T. Mukai, and S. Kokubun, Cold-dense plasma sheet and hot-dense ions in the inner magnetosphere, Adv. Space Res., 30(10), 2279-2288, 2002. Kokubun, S., T. Yamamoto, M. H. Acuna, et al., The GEOTAIL magnetic field experiment, J. Geomagn. Geoelectr., 46,7-21, 1994. Kozyra, J. U., V. K. Jordanova, J. E. Borovsky, et al., Effects of a high-density plasma sheet on ring current development during the November 2-6, 1993, magnetic storm, J. Geophys. Res., 103, 26,285-26,305, 1998. Seki, K., M. Hirahara, M. Hoshino, T. Terasawa, R. C. Elphic, Y. Saito, T. Mukai, H. Hayakawa, H. Kojima, and H. Matsumoto, Cold ions in the hot plasma sheet of Earth's magnetotail, Nature, 422, 589-592, 2003. Terasawa, T., M. Fujimoto, T. Mukai, et al., Solar wind control of density and temperature in the near-Earth plasma sheet: WIND/GEOTAIL collaboration, Geophys. Res. Lett, 24, 935-938, 1997. Thomsen, M. F., J. E. Borovsky, D. J. McComas, et al., Variability of the ring current source population, Geophys. Res. Lett., 25, 3481-3484, 1998. Williams, D. J., R. W McEntire, C. Schlemm II, et al., Geotail energetic particles and ion composition instrument, J. Geomagn. Geoelectr., 46, 39-57, 1994. E-mail address of M. Nose
[email protected]
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RELATIONSHIP BETWEEN PLASMA AND MAGNETIC FIELD PARAMETERS IN THE DISTANT PLASMA SHEET Oleg Troshichev Arctic and Antarctic Research Institute, St.Petersburg, Russia
ABSTRACT Measurements of the magnetic field and low-energy plasma by Geotail have been used to study mean and instantaneous characteristics of plasma and magnetic field in the distant tail X = -(79-200) RE under extremely quiet and weakly disturbed conditions. The analysis has been carried out separately for the tail lobes and the plasma sheet. A good consistency between variations of the plasma and magnetic parameters is observed only in the tail lobes, the correspondence being considered for appropriate components (Vx and B x , Vy and By, Vz and Bz). The distant plasma sheet seems to be constantly in regime of turbulence, i.e. such a regime when the magnetic field and plasma parameters are steadily subjected to fluctuations of different periods, the variations in magnetic field and plasma velocity being inconsistent in time. The mean values of V and B in the plasma sheet are not compatible with their instantaneous characteristics at all. The coefficient of diffusion across the plasma sheet, estimated from experimental data, is in harmony with predictions of theory for the plasma sheet with turbulence. The applicability of the frozen-in approximation to description of processes in the plasma sheet is discussed.
INTRODUCTION The first observations of magnetic field and plasma in the distant tail have been fulfilled on board ISEE 3 spacecraft (Siscoe et al, 1984; Slavin et al, 1987). They showed that magnetic field Bz component at distances X < -90 RE demonstrates a dependency on geomagnetic activity, being positive for quiet conditions and negative for storm periods. The combination of northward B z and slow tailward plasma velocity has been revealed for intervals with low activity, and Slavin et al. (1987) argued that these results suggested continued slow stretching of closed field lines during quiet periods due to quasi-viscous interaction in agreement with the model of Heikkila (1984). Heikkila (1988) has analyzed the magnetic data from ISEE 3 exactly in the cross-tail current sheet (neutral sheet) and came to the conclusion that the field lines are closed out to X = -220 RE with moderate values of plasma velocity at all levels of geomagnetic activity; plasmoids consistent with negative B z may be present but surrounded by closed field lines. The interpretation of Heikkila (1988) was criticized by Nishida et al. (1994), who noted that it implies that unrealistically large amounts of the magnetic flux are transported tailward. The most comprehensive data set permitting study of the structure of the distant magnetospheric tail has been obtained on the basis of the 12-s averages of magnetic field (Kokubun et al., 1994) and particle measurements (Mukai et al., 1994) produced by GEOTAIL spacecraft in period between October 1, 1993 and October 31, 1994 at X=-(100-200)RE. Yamamoto et al. (1994) showed statistically that the high plasma density current sheet exists with a northward magnetic field. The analysis of magnetic and electric field data from Geotail carried out by Nishida et al. (1994, 1995) for geomagnetically quiet periods has shown that there is a northward component in the distant neutral sheet while the drift motion of the plasma is directed antisunward. To eliminate dilemma of unrealistically large amounts of the magnetic flux taking away by tailward plasma flow Nishida et al. (1998) suggested model of the distant tail twisted under influence of the IMF azimuthal component |BY| > B z . In this model the IMF field lines convect parallel to the neutral sheet and across the tail in the twisted tail geometry at the same time as they moved tailward. When these open lines reach the plasma sheet they are reconnected; as a result the northward field lines moving tailward would be observed in the distant plasma sheet, but these field lines are connected with IMF, not with geomagnetic field. The distant tail was examined as a "single entity" in discussed above studies. However, the further statistical analysis {Troshichev et al, 1999) has demonstrated the essential difference in behavior of plasma and magnetic field near the neutral sheet and far from the neutral sheet. The following criteria have been taken by Troshichev et al. (1999) to identify the plasma regimes in the plasma sheet core and in the tail lobes: the tail lobes are regions in the distant tail with a large magnetic B x component (|BX| >5 nT) and low ion temperature (T, <300eV); the plasma sheet is region with low magnetic B x component (|BX| < 5 nT) and large ion temperature (T; >300eV). The same -54-
criteria have been used while examining the relationship between the plasma and magnetic field parameters in particular events (Troshichev et al, 2000). The present paper summarizes peculiarities of the plasma and magnetic field behavior in the distant plasma sheet and tail lobes. AVERAGED CHARACTERISTICS OF PLASMA AND MAGNETIC FIELD Only quiet periods of continuous observations with duration from 30 to 90 min have been examined in statistical analysis of Troshichev et al. (1999). All events were divided into two categories: "extremely quiet" intervals without any signatures of magnetic activity (the Kp index averaged over the chosen time intervals was less than 1+, with the maximum hourly averaged AE index less than 200 nT), and "weakly disturbed" periods with noticeable substorm activity in the auroral oval (1+
200 nT). As a rule, extremely quiet intervals were associated with steady northward or close to zero IMF B z components, whereas the weak disturbances are observed in association with southward or fluctuating IMF. The "empty" tail lobes with plasma density lower than 0.05 cm"3 were not examined. The plasma sheet boundary layer (PSBL) with combinations (|Bx| < 5nT, Tj <300eV) or (|Bx| >5 nT, Tj> 300eV) has been excluded from examination. Figure 1 demonstrates the fundamental difference in the mean ion temperature for |BX| >5 nT and |BX| <5 nT taken as the basis for the identification of the tail lobes and plasma sheet in the distant tail. One can see that mean plasma energy in the tail lobes does not exceed 100 eV, irrespective of whether quiet or weakly disturbed conditions prevail, while the mean plasma energy in the core of the plasma sheet is 5-15 times higher than in the lobes. Statistical features of plasma and magnetic field typical of the plasma sheet and tail lobes are presented in Figures 2-5.
Figure 1. Distribution of mean energy (ion temperature) in the distant tail by consistent B x used as the basis for definition of the tail lobes and plasma sheet for extremely quiet and weakly disturbed conditions (from [Troshichev et al, 1999]).
Figure 2 shows distribution of radial component of plasma velocity V x in both regions under extremely quiet and weakly disturbed conditions. The average X component of the plasma velocity in the entire distant tail is tailward (-Vx), and this fact has been noted by Slavin et al. (1987), Heikkila (1988), and Nishida et al. (1995). Figure 2 demonstrates that plasma in the lobes flows almost exclusively tailward with velocities in range 0-300 km/s irrespective of level of magnetic activity. In the plasma sheet the peak of the statistical distribution under the extremely quiet conditions also falls within the range -100
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generally close to 0 for large tailward Vx velocities, suggesting that the vertical component of the magnetic field then fluctuates near zero level. By contrast, for earthward velocities and small tailward velocities (V x > - 400 km/s) the average B z component in the plasma sheet is consistently positive, representing the predominance of the northward B z in these conditions.
Figure 2. Histograms of the V x component in tail lobes and plasma sheet under extremely quiet and weakly disturbed conditions (from [Troshichev et at, 1999]).
Figure 3. Average magnetic field B z component plotted by consistent B x for the tail lobes and plasma sheet {from [Troshichev etal., 1999]).
The magnetic pressure is maximum near the external boundaries of the tail lobes (up to 350 eV cm"3) and decreases on approaching the plasma sheet, with a minimum value of 20 eV cm'3 at the neutral sheet (Figure 5). The plasma pressure, being small and almost unvarying in the tail lobes, increases quickly within the plasma sheet
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and reaches a maximum at the neutral sheet. The total pressure in the plasma sheet tends to be constant, with increase of the plasma pressure being balanced by decrease of the magnetic pressure and vice versa. So, the total pressure in the tail lobes is determined by the magnetic pressure, but the total pressure in the plasma sheet is determined mainly by the plasma pressure. The plasma pressure in both the tail lobes and the plasma sheet is 1.5-2 times greater for extremely quiet conditions than for weakly disturbed conditions owing to the plasma density is 25 times higher under the quiet conditions.
Figure 4. Average magnetic field B z component in the plasma sheet plotted by consistent Vx.
10
15
Figure 5. The magnetic, plasma, and total pressure values in the plasma sheet and tail lobes plotted by consistent B x for extremely quiet and weakly disturbed conditions (from [Troshichev etai, 1999]).
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Thus, the statistical analysis of Troshichev et al. (1999) showed that the averaged characteristics of plasma velocity and magnetic field in the plasma sheet and tail lobes are rather similar under the quiet conditions: plasma in both regions flows tailward and the vertical component of magnetic field is northward on the average. A particularly remarkable feature of the tail lobes and plasma sheet is the ratio of the magnetic and plasma pressure values. The magnetic pressure in the tail lobes (100 - 350 eV cm-3) is much higher than the plasma pressure (10 - 50 eV cm-3). On the contrary, in the plasma sheet the plasma pressure exceeds the magnetic pressure (about 10 times near the neutral sheet, B x ~ 0). Differences in the plasma characteristics turn out to be more obvious while going from extremely quiet and to disturbed conditions. However, these changes concern mainly plasma sheet. Indeed, the plasma velocity and magnetic field in the tail lobes do not show noticeable changes with magnetic activity, as well as the magnetic pressure is almost insensitive to changes in magnetic activity. On the contrary, in the plasma sheet the plasma density is higher and the plasma energy is lower under extremely quiet conditions than under disturbed conditions. The most evident changes occur in the average vertical B z component in the plasma sheet, which is consistently northward under extremely quiet conditions, tends to be 0 under weakly disturbed conditions, and, according to Slavin et al. (1987), Heikkila (1988), and Nishida et al. (1995), is southward during disturbed periods. However, reverses in Bz from northward to southward in the distant plasma sheet are unambiguously associated with highspeed tailward plasma flows relevant to the passage of plasmoids {Hones et al., 1984; Moldwin and Huges, 1992). This implies that specific plasma and magnetic field features observed in the distant tail during disturbed periods are related to processes taking place in the near-Earth tail.
Figure 6. Magnetic field in tail lobes ( B x I, B z components and total field B T O T ) at distances from -100 RE to - 200R E .
The passage of plasmoids apart, it seems that processes going on in the distant tail, as such, are the same for extremely quiet and disturbed periods. Since these periods are consistent with periods of northward and southward IMF influence, it means that the distant tail is unaffected by direction of the IMF (northward or southward) and, therefore, the concept of Dungey (1961) is not applicable for the distant tail. In other words, the tail formation proceeds irrespective of interconnection between the interplanetary and geomagnetic fields. Figure 6 shows the manner in which B x , B z components in tail lobes and total field Btoiai change at distances from -100 RE to -200 RE.
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Taking into account the peculiarities of the spacecraft orbit the all data have been grouped in four ranges: X=-(95118)RE, X=-(134-136)RE, X=-(166-172)RE, and X=-(198-200)RE. Data for southern and northern parts of the tail have been united, and mean values of B z , |BX|, and Btolai for each range have been calculated. As Figure 6 shows, the total magnetic field in tail lobes decreases at distance from 100 RE to 200 RE by about 11%, similar to B x . At the same time the decrease of B z is so large as 30%. Thus, the main Bx component in the tail lobes decreases with distance three times slower than the vertical component Bz. It implies that reconnection between the southern and northern tail lobes tends to zero while moving tailward and the lobe magnetic field lines tend to stretch along the neutral sheet. Such structure of the tail magnetic field can be stable on condition that currents flowing across the plasma sheet (from dawn to dusk) envelop the northern and southern tail lobes, and these currents are generated constantly, irrespective of the IMF polarity. RELATIONSHIP BETWEEN MAGNETIC AND PLASMA PARAMETERS IN PARTICULAR EVENTS Although the average plasma sheet velocity and magnetic field characteristics in the plasma sheet look similar to those in the tail lobe, the situation principally reverses if we examine relationship between plasma velocity and magnetic field in particular events. As an example, Figure 7 shows the relationship between the absolute value of magnetic field tail and plasma velocity in tail lobes (upper panel) and plasma sheet (lower panel) in the distant tail (X=-184 RE) on December 09, 1993. One can see that plasma velocity in the tail lobe is connected with magnetic field by approximately linear dependence: the lesser velocity, the higher is magnetic field, and vice versa. The determined relationship between the plasma velocity and magnetic field is observed in almost all crossings of the tail lobes under extremely quiet conditions. In the plasma sheet, on the contrary, the absolute values of plasma velocity and magnetic field are not associated to one another, and quite different combinations of these quantities take place.
Figure 7. Relationship between the absolute values of magnetic field and plasma velocity in the tail lobes and plasma sheet in the distant tail (X= -184RE) on December 09, 1993.
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Figure 8. Changes of the corresponding plasma and magnetic field parameters observed on September 23, 1994 in the northern tail lobe and in on July 12, 1994 the southern plasma sheet.
It is significant that linear correlation between magnetic field and plasma velocity in the tail lobes is typical not only of absolute values of B and V, but also of the corresponding components plasma velocity and magnetic field (Vx/B x , VY/B Y , VZ/BZ). AS an example, Figure 8(a) shows changes of plasma and magnetic field parameters observed in the northern tail lobe in case of September 23, 1994. One can see that components of plasma velocity in the tail lobe change consistently with the corresponding components of magnetic field, the agreement being kept sometimes up to the small fluctuations. As a result, the relationships between the components Vx and B x , VY and BY, VZ and B z can be approximated by linear law with coefficient of correlation R >0 75. There is also a good correlation between plasma density and B x : the lower magnetic field (and larger -Vx), the higher is plasma density. By the contrast, in the plasma sheet the corresponding pairs of components show rather poor correlation, even if spacecraft measurements are taken in limits of solely northern or southern half of the plasma sheet (it excludes uncertainties associated with crossings of the neutral sheet). As an example, Figure 8(b) shows behavior of the corresponding plasma and magnetic field parameters in the southern plasma sheet in case of July 12, 1994. The representative nature of these examples is demonstrated by Figure 9 showing correlation between plasma and magnetic parameters calculated on the assumption of linear dependence between them (B^cb+fiV,). Results presented in Figure 9 establish that linear correlation (|R|>0.5) between the corresponding components of plasma velocity and magnetic field (and between the total values of V and B) is typical only of the lobe crossings. Statistical analysis also confirmed the negative correlation between the plasma density and magnetic field in the tail lobes (not shown in Fig.9). A distinguishing feature of the plasma sheet crossings is a lack of any consistency between variations of plasma parameters and magnetic field. The X component of magnetic field serves in analyses (Troshichev el al., 1999, 2000) as a measure of the distance from the neutral sheet (i.e. B x is equal to zero at the neutral sheet and reaches the maximum negative or positive, on the external border of lobe). Therefore, the negative linear correlation between -V x and | B X | values, typical of tail lobes, means that tailward velocity in the lobe increases while approaching the plasma sheet border. The tail lobe plasma is usually considered as formed while expanding from the plasma mantle. In such a case, the rise of the lobe plasma density going toward the plasma sheet border can be regarded as indication on the lobe plasma piling up in vicinity of the border, while the tailward plasma velocity continues to increase. It is widely believed that validity of frozen-in approximation suggests correlation between the orthogonal components of plasma velocity and magnetic field (i.e. between V x and BYz, VY and B x z , V z and BXY). Figure 10 shows typical example of behavior of the orthogonal plasma and magnetic components in the distant tail lobe and plasma sheet. One can see the absence of any determined relationship between the examined quantities even in the
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tail lobes, where the frozen-in approximation is regarded as undoubtedly valid. It implies that relationship between the orthogonal components of plasma velocity and magnetic field has no any relation to the idea of the frozen-in approximation. So, the high variability of the magnetic field and plasma parameters is typical of the entire distant tail, but correlation between changes in plasma velocity and magnetic field (V x / B x , Vy / By, V z / B z ) is observed only in the tail lobes being unusual in the plasma sheet.
Figure 9. Histograms of coefficients of correlation between the corresponding components of plasma velocity and magnetic field in the tail lobes and in the plasma sheet.
Figure 10. Behaviour of the orthogonal plasma and magnetic field parameters observed in the northern tail lobe and in the southern plasma sheet.
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PLASMA SHEET AND TURBULENCE A distinguishing feature of the plasma sheet crossings is lack of any consistency between variations of plasma velocity and magnetic field. Inconsistency is observed for any combinations of plasma velocity and magnetic field parameters. Variations of plasma density and velocity components are also unrelated each to other in the plasma sheet. It implies that magnetic field in the plasma sheet is generally disconnected from the plasma, and regime of turbulence seems to be realized in the distant plasma sheet. Investigations of turbulence in the plasma sheet were initiated with work of Angelopoulos et al. (1992, 1993), which demonstrated that the random velocities in the central plasma sheet can be much higher then the regular velocity. Theoretical model of plasma sheet has been put forward by Antonova and Ovchinnikov (1996, 1999). They suggest that existence of the inconsistent fluctuations of the plasma sheet velocity and magnetic field implies the existence of plasma transport within the plasma sheet, and that equilibrium structure of plasma sheet is supported by the medium-scale turbulence against the regular plasma drift from the tail lobes. Plasma transport can be described, in the quasi-diffusion approximation, by the expression = < n>< V> - DVn, where F is the particle flux, n is density, V is the velocity, D is the eddy diffusion coefficient. The first quantitative investigation of the plasma sheet turbulence was carried out by Borovsky et al. (1997, 1998) who derived the plasma sheet diffusion coefficient as D = VRMS2- TAUTO/2, where VRMs is mean value of random velocity, and TAUTO is time of auto-correlation for components of the plasma velocity. Borovsky (1998) could estimate the diffusion coefficient as of order of 2.6 105 km /s within the plasma sheet. Antonova and Ovchinnikov (1996) determined the diffusion coefficient DZz from theoretical examination as falling in the range 0.1 - 1.5-105 km2/s, depending on the level of geomagnetic activity. Expression of Borovsky (1998) applied to data from Geotail gave the value of DZz in range from 0.3- 10s to 1.2-105 km2/s (Troshichev et al., 2000). Thus, all three independent estimations of DZz ensure the agreed values of the diffusion coefficient in the plasma sheet. This result provides evidence on regime of turbulence in the plasma sheet. DISCUSSION AND CONCLUSION A good consistency between variations of the plasma and magnetic parameters is observed only in the tail lobes, the correspondence being kept only for appropriate components (V x and Bx, VY and B Y , V z and B z ). The distant plasma sheet seems to be constantly in regime of turbulence, i.e. such a regime when the magnetic field and plasma parameters are steadily subjected to fluctuations of different periods, the variations in magnetic field and plasma velocity being inconsistent in time. Resolution of measurements (the 12-s averages) does not make it possible to show cascade in full temporal scales. Nevertheless we can argue that inconsistency of the magnetic field and plasma velocity variations in the plasma sheet asserts for periods from 30 seconds to 30 minutes. Zelenyi et al. [1998] have shown that form of the spectra of magnetic field fluctuations in the distant tail region can be explained if the observed magnetic field fluctuations have a fine fractal structure. Different auto-correlation times for the magnetic field and velocity fluctuations {Troshichev et al, 2000) imply the non-MHD character of observed turbulence. So, the inapplicability of simple MHD approach seems quite natural. The absence of any correlation between the observed fluctuations of the velocity and magnetic field in the plasma sheet is difficult to combine with an often-used suggestion that the frozen-in approximation is valid for the entire plasma sheet. The development of different plasma instabilities can lead the turbulization of the plasma motion implying the inadequacy of the frozen-in approximation for description of the plasma sheet processes. In particular, the problem of the northward magnetic field associated with tailward flow in the distant tail being unresolved in the framework of reconnection model of the plasma sheet formation, is simply eliminated if magnetic field be decoupled to the plasma velocity. In such a case the large-scale plasma flows take place in the conditions of the nonconservation of the magnetic flux through the selected plasma contour, and the background northward field, available in the plasma sheets can have no relation to the mean antisunward plasma velocity observed there. ACKNOWLEDGMENTS Author is very grateful to Atsuhiro Nishida, Toshifumi Mukai, Susumu Kokubun, and Yohsuke Kamide who provide him with GEOTAIL data and help by invaluable advises in analyzing these data. I am also thankful to Elisabeth Antonova for the fruitful discussions and comments. REFERENCES Angelopoulos, V., Baumjohann W., Kennel C.F., Coroniti F.V, Pellat R, et al., Bursty bulk flows in the inner central plasma sheet. J. Geophys. Res., 97, 4027 (1992).
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Angelopoulos, V., C.F.Kennel, F.V.Coroniti, R.Pellat, H.E.Spence, et al., Characteristics of ion flow in the quiet state of the inner plasma sheet. Geophys. Res. Lett., 20, 1711 (1993). Antonova, E.E., I.L.Ovchinnikov, The equilibrium of turbulent current sheet and the current sheet of the Earth's magnetotail, Geomagn. Aeronomy (in Russian), 38, 7 (1996). Antonova, E.E. and I.L.Ovchinnikov, Magnetostatically equilibrated plasma sheet with developed medium-scale turbulence: Structure and implications for substorm dynamics, J. Geophys. Res., 104, 17289 (1999). Borovsky, J.E., R.C.Elphic, H.O.Funsten, and M.F.Thomsen, The Earth's plasma sheet as a laboratory for flow turbulence in high p MHD, J. Plasma Phys., 57, 1 (1997). Borovsky, J.E., M.F.Thomsen, and R.C.Elphic, The driving of the plasma sheet by the solar wind, J. Geophys. Res., 103, 17617(1998). Dungey, J.W., Interplanetary magnetic field and the auroral zones, Phys. Rev. Lett., 6, 47, 1961. Heikkila, W.J., Magnetospheric topology of fields and currents, in Magnetospheric currents, Geophys. Monogr. Sen, 28, 208, (1984). Heikkila, W.J., Current sheet crossings in the distant tail, Geophys. Res. Lett., 15, 299 (1988). Hones, E.W., J.Bim, D.N.Baker, S.J.Bame, W.C.Feldman et al, Detailed examination of a plasmoid in the distant magnetotail with ISEE 3, Geophys. Res. Lett., 11, 1046 (1984). Kokubun, S., T.Yamamoto, M.H.Akuna, K.Hayashi, K.Shiokawa, and H.Kawano, The Geotail magnetic field experiment, J. Geomagn. Geoelectr., 46, 7 (1994). Moldwin, MB., and W.J.Huges, On the formation and evolution of plasmoids: A survey of ISEE 3 geotail data, J. Geophys. Res., 97, 19259 (1992). Mukai, T., S. Machida, Y.Saito, M.Hirahara, T.Terasawa, et al., The low energy particle (LEP) experiment on board the Geotail satellite, J. Geomagn. Geoelectr., 46, 59 (1994). Nishida, A., T.Yamamoto, K.Tsuruda, H.Hayakawa, A.Matsuoka, et al., Structure of the neutral sheet in the distant tail (X=-20 RE) in geomagnetically quiet times, Geophys. Res. Lett., 21, 2951 (1994). Nishida, A., T.Mukai, T.Yamamoto, Y.Saito, S.Kokubun and K.Maezawa, Geotail observations of magnetosphere convection in the distant tail at 200 RE in quiet times, J. Geophys. Res., 100, 23663 (1995). Nishida, A., T.Mukai, T.Yamamoto, S.Kokubun, and K.Maezawa, A unified model of the magnetotail convection in geomagnetically quiet and active times, J. Geophys. Res., 103, 4409 (1998). Siscoe, D.F., D.E.Jones, G D.Sibeck, J.A.Slavin, E.J.Smith, et al., ISEE 3 magnetic field observations in the magnetotail: implications for reconnection, in Magnetic Reconnection in Space and Laboratory Plasmas, Geophys. Monogr. Ser., 30, 240 (1984). Slavin, J.A., P.W.Daly, E.J.Smith, T.R.Sanderson, K.Pwenzel, et al., Magnetic configurationof the distant plasma sheet: ISEE 3 observations, in Magnetotail Physics, edited by A.T.Y.Lui, p.59, Johns Hopkins Uni. Press, Baltimore, Md. (1987). Troshichev, O.A., S.Kokubun, Y.Kamide, A. Nishida, T.Mukai, and T.Yamamoto, Convection in the distant magnetotail under extremely quiet and weakly disturbed conditions, J. Geophys. Res., 104, 10249 (1999). Troshichev, O.A., E.E. Antonova, Y.Kamide, Inconsistency of magnetic field and plasma velocity variations in the distant plasma sheet, in Proc. 5th International Conference on Substorms, St.Petersburg, 16-20 May 2000, ESA SP-443, 209 (2000). Yamamoto, Y., A.Matsuoka, K.Tsuruda, H.Hayakawa, A.Nishida, et al, Dense plasma in the distant magnetotail as observed by Geotail, Geophys. Res. left., 21, 2879 (1994). Zelenyi, L.M., A.V.Milovanov and G.Zimbardo, Multiscale magnetic structure of the distant tai: Self- consistent fractal approach, in New Perspectives on the Earth's Magnetotail, Geophys. Monogr. Ser., 105, 321 (1998). O.A.Troshichev, Departments of Geophysics, Arctic and Antarctic Research Institute, 38 Bering Str., 199397, St Petersburg, Russia, oleRtro(£;aari.nw.ru
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COMPRESSIONAL VARIATIONS PROPAGATING IN THE DISTANT TAIL OBSERVED BY GEOTAIL ON OCTOBER 2, 1994 Ayako Matsuoka Institute of Space and Astronautical Science
ABSTRACT Static pressure variations in the distant magnetotail lobe caused by the passage of a plasmoid are investigated. The traveling speed of the plasmoid is estimated to have been faster than the concurrent magnetosonic speed in the lobe. The magnetic field variation along the maximum variance direction was linearly related to the variation in the field strength, which suggests that a magneto-hydrodynamic compressional mode might have occurred. The propagation direction of the variation is determined from the background field direction and the maximum variance direction of the field. Shortly after the passage of the plasmoid, the relation between the field and velocity variations is consistent with the fast mode. Pressure variation in the fast mode was possibly generated in the trail of the plasmoid to restore equilibrium. INTRODUCTION The equilibrium of the magnetosphere is often disturbed by changes in the condition of the solar wind and by geomagnetic activities such as magnetic substorms. Equilibrium is restored by re-distribution of the plasma and magnetic field flux, which is always carried by variation in the static pressure in the magnetosphere. The magneto-hydrodynamic (MHD) compressional mode is one of the most effective processes to restore equilibrium in the magnetosphere. In the near-Earth magnetosphere, at distances < 30RE, where the magnetic field is intense and the plasma is dense, pressure variations have been well investigated by accurate in-situ measurements of the field and charged particles. These data have been compared with ground observations and found to be well interpreted by existing models. At distances > 30/? E, on the other hand, analysis becomes relatively difficult because the background field is often affected by the condition of the solar wind and by magnetospheric activity (Fairfield and Jones, 1996). Moreover, few ion data were available until 1993, when Geotail started to probe the mid and distant tail region. Collier et al. (1998) reported that the magnetic field strength in the lobe at GSM X - - 3 5 30RE was changed by a sudden increase in the solar wind density, and concluded that MHD equilibrium was satisfied. A travelling compression region (TCR) has often been found in the magnetotail simultaneously with high energy ions (see the review by Slavin, 1998). It has been interpreted as compression of the lobe field caused by plasmoids traveling tailward from the near-Earth reconnection site. Statistical studies suggest that the field compression by the plasmoids propagates in the great volume of the lobe. In this paper, a Geotail observation in the distant magnetotail at GSM X = —160RE, on October 2, 1994, is examined in detail. We chose this event because the ion number density in the lobe was relatively high, nearly 1 cm" 3 , and the plasma moment data with 12 seconds resolution were expected to be very reliable. The solar wind was measured simultaneously by IMP 8, and the ion number density was found to be several tens cm" 3 . The high density in the lobe obviously reflects the exceptionally high density of the solar wind. Moreover, as we will show in the next section, there was a rapid drop in the number density of the solar wind; we were initially interested in the response of the magnetosphere to this drop. We found that a plasmoid traversed the position of Geotail, and that the static pressure in the lobe varied significantly after its passage. The polarization of the magnetic field variation, as well as the relation between the magnetic field and velocity variations, are investigated from the viewpoint of the MHD compressional mode.
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DATA FROM IMP 8 AND GEOTAIL On October 2, 1994, IMP 8 observed that the solar wind around the Earth's magnetosphere was exceptionally dense. Figure 1 shows the three components of the magnetic field and ion velocity in GSM coordinates, and the number density and temperature of ion, dynamic, magnetic, and static pressures in the period 1400-1800 UT. The IMP 8 position in GSM coordinates was (X, Y, Z) = (27.7, -12.3, -0.1) RE and (28.0, -9.7, -5.1) RE at the beginning and end of this period, respectively. The static pressure is denned as the sum of the magnetic, ion thermal, and electron thermal pressures. Since the electron data are not available, the electron density is assumed to have been equal to the ion density and the electron temperature is assumed to be 141,000 K (Newbury et al., 1998). The ion density was about 60 cm" 3 at 1400 UT, increased gradually and reached a maximum of 92 cm" 3 at 1622. After decreasing slightly, it dropped rapidly from 59 cm" 3 to 21 cm" 3 around 1710. The ion temperature increased from 3.0 eV to 4.5 eV at the same time. The mean temperature for 1400-1709 was 1.7 eV while that for 1711-1800 was 8.3 eV. Before 1710, the solar wind speed was nearly constant at about 350 km/s; it increased slightly around and after 1710, and was about 370 km/s by 1800. The magnetic field was rather variable throughout the period shown in Figure 1. Before 1710, the dynamic pressure and the static pressure varied in the same manner as the number density, because the velocity and temperature were nearly constant. The contribution of the magnetic pressure was negligible. Simultaneously with the density drop at 1710, the dynamic pressure decreased from 15 nPa to 5 nPa. The static pressure did not drop to a similar extent at this time, because the decrease in the thermal pressure was compensated by an increase in the magnetic pressure. However, the static pressure fell from 0.18 nPa to 0.15 nPa later, at 1722 UT, when the magnetic pressure decreased. In the near-Earth magnetosphere, magnetic field strength is strongly influenced by the dynamic pressure of the solar wind. At 1715 UT, GOES 6, a geosynchronous satellite, observed the magnetic field intensity to decrease by 7% at the local time of about 8.7 h.
Fig. 1. The magnetic field and ion velocity in GSM coordinates, the number density and temperature of ion, dynamic (thick line), magnetic (thin line), and static (dot line) pressures observed by IMP 8 during 1400-1800 UT on October 2, 1994. The IMP 8 position was (27.7, 12.3, -0.1) RE at 1400 and (28.0,-9.7, -5.1) RE at 1800 in GSM coordinates.
Fig. 2 The magnetic field and ion velocity in GSM coordinates, the ion number density and temperature, and the static pressure observed by Geotail during 1730-1835 UT on October 2, 1994. The Geotail position at 1800 was (-160.6, 3.5, 6.0) RE in GSM coordinates, and (160.5, 7.7, 4.9) RE if the aberration of the solar wind by the Earth's orbital motion and the magnetic dipole tilt are taken into account. In the bottom panel the dynamic and static pressure of the solar wind at the same X position as Geotail are plotted by open circles and dots, respectively.
On that day, Geotail was located in the magnetotail at about 160 R E from the Earth. The solar wind takes about one hour to travel from the IMP 8 position to a position at the same X as Geotail. Figure 2 shows the three components of the magnetic field and ion velocity in GSM coordinates, the ion number density and temperature, and the static pressure during the interval 1730-1835 UT. The electron thermal pressure is assumed to be 14% of the ion thermal pressure. Geotail's position at 1800 was (-160.6, 3.5, 6.0) RE in GSM coordinates, and (160.5, -7.7, 4.9) RE if the aberration of the solar wind by the Earth's orbital motion and the magnetic dipole tilt are taken into account. The X component of the field, Bx, was nearly constant at -20 nT, except during the period 1753-1801, and Geotail was apparently located in the southern lobe. It is noteworthy that before 1820 the ion number
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density was always about 1 cm 3 . The relatively high density in the lobe evidently reflects the exceptionally high density of the solar wind seen in Figure 1. Between 1753 and 1801, on the other hand, Bx remained around zero, and the data show a clear plasmoid signature. The tailward ion flow speed was enhanced to about 500 km/s and reached 650 km/s at 1759. The bipolar variation of the field's Z component, Bz, indicates an O-type field configuration in the plasmoid. Magnetic Pi2 activity corresponding to the plasmoid seen in Figure 2 was seen at Kakioka Station, where the local time was about 03 h. A Pi2 pulsation with a period of about 50 seconds started at 1732 UT and lasted for 12 minutes. If the plasmoid is assumed to have been ejected tailward at X — —20RE at 1732, the traveling speed of the head would have been about 710 km/s. This is faster than the typical speed of plasmoids investigated in previous statistical studies, but is not exceptional. The variations in the dynamic and static pressure of the solar wind at the same X position as Geotail are extrapolated from the IMP 8 data by using a few assumptions; the speed of the solar wind does not change as it flows tailward and the flow speed is uniform in the Y — Z plane. In the bottom panel of Figure 2, the dynamic pressure is shown as open circles and the static pressure as dots, hi the distant tail (X < — 130RE), where tail flaring is not effective (Slavin et al. 1985), the static pressure in the lobe naturally balances the static pressure in the solar wind. The mean static pressure measured by Geotail in the lobe between 1730 and 1753 UT was 0.193 nPa. On the other hand, the mean static pressure observed by IMP 8 during the corresponding period was 0.191 nPa, which agrees with the Geotail measurement very well. The static pressure observed by Geotail increased in the first half period of the plasmoid, and fell rapidly in the second half period. It continued to decrease after Geotail returned into the lobe at 1801, and had a minimum value of 0.118 nPa at 1804. The static pressure observed by Geotail increased again after 1804, and returned to the same level seen in the lobe before the passage of the plasmoid. The drops in the dynamic and static pressures of the solar wind corresponding to those seen in Figure 1 occurred at 1806 and 1816, respectively. We should recall that some assumptions were used to derive these times, and that there is therefore some ambiguity associated with them. Nevertheless, the density drop in the solar wind must certainly have reached the same X position as Geotail before 1820, when the number density inthemagnetosphere fell toO.l cm" 3 . The static pressure in the magnetosphere started decreasing at 1817, and lower than that of the solar wind after 1820. ANALYSIS OF THE MAGNETIC FIELD AND VELOCITY VARIATIONS hi Figure 2 the static pressure in the tail lobe is found to have varied substantially between 1802 and 1820. There are two possibilities for the origin of this variation; one is the passage of the plasmoid and the other is the rapid fall of the solar wind density seen in Figure 1. Alternatively, the variation could be related to both these effects. The mean magnetosonic speed VMS = y/V% + C'g between 1802 and 1820 UT in the lobe at Geotail's position was 494 km/s, where VA is (he Alfven speed and Cs is the sound speed. Meanwhile, in the previous section, we estimated the traveling speed of the head of the plasmoid to have been 710 km/s by assuming that it was launched from X = —20RE at 1732. If we assume that the end of the plasmoid seen by Geotail at 1801 was also launched from X = —20RE at 1732, the estimated mean speed is 510 km/s. This number gives practically the lower limit of the traveling speed of the plasmoid end because we should consider that the plasmoid end must have been launched after 1732, although the exact time cannot be known. When the speed of the plasmoid end is faster than the magnetosonic speed, the pressure is not in equilibrium in the lobe just after the passage of the plasmoid. Equilibrium is re-established after a period of time sufficiently longer than the transit time of the magnetosonic variation has elapsed.
Fig. 3 The magnetic field in GSM coordinates, the field strength \B\ and the static pressure observed by Geotail during 18001830 UT on October 2,1994. The X component is inversely scaled.
The sudden drop in the solar wind pressure observed by IMP 8 at 1710 UT could naturally have caused a corresponding pressure decrease in the magnetosphere. Because the magnetosonic speed was faster than the solar wind speed of 350 km/s, the change in pressure might have propagated tailward in the magnetosphere before the pressure drop in the solar wind. In the lobe, where the plasma 0 is low, the variation in the static pressure comes mainly from variation in the magnetic field strength. Figure 3 shows the variations in the three field components (Bx,By, and Bz) in GSM coordinates, the field strength |fl| and the static pressure between 1800 and 1830 UT, as observed by Geotail in the lobe. Here, it can easily be seen that the field points mainly in the -X direction, and that the —Bx variation is similar to that of |B|. A considerable parallel component in the field variation is apparent. Regarding By, the variation is again similar to that of |B|, and the range is less than that of
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Bx • Bz varies significantly before 1815, but is nearly constant later ; the maximum field variance direction differs before and after 1815. We analyzed the data in more detail for two separate intervals: 1802-1813 UT (Interval I) and 1816-1820 UT (Interval II). Minimum/maximum variance analysis was applied to the field data. We derived the maximum, medium, and minimum variance directions and the corresponding eigenvalues, Ai, A2, and A3, respectively. In Interval I the ratio between Ai and A2 is 62, and A2/A3 is 1.7. In Interval II, Ai/A 2 = 41 and A 2 /A 3 = 25. Figure 4(la) and (2a) show the field along the maximum variance direction versus the field strength in Intervals I and II, respectively. The field is linearly related to the field strength, and therefore the magnetosonic mode may be proposed to interpret the nature of the field variations. Figure 4(1 a) shows that the amplitude of the field variation was about 10 nT which is half of the field intensity. Strictly speaking, when the intensity of the field perturbation is comparable to the background field intensity, we cannot use the formulas for the MHD magnetosonic mode straightforwardly. Rather, we must divide the interval into a number of shorter intervals and define the instantaneous background fields individually, to make the perturbation amplitude smaller than the background field intensity by a sufficient amount (Matsuoka et al., 2000). Nevertheless, in the following analysis, we consider the mean field in Interval I as the fixed background field throughout this interval for the following two reasons. First, in Interval I, the angle between the magnetic field and the mean field is less than 12.5 degrees, so the direction of the background field does not seem to change much. Second, we will not see the absolute intensity of the perturbation, but rather the differentials of the variation components. We must note that the variation components shown below have offsets that change slowly. Generally, both the propagation direction of waves and the normal direction of discontinuity structure are determined by minimum variance analysis when the field varies in a plane. Meanwhile, the high Ai/A 2 ratios in both intervals suggest that the field variations are polarized mainly in one direction. Hence we adopted a different method to determine the propagation direction of the variation in the case of the magnetosonic mode. When the variation is one dimensional, the field variation A B is naturally perpendicular to the propagation direction, k = (d/dx, d/dy, d/dz). In the magnetosonic mode, moreover, AB, k, and the background field Bo must be in the same plane (Miyamoto, 1989). The line along the direction k can be determined when Bo and A B are derived from the data. There are two possible k directions, which are parallel and antiparallel to this line.
Fig. 4 (la) The magnetic field along the maximum variance direction versus the field strength in Interval 1(1802-1803 UT on October 2, 1994). (lb) Open circles, dots and triangles respectively indicate the unit vectors along the mean field, the maximum variance direction of the magnetic field and one of the probable k directions, in Interval I. The upper and lower panels are the plots in the x — y and x — z planes, respectively, in GSM coordinates, (lc) (top) Vk — cos 0Vj| is plotted versus B\\ — Bo, where the subscripts k and || mean the components along the k and Bo directions, respectively, (bottom) The ordinate is scaled by the component of cos OV, which is perpendicular to Bo and closest to the maximum variance direction of the field in Interval I. The abscissa is scaled by the component of the field in the same direction. The dashed and dotted lines respectively represent the slopes expected from the theory of the fast and slow mode variations. The panels from (2a) to (2c) are plotted in the same format as the panels from (la) to (lc), but for Interval II (1816-1820 UT on October 2, 1994). The unit vectors along the mean field for Intervals I and II are (-0.98, -0.17, -0.08) and (-0.96, -0.15, 0.22) in GSM coordinates, respectively, and are shown as open circles in Figure 4(lb) and (2b). Meanwhile, the maximum variance directions for Intervals I and II are (-0.61, -0.23, 0.76) and (-0.84, -0.54, -0.07), respectively, and are shown as dots. Triangles represent the unit vectors along one of the probable k directions, which make acute angles with the background field. The components are (-0.77, -0.05, -0.63) and the angle with respect to the background field, 8, is 35.0 degrees for Interval I. The corresponding quantitates are
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(-0.47, 0.67, 0.58) and 61.1 degrees for Interval II. The background field direction is not much different between Intervals I and II. The field variance direction is close to the x — z plane in Interval I, and to the x — y plane in Interval II. The k direction changes by an amount corresponding to the change in the maximum variance direction between Intervals I and II. MHD theory suggests that the magnetic field and velocity variations, AB and A V, should satisfy Faraday's law, V x (A V x Bo) — dAB/dt. When the variation propagates in the k direction with velocity ui/k = Vp, we may replace d/dt and Vx with u and k x, respectively. The relations in the directions parallel and perpendicular to B o are then expressed by the following equations. AVfe-cosflAVj, = ^ A f l n Bo
(1)
cos9AVL=z-^AB1 (2) Bo In the above equations AVk — k • AV and cos 0 — k bo where k and bo are the unit vectors along k and Bo, respectively. The subscripts || and -L represent the components parallel and perpendicular to B o , respectively. If the variations propagate in the fast magnetosonic mode, the phase velocity Vp equals that of the fast wave, Vj = y {V^s + \fV^s In the slow mode, on the other hand, the phase velocity is V, = \J(V^S — \/V^s
- 4C|Vjf cos2 0)/2.
— 4C|Vjf cos2 0)/2.
In the upper panel of Figure 4(lc), Vk — cos 0Vj| is plotted versus S|| — Bo in Interval I, when we adopt the k direction of cos 0 > 0. From Equation (1) the slope of the plots must be Vj /Bo if the variations propagate in the fast magnetosonic mode. Conversely, they must have a slope of VJ Bo in the slow magnetosonic mode. The dashed and dotted lines in the panel show slopes of Vj I Bo and Vs/ Bo, respectively, when we adopt Vj =453 km/s, Vs =70 km/sand Bo = 19.1 nT, as determined in the above analysis. We find a clear positive correlation in the data, and the slope roughly agrees with the dashed line. In particular, for the data points where B|| changes rapidly from-1.5 nT to 0.8 nT (from 1806:02 to 1807:38 UT) the agreement is remarkably good. In the bottom panel, the ordinate is scaled by the component of cos 6Vj_ closest to the maximum variance direction of the field, and the abscissa is scaled by the component of B± in the same direction. The dashed and dotted lines show slopes of — Vj/Bo and —Vs/Bo, respectively, as expected from the relation represented by Equation (2). The correlation is negative, and the agreement with the dashed line is again good. This good agreement indicates that the variations possibly propagate in the fast mode and cos 8 > 0. The group velocity and the Poynting vector are found to have the same direction as that of k from the theory of fast waves. Figure 4(2c) demonstrates the same comparison as Figure 4(lc) for Interval II. The dashed and dotted lines in the panel show slopes of Vj I Bo and Vj /Bo, respectively, where Vj = 612 km/s, Vs = 27 km/s and Bo = 20.4 nT. Again we find substantial linear relations, positive in the top panel and negative in the bottom one. The slopes of the data in this interval are close to the dotted lines, so that it might be possible to interpret the variations as the slow mode. We must note, however, that the ranges in Vk — cos 0Vj| and cos 0V± are short, respectively only 14 km/s and 8 km/s. Because of the limitation of the accuracy of the plasma moment data, it is difficult to conclude that the variation is due to the slow mode. Meanwhile, we have tested another method which has been prevalently used in previous studies to identify the wave mode in higher /? plasmas. In compressional (fast and slow) waves, the perturbation in the thermal pressure APth and the perturbation in the field strength AB[| are related as (Song et al., 1994)
APth/Ptho = (1 -
C
yV y^A-Bii/Bo
(3)
P
where Ptho, Vp and 7 are the background thermal pressure, the phase velocity, and the polytropic index, respectively. For fast waves in the low-/? plasma (C's "C VA < Vj — Vp), this equation becomes APth/Ptho — "/AB^/Bo. For Interval I, By increased monotonously during 1806:02 - 1807:38 UT and ASy = 2.3 nT. If the variation is the fast mode, from the above equation, Pth is expected to increase by 20 % in the same period when 7 = 5/3 and Bo = 19.1 nT. In the Geotail observation, the thermal pressure did not increase monotonously; it varied randomly and the variance range was 13 %. Considering the measurement accuracy of the plasma temperature, it is difficult to identify the fast mode wave just by testing Equation (3). The former method, investigating the relation between the field and velocity variations, appears to be much more effective to identify the fast mode in Interval I. On the other hand, for slow waves in the low-/? plasma, the ratio between APth/Ptho and AB^/Bo strongly depends on the propagation direction, 6. For Interval II, Sj| increased monotonously during 1817:11-1819:11 and AS n =2.2nT. If the variation
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is the slow mode, APth/Ptho is calculated to be -38 when Cs = 56 km/s, Vs = 27 km/s, i?o = 20.4 nT and 6 = 61.1 degrees. It means that the thermal pressure should decrease at least by an order. In this period the measured thermal pressure decreased but not so significantly; the depression was 75 %. It seems that the test of Equation (3) denies the possibility of the slow wave for the variations seen in Interval II. DISCUSSION ON THE SOURCE OF THE VARIATIONS The pressure variations in the lobe seen in the previous section may be caused by either the passage of the plasmoid or the drop of the solar wind pressure. It is difficult to distinguish the source of the variations just from the expected arrival times of the plasmoid and the low-density solar wind to the same X position as Geotail, because they are almost simultaneous. In this section we show an interpretation for the source of the fast wave seen in Interval I by considering its propagating direction.
Fig. 5 The southern part of the plasmoid and the surrounding lobe are illustrated. The plasmoid is assumed to be elongated in the GSM Y direction, (a) The boundary of the southern and rear part of the plasmoid is represented simply by two planes, Plane 1 and 2. Vp and VPM denote the phase speed of the fast wave and the speed of the plasmoid, respectively, (b) The boundary of the plasmoid cross-section is represented by a curve. Solid curves represent the magnetic field lines. The dashed arrow shows the path of Geotail relative to the tail structure.
In the previous section the wave was shown to have propagated southward in the southern part of the lobe, and hence it is natural to interpret that the wave had been excited in the center of the tail, possibly by the plasmoid. The passage of the plasmoid through the lobe is considered to have left a low pressure region behind, in the central lobe. When the speed of the plasmoid end is faster than the magnetosonic speed and the pressure in the external region is relatively high, the magnetotail just after the passage of the plasmoid is presumably far from equilibrium. Field and velocity variations of the magnetosonic mode are naturally excited to restore equilibrium in the magnetotail. Here we assume that the plasmoid is elongated in the GSM Y direction. We cannot state the specific shape of the plasmoid cross-section in the XY plane from the Geotail data. Now we consider a simplified model in which the boundary of the southern and rear part of the plasmoid consists of two planes as illustrated in Figure 5(a). We assume that Plane 1 and 2 moves in the —X direction with the velocity of the plasmoid, VPM (> KM S ). The normal components of the velocity of Plane 1 and 2 are defined as Vnl and Vn2, respectively. When the angle between Plane 1 and the XY plane is small so that Vni is much less than VMS, the equilibrium is always satisfied in the lobe plasma in Region (i). On the other hand, when angle between Plane 2 and the XY plane is large so that Vn2 > VMS, rarefaction waves are excited in Region (iii), and propagate from Region (iii) to Region (ii). Just behind the wave front in Region (ii) (represented by a dash line), the propagation direction is southward and tailward, which is determined by kx = —VP/VPM. Of course the real boundary of the plasmoid cross-section is possibly not angular one as illustrated in Figure 5(a). However, even when the boundary is round as illustrated in Figure 5(b), we may consider in the same way; the normal speed of the boundary would be faster than VMS on the sunward side of the plasmoid, while it is much slower than VMS on the southward side. The rarefaction waves propagate southward and tailward in the southern part of the lobe behind the plasmoid except the region close to the neutral sheet (Figure 5(b)). Geotail has possibly observed these waves after getting away from the plasmoid (Figure 5(b)). In the event shown here, the traveling speed of the plasmoid end is estimated to have been faster than 510 km/s and slower than 710 km/s. Therefore —Vp/VpM is estimated to be between 0.64 and 0.89 where Vp is the phase speed of the fast wave, 453 km/s. The —X component of the k determined in Figure 4(lb) is 0.77, which is within this range. In fast waves, the Poynting vector is in the same direction as the phase velocity. As Slavin (1998) stated, the field and velocity variations in the magnetosonic mode must be excited after the passage of plasmoids through the usual magnetotail when the plasmoid traveling velocity exceeds the magnetosonic speed. However, with Geotail data, it would not always be possible to analyze such variations using the method that we have used here. The number density in the distant lobe decreases when approaching the plasma sheet (Siscoe and Sanchez, 1987); thus the density at Geotail, located near the distant plasmasheet, is often very low. Generally, the velocity data of Geotail show significant scatter and are less accurate when the number density is under 0.1 c m " 3 . Moreover, when the number density is 0.1 cm" 3 , the field intensity must be less than 7 nT to make the magnetosonic speed slower than the plasmoid traveling speed of 500 km/s. When the field variation is less intense than the case investigated here, definition of the maximum variance direction and the k direction may be more difficult. Selection of data acquired when the number density in the lobe is relatively high and the field is simultaneously intense seems to
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be essential to determine the k direction. In the previous section, the relation between the field and velocity variations in Interval II appears to be similar to that expected for slow mode variations, but the pressure variation is found to be too small to interpret it as the slow mode. Moreover, if we consider that these changes were generated by the pressure drop in the solar wind, it is impossible to interpret the variations as slow waves for a simple reason. When we assume that the distance between Geotail and the magnetopause is 30 RE and the wave velocity is 27 km/s everywhere, a slow wave generated on the magnetopause would take about two hours to travel to the Geotail position. This is inconsistent with the observation result in which the pressure change at the Geotail position occurred nearly simultaneously with the pressure drop in the solar wind at the same X position as Geotail. On the other hand, it is also possible to interpret the results as showing that equilibrium was established in the lobe in a short time before and after the arrival of the pressure drop in the solar wind. If this interpretation is correct, the real arrival time of the pressure drop at the same X position as Geotail must have been between 1817 and 1820 UT, later than plotted in Figure 2. ACKNOWLEDGMENTS We are grateful to Drs. T. Mukai and T. Nagai for providing the LEP and MGF data from Geotail. The IMP 8 solar wind data were provided by the MIT Space Plasma Physics Group. The IMP 8 magnetic field data are from the Goddard Space Flight Center, NSSDC, courtesy of Dr. R. Lepping. The GOES 6 magnetic field data were provided by NOAA/SPIDR. The Kakioka magnetometer data were provided by the Kakioka Geomagnetic Observatory. REFERENCES
Collier, M. R., J. A. Slavin, R. P. Lepping, K. Ogilvie, A. Szabo, H. Laakso, and S. Taguchi, Multispacecraft observations of sudden impulses in the magnetotail caused by solar wind pressure discontinuities: Wind and IMP 8, J. Geophys. Res., 103, 17,293-17,305, 1998. Fairfield, D. H., and J.Jones, Variability of the tail lobe field strength, J. Geophys. Res., 101, 7785-7791, 1996. Matsuoka, Ayako, David J. Southwood, Susumu Kokubun, and Toshifumi Mukai, Propagation sense of low-frequency MHD waves in the magnetosheath observed by Geotail, J. Geophys. Res., 705,18,361-18,376,2000. Miyamoto, K., Plasma physics for nuclear fusion, Rev. ed., Cambridge, Mass. MIT Press, 1989. Newbury, J. A., C. T. Russell, J. L. Phillips, and S. P. Gary, Electron temperature in the ambient solar wind: Typical properties and a lower bound at 1 AU, J. Geophys. Res., 705,9553-9566,1998. Siscoe, G. L., and E. Sanchez, An MHD model for the complete open magnetotail boundary, J. Geophys. Res., 92, 7405-7412, 1987. Slavin, J. A., Traveling compression regions, New Perspectives on the Earth's Magnetotail, AGU, 225-240,1998.
Geophysical Monograph 105,
Slavin, J. A., E. J. Smith, D. G. Sibeck, D. N. Baker, R. D. Zwickl, and S.-I. Akasofu, An ISEE 3 study on average abd substorm conditions in the distant magnetotail, J. Geophys. Res., 90,10,875-10,895,1985. Song, P., C. T. Russell, and S. P. Gary, Identification of low-frequency fluctuations in the terrestrial magnetosheath, J. Geophys. Res., 99,6011-6025, 1994.
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DAYSIDE OUTER MAGNETOSPHERE ULF WAVES OBSERVED BY GEOTAIL T. Sakurai1, Y. Tonegawa1, Y. Shinkai2 and M. Nowada1 1
Department Aeronautics and Astronautics, School of Engineering, Tokai University, 1,117 Hiratsuka, 259-1292 Japan 2 Graduate University Advanced Study, National Institute of Polar Research, 1-9-10 Kaga, Itabashi, Tokyo, 1738515 Japan
ABSTRACT On the basis of our recent studies (Sakurai, et al., 1999a, 1999b and 2001) the present paper is intended to review oscillation and propagation characteristics of Pc 3 and Pc 5 ULF waves observed near the dayside magnetopause by the Geotail satellite. These characteristics are studied on the basis of the simultaneous observations of the magnetic and electric fields, and low energy plasmas measured with the excellent instruments on board the satellite. The study reveals that the dominant ULF waves observed near the dayside magnetopause are Pc 3 and Pc 5 oscillations. The Pc 3 oscillations appear with a peak power around noon at the frequency of 25 mHz in the azimuthal component of the magnetic field This frequency component shows clear resonant oscillations. In addition fast mode earthward propagation is recognized. The Poynting flux of Pc 3 waves is about 1-10 nW/m2 on average, the strongest being along the magnetic field-line. Pc 5 oscillations are also dominant, and are observed mainly in the dawn and dusk-side flanks. They appear as clear oscillations in the radial component of the electric field, suggesting that resonant oscillations along the magnetic field-line are well established. The Poynting flux of Pc 5 can be estimated as 10 - 100 nW/m2 in both directions across and along the magnetic field-line. This result suggests that Pc 5 wave energy is carried into both the ionosphere and the inner magnetosphere during an hour with the energy of 1010 - 1013 J, which is one or two orders of magnitude less than the substorm energy. By taking into account the continuous activation of ULF waves in the outer magnetosphere, these waves should play an important role in the energetics of the magnetosphere. INTRODUCTION The Geotail satellite frequently measures Pc 3 and Pc 5 ULF waves during the skimming of the dayside outer magnetosphere. The instrumentation on board the satellite is excellent for studying ULF wave characteristics since the satellite simultaneously measures the magnetic and electric fields and low energy plasma, whose interrelationships make possible new physical insights into ULF wave studies. A number of studies of ULF waves so far have been based on satellite data, which have been restricted only to magnetic field data (Anderson and Engebretson, 1995; Cao et al., 1994). A very few satellites have provided the data of the simultaneous electric field variations, such as DE-1, ISEE's and GEOS 2 satellites (Chi and Russell, 1998; Junginger et al., 1985). Therefore, our knowledge of the electric field and plasma signatures of ULF waves in the magnetosphere seems to be substantially important. In this study we examine wave characteristics of Pc 3 and Pc 5 observed in the outer magnetosphere. In addition, Poynting fluxes are used to clarify how these ULF play an important role in the energetics of the magnetosphere.
DATA AND ANALYSIS PROCEDURE Data used in this study are the magnetic and electric fields and plasma data when the satellite surveyed the dayside outer magnetosphere. The instruments and data processing techniques are described in detail for the magnetic and electric fields in Kokubun et al. (1994) and Tsuruda et al. (1994), respectively, and for the plasma data in
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Mukai et al. (1994). The latter instrument covers the energy from 32 eV to 39 keV. In Figures 1 and 2 data are shown in mean-field-aligned coordinates. The details are described in Sakurai et al. (2001). Pc 3 AND Pc 5 ULF WAVES OBSERVED IN THE OUTER MAGNETOSPHERE The magnetic and electric field oscillations of 14 January 1995 are examined when the Geotail satellite traversed from the dawn- to the post noon side outer magnetosphere with a path through the subsolar magnetosphere. The satellite crossed the magnetopause outbound on the afternoon side, and observed various types of magnetic and electric field oscillations characterizing each region of the outer magnetosphere. Double frequency oscillations of low energy plasmas One of the interesting phenomena is the large amplitude Pc 5 electric field oscillation appearing in the dawn side from 07h38m MLT (07h30m UT) to O8hO2m MLT (08h00m UT) and near L = 10, which is shown in Figure 1. It is very clear that the large amplitude Pc 5 oscillations appear in both the electric field and the velocity field. Magnetic field oscillations are not clear because this Pc 5 oscillation is observed in the vicinity of the magnetic equator and the oscillation is a fundamental odd mode oscillation with anti-node at this location. It is very interesting that this Pc 5 oscillation accompanies the oscillation of low energy plasmas with a frequency double the frequency of Pc 5 (lower panels). This is due to the electric field drift motion of the low energy plasmas. The large amplitude radial electric field oscillations of the Pc 5 wave cause the azimuthal drift motion of the low energy plasmas as shown in the top and middle panels. The density peaks correspond to troughs of the plasma temperature and/or vice versa. These density peaks are higher than the background density, while the temperature troughs are lower than the temperature of the background plasma because of the presence of plasma with energy less than the threshold of the plasma instrument. The detector measures the lower energy plasmas conveyed to the detector due to the wave electric field oscillations. The details are described in Sakurai et al. (1999a).
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-1 . 0
5D85.0
1358.0 08 ; 00
Fig. 1. Relations between the electric field and velocity field Pc 5 oscillations and the double-frequency low energy plasmas.
Fig. 2. Poynting flux of Pc 5 oscillations. The radial Sx, azimuth Sy, field-aligned Sz and perpendicular components are presented from the top to the bottom panels, respectively.
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Poynting flux of Pc 5 The Poynting flux of this Pc 5 oscillation is shown in Figure 2, again in mean-field coordinates. The Poynting flux along the magnetic field-line (Sz) is almost zero on average, suggesting that the Pc 5 energy bounces back and forth along the magnetic field-line. On the other hand, the Poynting energy propagates tailward perpendicular to the magnetic field-line, and its average value can be estimated as 100 - 1000 nW/m 2 , which may correspond to a total energy 107-109 W, if we assume that the area of the flank should be 2xlO15 m2 by considering the longitude and the latitude extents to be 90°x20° for the flank (Sakurai et al. , 2001). This result suggests that the wave energy of Pc 5 is carried into the inner magnetosphere during an hour with the energy of 1010 - 1013 J, which is one or two orders of magnitude less than the substorm energy of 1014 - 1015 J (Greenwald and Walker, 1980; Sakurai et al., 2001). Pc3 ULF waves The other interesting and frequently observed oscillations are Pc 3's in the frequency range from 20 to 30 mHz. A remarkable activity of Pc 3 signals appears with a lumped enhancement in the azimuthal component of the magnetic field, which is exhibited with a peak power around the sub-solar magnetopause. This result is consistent with the previous satellite study by Chi and Russell (1998) based on the ISEE 1 satellite magnetic field data. Dynamic spectral characteristics of Pc 3 oscillations The other point to be noted is the dynamic spectral characteristics studied by Sakurai et al. (1999b) in which the horizontal structure centered at the frequency of about 25 mHz is clearly recognized in the azimuthal component of the magnetic field. The frequency of the spectral band varies in a 10 to 20 minute interval with an intermittent enhancement of the activation. Poynting flux of Pc 3 The characteristics of Poynting flux were also studied by Sakurai et al. (1999b). The energy propagates earthward and in the field-aligned direction. Each component shows a value of 1-10 nW/m2 on average. In the direction along the magnetic field-line the energy repeatedly bounces back and forth along the magnetic field-line, thus establishing the resonance oscillation. If we consider the energy flux crossing the dayside magnetopause, we can estimate an hourly value flowing into the inner magnetosphere as 109-10u J, which is one or two orders of magnitude less than that of Pc 5. The details should be referred to Sakurai et al. (1999b). DISCUSSION AND CONCLUSIONS Many examples of Pc 3 and Pc 5 ULF waves have been examined by using the Geotail observations. Typical examples of Pc 5 and Pc 3 oscillations are discussed above on the basis of the 14 January 1995 events. Our study reveals interesting Pc 3 and Pc 5 phenomena characterizing each localized region of the outer magnetosphere. The most interesting and important thing is about the Poynting flux of Pc 3 and Pc 5 wave energy. As described in the previous sections by using the simultaneous observation of the magnetic and electric field oscillations Pc 5 waves carry the wave energy of about 1010-1013 J/hour, which is one or two orders of magnitude less than the substorm energy. Pc 3 waves convey the energy of one or two orders of magnitude less than that of Pc 5. However, it should be remembered that these wave energies are continuously flowing into the inner magnetosphere, and then, their role for energetics in the inner magnetosphere seems to be important. Similar examinations of the Poynting flux of Pc 3-4 waves in the outer magnetosphere have been reported in Chi and Russell (1998). As for the propagation of Pc 3 waves, we should recall, however, some of the evidence for the Pc 3 energy intrusion across the magnetopause into the magnetosphere, which has been examined by Song et al. (1993) based on the ISEE satellite data. They show that energy transmission from the magnetosheath into the magnetosphere might be expected to be mainly due to slow mode waves, and then the transmitted energy might excite the eigen mode of the local field-line resonant oscillation. This scenario can be verified also in our study. The most pronounced spectrum with a high Q value of the transverse component of the magnetic field oscillations is obtained when the satellite traverses inside along the magnetopause. This result seems to be reasonable when we consider the source wave propagates across the magnetopause into the magnetosphere from the magnetosheath.
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We can summarize our results as follows: Pc 3 oscillations are most frequently observed near the subsolar region, and Pc 5 oscillations are mainly observed in the dawnside of the outer magnetosphere. Their oscillations appear as a resonance along the magnetic field-line. Strong resonance occurs in the region far from the magnetopause. Closer to the magnetopause compressional and impulsive oscillations with the frequencies of Pc 5 are observed dominantly. Poynting fluxes of Pc 3 and Pc 5 ULF waves are estimated and found to be important for the energetics in the inner magnetosphere because of their continuous activation throughout the day. ACKNOWLEDGEMENTS We would like to express our sincere thanks to A. Nishida, S. Kokubun, K. Tsuruda, T. Mukai and all members of the GEOTAIL project team, and to two referees for their critical reading our manuscript and valuable suggestions. REFERENCES Anderson, B. J. and M. J. Engebretson, Relative intensity of toroidal and compressional Pc 3-4 wave power in the dayside magnetosphere, J. Geophys. Res., 100(A6), 9591-9603, (1995). Cao, M., R. L. McPherron, and C.T Russell, Statistical study of ULF wave occurrence in the dayside magnetosphere, J. Geophys. Res., 99, 8731-8753, (1994). Chi. P. J., and C. T. Russell, Phase skipping and Poynting flux of continuous pulsations, J. Geophys. Res., 103(A12), 29,479-29,491,(1998). Greenwald, R. A., and D. M. Walker, Energetics of long period resonant hydromagnetic waves, Geophys. Res. Lett. 7, 745-748, (1980). Junginger, H., G. Haerendel, and F. Melzner, A statistical study of wave Poynting vectors measured during longperiod magnetospheric pulsations at geostationary orbit, J. Geophys. Res., 90, 8301-8307, (1985). Kokubun, S., T. Yamamoto, M.H. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46, 7-21, (1994). Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particles (LEP) experiment onboard the GEOTAIL Satellite, J. Geomag. Geoelectr, 46, 669-692, (1994). Sakurai, T., Y. Tonegawa, T. Kitagawa, M. Nowada, A. Yamawaki,T. Mukai, S. Kokubun, T. Yamamoto, and K. Tsuruda, Double-frequency oscillations of low energy plasma associated with transverse Pc 5 pulsations: GEOTAIL satellite observations, Earth Planets Space, 51, 43-53, (1999a). Sakurai, T., Y. Tonegawa, T. Kitagawa, K. Yumoto, T. Yamamoto, S. Kokubun, T. Mukai, and K. Tsuruda, Dayside magnetopause Pc 3 and Pc 5 ULF waves observed by the GEOTAIL Satellite, Earth Planets Space, 51, 965-978, (1999b). Sakurai, T., Y. Tonegawa, Y. Shinkai, K. Yumoto, S. Kokubun, K. Tsuruda, and T. Mukai, Poynting vectors of Pc 5 pulsations observed by the GEOTAIL satellite in the dayside outer magnetosphere, Earth Planets Space, 53, 843-849, (2001). Song, P., C. T. Russell, R. J. Strangeway, J. R. Wygant, C. A. Cattell, R. J. Fitzenreitter, and R. R. Anderson, Wave properties near the subsolar magnetopause: Pc 3-4 energy coupling for northward interplanetary magnetic field, J. Geophys. Res., 98(A1), 187-196, (1993). Tsuruda, K., H. Hayakawa, M. Nakamura, T. Okada, A. Matsuoka, F. S. Mozer, and R. Schmidt, Electric field measurements on the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 693-712, (1994). E-mail address: sakuraifgjms.u-tokai.ac.jp
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APPLICATION OF SPACECRAFT POTENTIAL TO INVESTIGATE THE DISTRIBUTION OF LOW-ENERGY PLASMA IN MAGNETOSPHERE K. Ishisaka1, T. Okada1, H. Hayakawa2, T. Mukai2, and H. Matsumoto3 1
3
Faculty of Engineering, Toyama Prefectural University, Kosugi, Toyama 939-0398, JAPAN 2 Institute of Space and Astronautical Science, Sagamihara, Kanagawa 229-8510, JAPAN Radio Science Center for Space and Atmosphere, Kyoto University, Uji, Kyoto 611-0011, JAPAN
ABSTRACT
We have found that the spacecraft potential in the magnetosphere and the solar wind can be used to derive the electron number density of the plasma surrounding spacecraft. The relationship between the Geotail spacecraft potential and the electron number density as determined by the plasma wave observations in the solar wind and broader magnetosphere (except for the high-density plasmasphere) was investigated and an empirical formula correlating these values obtained. Using this empirical formula and plasma particle measurements, we have shown the distribution of low-energy plasma in the magnetosphere. INTRODUCTION Spacecraft in sunlight is generally charged to a positive potential relative to the ambient plasma in the solar wind and almost all the regions of the magnetosphere because of the photoelectron emission from the spacecraft conductive surface. Several studies have shown the relation between the spacecraft potential and the electron number density surrounding the spacecraft (Pedersen et al., 1984; Schmidt et al., 1987; and Escoubet et al., 1997). Escoubet et al. (1997) have determined this relationship for potentials up to 30 V. These studies have examined this relation in the solar wind, magnetosheath, and magnetosphere near the Earth. The present authors have investigated the relation between the Geotail spacecraft potential and electron number density derived from the plasma waves during the period from September 1992 to December 1995 and an empirical formula was obtained to show their relation in which the range of potentials was from a few volts up to 80 V (Ishisaka et al., 2001). In this study, we demonstrated that the electron number density derived from the spacecraft potential can be used to investigate the distribution of the low-energy plasma in the magnetosphere. GEOTAIL SPACECRAFT POTENTIAL The spacecraft potential (Vs/C) of Geotail is measured by a single probe system, which is one of the subsystems of the electric field detector (EFD) (Tsuruda et al., 1994). The single probe system measures the difference of potential (AV) between the spacecraft and the probe located at the end of a long wire boom; AV is given by AV = Vp — Vs/C, where V^, is the probe potential for the ambient plasma. The probe potential can be set nearly equal to the ambient plasma potential by applying the appropriate bias current to the probe. Therefore the voltage measured by the single probe system is approximately equal to the spacecraft potential relative to ambient plasma, that is Vp ~ 0, AV ^ —Va/C. The upper panel of Figure 1 shows the spacecraft potential and the plasma wave spectrogram (lower panel) between 14:00 UT and 16:00 UT on Jan. 12, 1994, when Geotail was in the plasma sheet between 14:00 -75-
„ , _,, . . . ,. Fig. 1. T h e spacecraftri potential (upper panel) , , , , . j.i , j . ,, , measured by electric field detector a n d t h e plasma ,, ,. , , wave spectrogram (lower panel) measured by . ? f' ,' ., , . L plasma wave instrument onboard Geotail during the period from 14:00 t o 16:00 U T , January 12, 1994.
Fig. 2. Characteristic between spacecraft poten° . . , . _ . . . . . . tial and electron density. O u r empirical formula is j ; . . , . shown as a solid line. T h e empirical formula given , _ , , ,.„-_. . , , , ° . by Escoubet et al. (1997) is shown as a dashed line ,.: . U s h l s a k a e t aJ -- 2^)-
UT and 15:10 UT, in the tail lobe between 15:10 UT and 15:40 UT and in the plasma sheet between 15:40 UT and 16:00 UT. The plasma wave spectrogram was observed by the plasma wave instrument (PWI) (Matsumoto et al., 1994). The continuum radiation (CR) was in the frequency range from 196 Hz to 12.5 kHz. The white line in the plasma wave spectrogram indicates the lower cutoff frequency (/ cut ) of CR. The fcut of CR indicates the electron plasma frequency surrounding the spacecraft. A good correlation between the spacecraft potential and the / c u t of CR is seen in Figure 1. The electron number density (Nc) is estimated by Nc = /,? ut /81 using the / c u t of CR and the center frequency of the Langmuir waves. In order to obtain the relation between them, the present authors investigated the Vs/C - Nc characteristics in the region where the lower cutoff frequency of CR or the Langmuir waves can be clearly seen during the period from September 1992 to December 1995 (Ishisaka et al., 2001). The Va/C - Nc relation is shown in Figure 2. A high correlation is clearly seen in the range of spacecraft potentials from a few volts up to 80 V, corresponding to an electron number density from 0.002 to 30 cm" 3 , although the measured values are scattered around the mean value. The best fit function to the data is represented by an empirical formula with four exponentials as shown by the solid line in Figure 2.
The degree of scatter of the measured value from the empirical formula is about 40%, which, because the spacecraft potential depends on the electron temperature, is inferred due to the variation of electron temperature surrounding spacecraft. INVESTIGATION OF LOW-ENERGY PLASMA USING SPACECRAFT POTENTIAL The investigation of low-energy plasmas below 32 eV using data obtained by the Geotail was reported by Matsui et al. (1999). They used the electron number density determined by the plasma waves and data obtained by the low-energy particle instrument (LEP) onboard Geotail (Mukai et al., 1994). The electron number density estimated by plasma waves includes cold components that were not measured by the particle detectors. On the other hand, from the LEP data, the energy range of ions were between 32 eV and 39 keV while the energy range of electrons were between 60 eV and 38 keV. The density and temperature of electrons measured by the LEP were not actually used, because the electrons surrounding the spacecraft are
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Fig. 3. The density profile and energy-time (E-t) spectrogram during the period from 15:00 to 21:00 UT, November 26, 1995. Top panel is the density profile. Ns/C and A^LEP are shown as solid lines and Npvli is shown as open circles. Middle and bottom panels are omnidirectional E-t spectrogram of electrons and ions. very sensitive to spacecraft potential and photoelectrons emitted from the spacecraft surface. The number density (TV) of low-energy plasma was estimated from N = Npwi — -/VLEPI where JVpwj is the total density acquired from the plasma waves and A^LEP is the partial density acquired from the LEP. However, the electron number density cannot be estimated by the plasma waves when the / c u t of CR and the center frequency of the Langmuir waves cannot be clearly seen. Moreover, there were errors that resulted from visual interpretation of the plasma wave spectrogram. In this study, we attempt to substitute the electron number density obtained by the spacecraft potential [Ns/C) for that estimated by the plasma waves. The electron number density obtained from the spacecraft potential is affected by the electrons in a wider energy range. Figure 3 shows the density profile (top panel) and the energy-time (E-t) spectrogram of electrons (middle panel) and ions (bottom panel) during the period from 15:00 to 21:00 UT, November 26, 1995. The Geotail spacecraft observed the region between the outer magnetosphere and the inner magnetosphere on the dusk-side. In the top panel of Figure 3, Ns/C is equal to Npwi between 15:00 UT and 19:30 UT. We could not obtain 7Vpwi after 19:30 UT because the weakening CR made it too difficult to determine the / c u t of CR, but we were able to use Ns/C in order to examine the low-energy plasma after 19:30 UT. There was an interval where there was a large difference between Na/C and A^LEP of more than 10 cm" 3 during the period from 15:00 to 17:40 UT. In the E-t spectrogram of Figure 3, electrons with energy ranges detectable by LEP were not observed in this period. On the other hand, ions with energies higher than 1 keV were observed throughout most of the interval. Ions above 1 keV have their origin in the plasma sheet (Matsui et al., 1999). Ions with energies lower than 1 keV were not observed. The difference between Ns/C and A^LEP indicate the existence of a large amount of low-energy plasma. This low-energy plasma represents the plasma directly supplied by the ionosphere. -77-
In the electron E-t spectrogram, it is clear that enhancements of counts of lower-energy electrons were observed after 17:40 UT. This enhancement of electrons indicates that Geotail was located in the magnetosphere. iVs/c decreased down to about 2 cm" 3 in this time. Subsequently, Ns/C increased suddenly around 19:50 UT because Geotail observations occurred outside of the magnetosphere. There was a difference between Ns/C and A^LEP of about 1 cm~3 during the period from 20:00 to 21:00 UT, however we cannot infer from the difference of 1 cm" 3 that a low-energy plasma existed in this region because Ns/C is affected by the increase in the counts of electrons with energies higher than 1 keV. CONCLUSIONS We have shown that the spacecraft potential of Geotail gives a good estimate of the electron number density in the solar wind and almost all the regions of the magnetosphere except for the high-density plasmasphere. The data have been fit to an empirical formula that gives electron number density as a function of spacecraft potential. The empirical formula is obtained by extending the relation of Escoubet et al. (1997) up to 80 V and consists of four exponential terms. The empirical formula is applied to obtain the electron number density in the solar wind and the magnetosphere in the range of spacecraft potentials from a few volts to about 80 V corresponding to an electron number density from 0.002 to 30 cm" 3 . We have also examined whether or not we should substitute Na/C for iVpwi. Consequently, we found that Ns/C can be used in order to investigate the low-energy plasma in the region where the electron and ion temperatures are lower than 1 keV. Ishisaka et al. (2001) showed that Ns/C from the empirical formula (see Equation (1)) is overestimated in the region where the plasma temperature is higher than 1 keV. Therefore, we are able to investigate the low-energy plasma in the bowshock and the magnetosheath using the difference between Ns/C and TVLEP- In future work we intend to discuss the low-energy plasma quantitatively in the bowshock and the magnetosheath. In particular, we plan to investigate the low-energy plasma in the dawn-side region where it is very difficult to obtain the electron plasma frequency because the electron cyclotron harmonic waves merge with the / c u t of CR. We will then attempt to construct a model that can explain the region supplied by the ionospheric plasma directly. With the completion of these models, we will be able to understand the details of the existence of low-energy plasma in the outer magnetosphere. REFERENCES Escoubet, C. P., A. Pedersen, R. Schmidt, and P. A. Lindqvist, Density in the Magnetosphere Inferred from ISEE 1 Spacecraft Potential, J. Geophys. Res., 102, 17,595-17,609, 1997. Ishisaka, K., T. Okada, K. Tsuruda, H. Hayakawa, T. Mukai, and H. Matsumoto, Relationship between the Geotail Spacecraft Potential and Magnetospheric Electron Number Density including the Distant Tail Regions, /. Geophys. Res., 106, 6309-6319, 2001. Matsui, H., T. Mukai, S. Ohtani, K. Hayashi, R. C. Elphic, M. F. Thomsen, and H. Matsumoto, Cold Dense Plasma in the Outer Magnetosphere, J. Geophys. Res., 104, 25,077-25,095, 1999. Matsumoto, H., I. Nagano, R. R. Anderson, H. Kojima, K. Hashimoto, M. Tsutsui, T. Okada, I. Kimura, Y. Omura, and M. Okada, Plasma Wave Observations with Geotail Spacecraft, J. Geomag. Geoelectr., 46, 59-95, 1994. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particle (LEP) experiment onboard the Geotail satellite, J. Geomag. Geoelectr., 46, 669-692, 1994. Pedersen, A., C. A. Cattell, C. G. Falthammar, V. Formisano, P. A. Lindqvist, F. Mozer, and R. Torbert, Quasistatic Electric Field Measurements with Spherical Double Probes on the GEOS and ISEE satellites, Space Sci. Rev., 37, 269-312, 1984. Schmidt, R. and A. Pedersen, Long-Term Behavior of Photo-Electron Emission from the Electric Field Double Probe Sensors on GEOS-2, Planet. Space Sci., 35, 61-70, 1987. Tsuruda, K., H. Hayakawa, M. Nakamura, T. Okada, A. Matsuoka, F. S. Mozer, and R. Schmidt, Electric Field Measurements on the GEOTAIL Satellite, J. Geomag. Geoelectr., 46, 693-711, 1994.
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GEOTAIL STUDY OF COMPARISON BETWEEN THE DOUBLE-PROBE ELECTRIC FIELDS AND THE CONVECTION ELECTRIC FIELDS IN THE DISTANT TAIL Y.Takei1. T. Mukai2, Y. Saito2, H. Hayakawa2, K. Tsuruda2 1
Department of Earth and Planetary Science, University of Tokyo, Tokyo, 113-0033, Japan 2 Institute of Space and Astronautical Science, Sagamihara, 229-8510, Japan ABSTRACT
The double probe technique is one of the standard techniques for in situ DC electric field measurements in tenuous plasmas, but it is known for its sensitiveness to the plasma environment around a spacecraft. In the present study we compare the double-probe electric fields obtained by Geotail/EFD-P with the convection electric fields —V x B in various regions of the distant tail. We have found that the sensitivity (effective length of the double probes) varies as a function of the spacecraft potential, and the offset weakly depends on the electron temperature of ambient plasma. Using this result, an empirical calibration formula for the double-probe electric fields is obtained.
INTRODUCTION The double probe technique is one of the standard techniques for in situ DC electric field measurements in space plasmas, and has been used by many spacecraft [Pedersen et al., 1998]. However, this technique is known for its sensitiveness to the plasma environment around a spacecraft, and the reliability of the doubleprobe data is a matter of concern in the space plasma physics community. Mozer et al. [1983] reported that the double probe technique is useful, based on evaluation of the data in tenuous plasmas having the density of 0.1 to 1/cc. The GEOTAIL EFD-P [Tsuruda et al., 1994] provides the double-probe electric field in the magnetotail over a wider range of plasma parameters than studied by Mozer et al. In this paper we compare the GEOTAIL double-probe electric fields with the — V x B electric fields calculated with the plasma bulk flow (V) and the magnetic field (B). Assuming that the frozen-in theory holds in the distant tail of the magnetosphere, we will derive an empirical calibration formula for the dawn-to-dusk electric fields (Ey). DATA SET Data obtained by the low-energy particle experiment (LEP) [Mukai et al., 1994], the magnetic field experiment (MGF) [Kokubun et al., 1994] and the electric field detector (EFD) [Tsuruda et al., 1994] onboard the GEOTAIL spacecraft are used in the present study. The three-dimensional velocity distribution functions of ions and electrons measured by LEP are used to derive their velocity moments every four-spin period (about 12s). The MGF can measure magnetic field vectors at a rate of 16 samples per second, but we use the data averaged over four spin periods in the present study. The double probes in the EFD-P are placed in the spin plane at the tips of two thin wires, each having the length of 50m from the spacecraft sidewall. By the double probes two components, i?x and Ey. of the DC electric fields in the spacecraft coordinate system are available every spin period (about 3s), but we use the data averaged over four spin periods. Note that the GEOTAIL spin axis is almost parallel to the 7, axis of the GSE coordinate system, and hence the £?x and Ey are almost identical to the sunward and the duskward components, respectively, -79-
Fig 1. Plasma properties of the selected data for four regions.
Table 1. The definition for the region identification in the present study.
Region in the distant tail Plasma Sheet (PS) Plasma Sheet Boundary Layer (PSBL) Lobe/Mantle (L/M)* Cold Dense Plasmas (CDPs)
A > l 0.05 < ft < 1 <1 >1
Ti(eV) >300 >300 <300 <300
*Note that the criterion of Lobe/Mantle contains |B| > 5nT.
in the GSE coordinates. Also we use the data of the spacecraft potential averaged over four period periods. The spacecraft potential is a function of the ambient plasma density and increases as the plasma density becomes lower in the tenuous plasma. We select data from September 1993 to November 1994 to cover the distant magnetotail region (Xgsm < -50i?e) with following criteria: (1) 0.01 < N; < 2 [/cc], (2) T; > 50[eV], (3) 0.9 < Ne/N[ < 1.1, where A/j, Ne and T\ represent the ion density, the electron density and the ion temperature, respectively. Data for time periods with active control of the spacecraft potential are rejected so that the natural potential of the spacecraft can be used for the analysis. The first and second criteria are used for selection of reliable ion moments: The first criterion is to assure enough counting statistics but without saturation of counts for a large geometrical factor of LEP-EA ion analyzer, while the second one is for using reliable ion temperature, taking into account the lowest energy (32eV) of the ion measurement [Mukai et al., 1994]. The third one is a crude criterion for reliability of the electron moments. It is noted that the ion and electron moments as well as the EFD-P data in the near-Earth tail region seem to have a different properties from those in the distant tail, but it is beyond the scope of this paper. As shown in Figure 1, the selected data are sorted into four regions with ion beta and ion temperature; plasma sheet (PS), plasma sheet boundary layer (PSBL), lobe/mantle (L/M), cold dense plasmas (CDPs) which correspond to the magnetosheath region or Low Latitude Boundary Layer (LLBL). Table 1 shows criteria for the region identification. RAW DATA OF DOUBLE-PROBE ELECTRIC FIELDS IN THE DISTANT TAIL Figure 2 shows histograms of the DC electric fields observed by the EFD-P onboard GEOTAIL. The left and right panels show the occurrences of EY and Ey components of the double-probe electric fields, respectively. AVE, MED and DEV represent the average value, the median value and the standard deviation for each component, respectively. The average value of raw data of Ex is 1.22 mV/m. It is not the real electric field in the ambient plasma, but may be generated by photoelectrons emitted from the surface of -80-
Fig. 2. Histograms of the double-probe electric fields in the distant tail. The right panel and the left panel show the occurrence of the x component and the y component of the double-probe electric fields, respectively.
the GEOTAIL satellite due to the solar EUV radiation. The average value of raw data of Ey is -0.47 mV/m, but the y component of natural convection electric field must be positive because of the existence of the dawn-dusk electric field in the distant tail [Nishida et al.. 1998]. This means that raw data of Ey contain the negative offset.
ANALYSIS METHOD If the frozen-in theory holds in the region of our interest (distant magnetotail), the relationship between the DC electric field (E) and the convection electric field is expressed by the following formula:
E + V; x B = 0,
(1)
where the convection electric field is represented by -V; x B. and V; and B are the ion bulk velocity and the magnetic field, respectively. The correlation between the convection electric fields and the observed double-probe electric fields for the y component is examined for the whole data set (not shown), and the regression line is expressed by Ey = 0.921(-V; x B) y -0.547. The correlation coefficient is 0.932. This result suggests that the observed double-probe electric fields contain errors because the slope of the regression line is not unity and the negative offset exists. The correlation coefficient, however, is close to unity, so we believe that the frozen-in theory holds in most of the distant tail. The sensitivity and the offset are evaluated with comparison between the double-probe electric field and the convection electric field.
RESULTS We have made the correlation analysis of the convection electric field and the double-probe data to study how the correlation coefficient can be improved by separating the data with a limited range of plasma parameters in various regions (Table 1). As a result, we have found that spacecraft potential (Vg/c) is a key factor to control the sensitivity, equivalently the effective length of the double probes. Figure 3 shows an example of the correlation analysis for Ey data at the spacecraft potential (Vg/C) of 15V to 20V in the plasma sheet. The regression line is Ey = 0.837(-Vj x B) v - 0.545.
(2)
The unit is mV/m. The correlation coefficient R is 0.960 and the number of events is 3133. Using the result, the sensitivity and the offset are estimated to be 83.7% and -0.545mV/m, respectively. Figure 4 shows the dependences of the sensitivity and the offset on the spacecraft potential and the correlation coefficient. The panel (d) shows that the sensitivity of the Ey component has a systematic correlation with the spacecraft potential and the correlation is independent of the region. The panel (e) -81-
Fig. 3. An example of regression analysis between the double-probe electric fields and the convection electric fields for the plasma sheet in the spacecraft potential range of 15 to 20 V.
shows that the Ey offset is almost a constant value but has a little different trend in the lobe/mantle region. It is noted that the sensitivity and the offset for the spacecraft potential less than 5V show different tendencies (not shown), and it is probably caused by a different property of the photoelectron sheath around the GEOTAIL spacecraft [Nakagawa et al., 2000]. The panel (f) shows the correlation coefficient of the regression analysis. The correlation coefficients are highest (> 0.95) in the plasma sheet, while it is still above 0.9 in other regions except for the spacecraft potential of 27.5V (25-30V) in the lobe/mantle region. Since all the ion species are assumed to be protons in the calculation of the ion velocity moments, non-negligible contribution of heavy ions, for example O + , may also make the correlation coefficient worse in the lobe/mantle region. Note that oxygen ions tend to be more abundant in the mantle region rather than in the empty lobe [Seki et al., 1998]. We have also found that the offset can be well represented by a function of the electron temperature (Figure 5). Based on these results, an empirical calibration formula for the GEOTAIL double-probe electric fields, Ec. is derived as follows: Ec = {Ey + 0.0393 log Te + 0.557) /S
(3)
5 = 0.0138Vs/c +0.59
(4)
where S. V$/Q and Te represent the sensitivity of the double probe, the spacecraft potential and the electron temperature in the unit of keV, respectively. It is often said that the double-probe technique is difficult to measure the electric field in tenuous plasmas, in which the Debye length is much longer than the separation distance between the probes. The spacecraft photoelectrons may also disturb the natural electric field in the ambient plasmas. The spacecraft potential in the tenuous plasma depends on the current balance between the ambient electrons into the spacecraft and the electrons emitted from sunlit surface of the spacecraft [Garret, 1981]. In the case of the high spacecraft potential, emitted low-energy photoelectrons are hard to escape from the spacecraft, and confined to the vicinity of the spacecraft. Therefore the photoelectron cloud may not make an impact on the probes, which are beyond 50m from the spacecraft. We have also made similar analyses for the Z?x component. However, as shown in the panel (a), (b) and (c) of Figure 4. the trends for the £ x component are quite different from those for the Ey, and we could not derive the reliable calibration formula for the Ex component. Probably the £ x component is strongly -82-
Fig. 4. The dependences of the sensitivity and the offset on the spacecraft potential, and the correlation coefficient. Panel (a), (b), and (c) represent the sensitivity and offset of x component of double-probe electric fields, and correlation coefficient of the regression analysis, respectively. Panel (d), (e), and (f) are shown in same format for y component of double-probe electric fields.
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Fig. 5. Relationship between the offset and the electron temperature.
affected by the asymmetrical distribution of the photoelectron cloud around the spacecraft, and the study should be pursued in the future. SUMMARY We have derived the empirical formula to correct the y component of the double-probe electric fields observed by GEOTAIL. The sensitivity increases as the spacecraft potential becomes higher, while the offset weakly depends on the electron temperature. The good correlation coefficient means that the doubleprobe technique is useful even for very tenuous plasmas, and the magnetotail plasmas are frozen-in in most cases. ACKNOWLEDGMENTS The authors thank T. Terasawa for useful comments. REFERENCES Garret, H. B., The charging of spacecraft surfaces, Rev. Geophys. Space Phy., 19, 577, 1981. Kokubun, S., T. Yamamoto, M. H. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46, 7, 1994. Mozer, F. S., E. W. Hones, Jr., and J. Birn, Comparison of spherical double probe electric field measurements with plasma bulk flows in plasmas having densities less than 1 cm""3, Geophysical Res. Lett., 10, 8, 1983. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particle (LEP) experiment onboard the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 669, 1994. Nakagawa, T., T. Ishii, K. Tsuruda, H. Hayakawa, and T. Mukai, Net current density of photoelectrons emitted from the surface of the GEOTAIL spacecraft, Earth Planets Space, 52, 283, 2000. Nishida, A., T. Mukai, T. Yamamoto, S. Kokubun, K. Maezawa, A unified model of the magnetotail convection in geomagnetically quiet and active times, J. Geophys. Res., 103, 4409, 1998. Seki, K., M. Hirahara, T. Terasawa, T. Mukai, Y. Saito, S. Machida, T. Yamamoto, and S. Kokubun, Statistical properties and possible supply mechanisms of tailward cold O+ beams in the lobe/mantle regions, J. Geophys. Res., 103, 4477, 1998. Pedersen, A., F. Mozer, and G. Gustafsson, Electric field measurements in a tenuous plasma with spherical double probes. Geophysical monograph, 103, 1. 1998. Tsuruda, K., H. Hayakawa, M. Nakanura, T. Okada, A. Matsuoka, F. S. Mozer, and R. Schmidt, Electric field measurements on the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 963, 1994. E-mail address of Y. Takei: [email protected]
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SECTION 2: Current Sheet and Magnetic Reconnection
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MAGNETIC RECONNECTION IN THE MAGNETOTAIL: REVIEW OF THE JAPANESE CONTRIBUTION WITH THE SPACECRAFT GEOTAIL T. NAGAI1 'Tokyo Institute of Technology, Tokyo 152-8551, Japan
ABSTRACT This paper reviews the Japanese contribution to our current understanding of magnetic reconnection with the spacecraft Geotail. Past spacecraft observations have examined mainly the MHD characteristics of magnetic reconnection in the Earth's magnetotail, whereas Geotail has revealed various characteristics of magnetic reconnection at the kinetic level as well as the MHD level. At the MHD level, the role of magnetic reconnection in substorm physics is studied. It is firmly established that magnetic reconnection for substorm onsets takes place in the spatially limited region in the premidnight magnetotail at radial distances of 20-30 RE. Geotail observations provide strong evidence that reconnection observations in the magnetotail precede common substorm indicators on the ground. At the kinetic level, ion dynamics in magnetic reconnection is examined, and Hall physics (ionelectron decoupling) is confirmed first in the vicinity of the magnetic reconnection site.
SIGNATURES OF MAGNETIC RECONNECTION Magnetic reconnection is defined as the physical processes in which magnetic field energy is converted into particle kinetic and thermal energies in association with a change in magnetic field topology. In past spacecraft observations, fast tailward flows with negative Bz are attributed to the occurrence of magnetic reconnection in the magnetotail (e.g., Hones, 1979, 1980; Nishida, 1984). These fast tailward flows are expected to form plasmoids in the distant tail, and plasmoids are observed in the distant tail beyond 100 RE (e.g., Hones et a l , 1984). Plasmoids are also observed as traveling compression regions (TCRs) in the magnetic field in the tail lobes (e.g., Slavin et al., 1984). It is indeed difficult to identify any change in magnetic field topology with a single point observation. However, it is known from spacecraft observations that Bz regularly becomes negative in the midtail equatorial plane in association with fast tailward flows during substorms. This signature is strong evidence of magnetic reconnection, and we believe it excludes other explanations (e.g., Nishida, 1984). Observations of plasmoids strongly support our belief. This interpretation is based on our knowledge of magnetic reconnection at the MHD level. Various simulation studies of magnetic reconnection have been carried out using an MHD code (e.g., Ugai and Tsuda, 1977; Hayashi and Sato, 1978). Several distinguishing characteristics of plasmas and fields have been found in the MHD simulation studies (e.g., Sato, 1979; Sato et al., 1982), and these should be compared with in situ observations. These characteristics are as follows: (1) Fast outward flows show a pure convection motion in the equatorial plane. (2) Fast outward flows tend to be field-aligned off the equatorial plane. (3) Near the separatrix layer, slow-mode shock forms and ions are accelerated and heated. (4) Electrons are accelerated near the X-type neutral line. GEOTAIL OBSERVATIONS AND SIMULATIONS The spacecraft Geotail was launched on July 24, 1992 (Nishida, 1994). Geotail surveyed mostly the distant tail in 1992-1994. Geotail changed its apogee near 50 RE in 1994 and near 30 RE in 1995. Geotail has surveyed the midtail at radial distances of 20-30 RE over seven years. Geotail has added significantly to our knowledge of
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magnetic reconnection, not only at the MHD level but also at the kinetic level. Nishida (2000) summarized the overall dynamics of the magnetotail. Various new characteristics of magnetic reconnection at the kinetic level are revealed by recent simulation studies using a hybrid code and a full-particle code (e.g., Hesse and Winske, 1994, 1998; Mandt et al, 1994; Krauss-Varban and Omidi, 1995; Tanaka, 1995; Lin and Swift, 1996; Horiuchi and Sato, 1997; Nakabayashi and Machida, 1997; Hoshino et al., 1998; Nakamura and Fujimoto, 1998; Nakamura et al., 1998; Shay et al., 1998, 2001; Hesse et al., 1999, 2001a, 2001b; Yamade et al., 2000; Pritchett, 2001). In the magnetic reconnection site, the ideal MHD condition breaks down, and the scale length there becomes smaller than ion inertial length and electron inertial length. Ion inertial length, typically 600 km in the midtail plasma sheet, is approximately 40 times longer than electron inertial length. Ions are easily unmagnetized, even when electrons are still coupled to the field. An ion diffusion region forms, and ion-electron decoupling takes place there. In the inflow region, electrons can be transported with the magnetic field lines very close to the electron diffusion region, well inside the ion diffusion region, whereas ions easily escape from the magnetic field lines and do not approach the electron diffusion region. This relative motion produces currents, namely the Hall current system (Sonnerup, 1979). Furthermore, the relative motion of electrons and ions is expected to form the electric field, whose direction is toward the neutral sheet in the inflow region. The electric field Ez, which is another good indicator of magnetic reconnection, is expected to produce a dawnward convection motion for incoming magnetized plasmas in both the northern and southern tail lobes. In the MHD simulations, the speed of outflows is considered to become the tail lobe Alfven velocity (typically 3000 km/s in the midtail). In the full-particle simulations (e.g., Hesse et al., 1999), electron outflow velocity exceeds the Alfven velocity and is higher than ion outflow velocity in the immediate vicinity of the electron diffusion region. This ion-electron decoupling produces the electric field Ex. These high-speed electron outflows are decelerated, and then the electron outflow velocity becomes comparable to the ion outflow velocity. This deceleration process is expected to produce a large Bz field. It is important to note that only a weak Bz field forms in the outflow region at the MHD level. Hence, a large Bz field is another good indicator of magnetic reconnection at the kinetic level. IN SITU OBSERVATIONS OF MAGNETIC RECONNECTION It is not evident how to find a position where magnetic reconnection takes place with a single spacecraft observation. To find the magnetic reconnection site, we can utilize a so-called velocity filter effect. For a plasmoid in the distant tail (e.g., Machida et al., 1994a), high-energy electrons are first observed, and observed energies of electrons decrease. Then, high-energy ions are observed. Finally a plasmoid itself, which contains low-energy plasmas with a bipolar Bz signature, is observed. This velocity filter effect indicates that magnetic reconnection for substorm onsets starts rather abruptly in a spatially limited region. In the case of high-energy flowing electrons and high-energy flowing ions being observed simultaneously, the observations are thought to be carried out in the immediate vicinity of the magnetic reconnection site. Indeed, Geotail has found several events in which electrons are strongly accelerated and heated and coexist with high-energy ions. Strong electron heating is seen in previous spacecraft observations (Bieber et al., 1984). The characteristics observed in plasmas and fields are consistent with several expectations in the simulation studies, and they provide new insight into magnetic reconnection. We have studied the plasma and field characteristics in the immediate vicinity of the magnetic reconnection site in detail using the data obtained in the period 1400-1410 UT on January 27, 1996, with Geotail (Nagai et al., 1998a; 2001). Geotail is located at (-28.9, 5.7, -2.5 RE) in the GSM coordinates. Ion and electron velocity distribution functions reveal ion and electron behaviors at the kinetic level. Ion and electron distribution functions in three representative time intervals are shown in Figure 1. The distribution functions are presented in the plane that includes the magnetic field vector and the ion flow vector. The B (vertical) axis is the magnetic field direction, and the C (horizontal) axis is the flow direction perpendicular to the magnetic field. Ion phase space densities from 10"170 to 10"135 s3m~6 and electron phase space densities from 10"220 tolO"185 s3m"6 are color-coded according to the color bar at the right. The magnetic field observations (with a time resolution of 1/16 s) in the GSM coordinate system are shown in Figure 2. The near-equatorial plane is defined as the region where Bx is almost zero and Bz is near -10 nT (12-s time interval A in Figure 2). At 1404:59 UT, the positive B direction is almost southward and the positive C direction is almost tailward. The ion distribution (Figure la) consists of two hot-ion components with a peak at Vperp = 2500 km/s (tailward). This indicates a bulk tailward velocity of 2500 km/s. The ions are counterstreaming parallel
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Fig. 1. Ion and electron distribution functions at 1404:59 UT (time interval A, a and d), 1404:34 UT (time interval B, b and e), and 1406:12 UT (time interval C, c and f) on January 27, 1996, from Geotail.
January 27, 1996 Fig. 2. Magnetic field variations (with a time resolution of 1/16 s) in the GSM coordinate system for the period 1400-1410 UT on January 27, 1996, from Geotail. The 12-s time intervals starting at 1404:59 UT (A), 1404:34 UT (B), and 1406:12 UT (C) are indicated.
to the magnetic field (almost perpendicular to the equatorial plane) with a speed of 800-1200 km/s. High-speed counterstreammg ions are particles that just come from the north and south tail lobes and make a meandering motion in the magnetic reconnection site. Electrons (Figure 1 d) have a rather flat-top distribution, indicating strong acceleration and heating (see also Shinohara et al., 1998). The center of the distribution is shifted tailward, and this shift corresponds to an electron convection velocity of approximately 4000 km/s (tailward), which is higher than the estimated Alfven velocity. These characteristics are consistent with the results of the simulations (e.g., Nakamura et al., 1998; Hoshino, 1998; Hesse et al., 1999). The off-equatorial plane is defined as the region where Bx is less than 10 nT and Bz is near -5 nT (time interval B in Figure 2). At 1404:34 UT, the positive B direction is almost earthward and the positive C direction is
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toward the equatorial plane. Ions (Figure lb) have a field-aligned component and a convection component. The field-aligned component shows fast tailward flows almost along the field line, and its velocity exceeds 2800 km/s (the upper limit of the instrument). The convection component consists of warm ions, and the convection velocity is in the wide range of 500-1500 km/s toward the equatorial plane. The Bx magnitude is regarded as the distance from the equatorial plane. Just off the equatorial plane, the field-aligned flow velocity is high, the convection velocity is high, and the convection component has a higher temperature. Far from the equatorial plane, the fieldaligned flow velocity is still high, the convection velocity is lower, and the convection component has a lower temperature. These characteristics indicate ongoing acceleration and heating processes for ions. There are cases in which only part of Larmor motion is observed for ions (e.g., Nagai et al., 1998b). The electron distribution (Figure 1 e) is an almost flat-top distribution with an excess of higher-energy electrons in the field direction. This excess means tailward escaping of high-energy electrons. The electrons are more energetic off the equatorial plane than those in the equatorial plane. The ion-electron decoupling is most evident in the plasma sheet/tail lobe boundary (Nagai et al., 2001). The time interval C in Figure 2 is taken in the plasma sheet/tail lobe boundary. At 1406:12 UT, the positive B direction is almost earthward and the positive C direction is toward the equatorial plane. Ions (Figure lc) have two components: a field-aligned component and a convection component. The field-aligned component has a high tailward velocity. The electron distribution (Figure If) is elongated along the field line, and the excess of the highenergy electrons is more evident (the lower part of Figure If). Hence, ions and electrons make field-aligned highspeed outflows. A prominent feature in the electron distribution is a low-energy (3-keV) electron beam in the direction opposite to that of the ion outflows (the upper part of Figure If)- These low-energy electrons form fieldaligned inflows toward the magnetic reconnection site. The inflowing electrons are observed in the four quadrants in "the 2D magnetic reconnection plane" (Nagai et al., 2001). The low-energy electrons can become current carriers. Indeed, the current density is calculated to be 6-13 nA m"2, and the currents flow out of the magnetic reconnection region. In the original idea of the Hall current system (Sonnerup, 1979), outward currents are almost perpendicular to the field lines in the inflow region, and there are inward currents in the outflow region in order to keep the current continuity. The observed currents are almost parallel to the magnetic field (Figure 3). In the simulations, outward currents are confined to the plasma sheet/tail lobe boundary (e.g., Nakamura et al., 1998), and counterstreaming features are reproduced in electron distributions (Hoshino et al., 2001). Since electrons are fairly mobile in the field direction, it is likely that field-aligned currents are much more evident in the inflow region. However, it is not evident how current-carrying electrons have energies of a few keV. The electrons seem to be accelerated with parallel electric fields. As indicated in the simulations, Ez toward the equatorial plane forms in the inflow region, and tailward/earthward Ex forms in the outflow region. The potential structure results in an outer field-aligned electric field in the separatrix layer. This field-aligned electric field probably accelerates electrons up to a few keV (Nagai et al., 2001).
Fig. 3. Summary of Geotail observations near the magnetic reconnection site. The upper right shows ion behavior, while the upper left shows electron behavior. The lower panel shows the Hall currents, associated By deflections, and the Ez electric field.
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The effect of the Hall current system should be observed as the By deflection in the magnetic field (Sonnerup, 1979). The four-current loop of the Hall current system would result in a quadrupole structure in the By field in the vicinity of the magnetic reconnection site. Indeed, the consistent By deflection is found in the observations (Nagai et al., 2001). Thus, the existence of the Hall current system is firmly established by the direct detection of the currents and the detection of the magnetic field deflection as its effect. The Hall current signatures are confirmed in the magnetotail at radial distances of 60 RE with the spacecraft Wind; however, the Wind observations seem not to be associated with substorm activity (0ieroset et al., 2001). The extension of the Hall current system is observed in the near-Earth tail (Fujimoto et al., 2001; Ueno et al., 2002) and in the midtail (Fujimoto et al., 1997). The Ez field, as indicated in Figure 3, is observed as ion flows in the inflow region (Nagai et al., 2001). Far from the separatrix layer, incoming ions are cold and show a convection motion toward the equatorial plane. Near the separatix layer, incoming ions become hot and persistently show a dawnward motion. This can be interpreted as the ExB drift of incoming ions which are partially magnetized. Hence, this observation provides good evidence of the electric field Ez produced by ion-electron decoupling. Slow shocks are expected to form in the plasma sheet/tail lobe boundary in the course of magnetic reconnection. Slow shocks are identified in the distant magnetotail (Saito et al., 1995, 1996, 1998; Hoshino et al., 1997, 2000). However, they are not easily identified near the magnetic reconnection events in the near-Earth plasma sheet. Identification of slow shocks requires observations in both the upstream and downstream regions. It is possible that any simple crossing of the plasma sheet/tail lobe boundary is rare for observations in the near-Earth magnetotail. It is also possible that shock conditions are not easily satisfied with significant temporal variations in the near-Earth magnetotail. It should be noted, however, that ion heating is always observed near the plasma sheet/tail lobe boundary in the reconnection events (e.g., Nagai et al., 1998a; Ueno et al., 2001). Just before an onset of magnetic reconnection, strong intensification of the cross-tail current is observed in the near-Earth plasma sheet (Mukai et al., 2000). The intensification of the current is caused mainly by enhancement of the dawnward electron velocity, and fluctuations start in the magnetic field. It has been suggested that magnetic reconnection takes place in the fluctuating thin current sheet (Shinohara et al., 2001); however, any definite onset mechanism of magnetic reconnection is not revealed observationally. COUNTERSTREAMING IONS AS EVIDENCE OF MAGNETIC RECONNECTION Counterstreaming ions are observed in tailward plasma flows, inside plasmoids, and even at the head of fast earthward plasma flows. Counterstreaming ions form in the leading edge of jetting plasmas produced with magnetic reconnection, where the magnetic field lines pile up because of pre-existing stationary plasmas. The counterstreaming ions originate from cold ions in northern and southern tail lobe fields, and the transport of the reconnected field lines makes these cold ions flow into the equatorial plane of the magnetotail (e.g., Nakamura et al., 1998). These counterstreaming ions are kept in plasmoids (Mukai et al., 1996, 1998; Hoshino, 1998; Hoshino et al., 1998). Strong southward Bz is observed at the head of fast tailward flows (Nagai et al., 1998c), and counterstreaming ions are most evident in association with the large Bz (Fujimoto et al., 1996). It is noted that counterstreaming ions are rather hot in plasmoids and tailward flows. This fact may indicate that these ions are heated by slow shock produced by magnetic reconnection in the midtail. Cold counterstreaming ions are found at the front of fast earthward plasma flows in the recovery phase of substorms (Nagai et al., 2002). The counterstreaming ions for earthward flows are confined to the equatorial plane when Bz is strongly northward. The lower temperature of the counterstreaming ions in the earthward flows may indicate that slow shock does not develop strongly in magnetic reconnection in the distant tail. RELATIONSHIP OF MAGNETIC RECONNECTION TO SUBSTORMS Geotail has provided various characteristics of magnetic reconnection at the MHD level. Analyses of flow directions and fields associated with substorm onsets have shown that strong coupling between fast tailward (earthward) flows and negative (positive) Bz is persistent and that magnetic reconnection frequently takes place in the premidnight region at radial distances of 20-30 RE in association with substorm onsets (Nagai et al., 1998a, 1998d, 1998e; Nagai and Machida, 1998). An important finding is that the magnetic reconnection site is initially a spatially limited region, and then it expands both duskward and dawnward. However, the magnetic reconnection site is spatially limited in the magnetotail, and it rarely expands to the dawn or dusk tail boundary. The same results have been obtained from several analyses, using different procedures to determine substorm onsets and plasma and field signatures (Nakamura, 1999; Machida et al., 1999, 2000; Miyashita et al., 1999, 2000, 2001, 2002; Ueno e t a l , 1999).
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The start timing of magnetic reconnection relative to substorm onset signatures on the ground has been studied extensively (Table 1). The Geotail observations have provided strong evidence that magnetic reconnection in the magnetotail precedes substorm onset signatures on the ground. Since the magnetic reconnection site is spatially limited to the premidnight sector of the magnetotail, only a signature of magnetic reconnection observed in the premidnight sector should be compared with substorm onset signatures on the ground. The tailward flows with negative Bz in the plasma sheet are the most important signature of magnetic reconnection. The start time of this signature has been compared with the onset time of midlatitude Pi2 pulsations (Nagai et al., 1998a). When Geotail is in the premidnight sector beyond 20 RE, the tailward flows are observed consistently 2-3 minutes prior to the Pi2 onset. When Geotail is located in other regions of the magnetotail, the tailward flows are observed at or even after the Pi2 onset. A complementary study with other data sets (Nagai et al., 1998e) confirms that magnetic reconnection observations in the magnetotail precede substorm indicators on the ground. This indicate that magnetic reconnection starts in a very limited region in the premidnight sector of the magnetotail beyond 20 RE, and that the time delay should be interpreted as the propagation time of the flow in the radial direction or the expansion time of the magnetic reconnection site in the dawn-dusk direction. The fast earthward flows in the premidnight sector also precede the ground onset (Nagai et al., 1998a; Nagai and Machida, 1998). This result has been confirmed by subsequent studies (e.g., Miyashita et al., 1999; Machida et al., 1999, 2000). A clear observation of magnetic reconnection, even at X = -17 RE, also precedes the ground substorm onset determined with the Pi2 onset and the auroral kilometric radiation (AKR) onset by 5 minutes (Mukai et al., 2000). The TCR is a good indicator of the development of magnetic reconnection that is observed in the tail lobes. The start time of TCR observations precedes particle injections at geosynchronous orbit by 1-4 minutes (Taguchi et al., 2000; see also Taguchi et al., 1998, 2001). Table 1. Tail magnetic reconnection and substorm onset Tail reconnection signature Tailward flows with Bz < 0 and earthward flows with Bz > 0
Substorm onset signature Pi2
Tailward flows with Bz < 0 Traveling compression region Plasmoid Plasma and field changes
AKR Particle injection Auroral brightening Auroral breakup
References Nagai et al., 1998a, 1998e Nagai and Machida, 1998 Machida et al., 1999, 2000 Miyashita et al., 1999, 2000, 2001 Mukai et al., 2000 Taguchi et al., 2000 Ieda et al., 2001 Miyashita et al., 2002
Recently, the onset timing was determined with auroral observations from spacecraft. Figure 4 shows a representative example taken on January 12, 1997. Geotail was located at (-28.7, +6.3, -2.9 RE). Total pressure (plasma pressure + magnetic pressure) was almost equal to the plasma pressure (the bottom panel of Figure 4) for the period 0720-0725 UT, indicating that Geotail was located near the neutral sheet (in the central plasma sheet). Hence, Geotail is at the ideal place for observations of magnetic reconnection: near the equatorial plane in the premidnight sector of the magnetotail beyond 20 RE. A well-isolated substorm started after a prolonged quiet period. The Polar Ultraviolet Imager (UVI) observations show that the first sign of a substorm onset appeared in the 23 MLT region at 0724 UT (Brittnacher et al., 1999). The Visible Imaging System (VIS) on Polar also showed a substorm onset at 0726 UT (Frank et al., 1998). The timings of the observed auroral images from the Polar UVI are shown by vertical dashed lines in Figure 4 (courtesy of J. Spann). The exposure time is 18 s or 36 s. The observations starting at 0723:03 UT (time 3 in Figure 4) does not show any brightening in the premidnight sector. The observation starting at 0724:35 UT (time 4 in Figure 4) shows a tiny brightening near 23 MLT. The observations starting at 0724:54 UT (time 5 in Figure 4) shows a clear brightening, which covers the 22.5-23.5 MLT sector. The ground magnetic field data from the CANOPUS chain (courtesy of G. Rostoker) shows a start time of the westward electrojet and that of Pi2 pulsation at 0726 UT at Fort Smith (magnetic latitude of 67°, 23 MLT). The photometer data from Fort Smith (courtesy of G. Rostoker) shows a weak brightening at 0725 UT (time resolution of 1 minute). A dipolarization in the magnetic field at geosynchronous spacecraft GOES 8 (02.5 MLT) and GOES 9 (22.5 MLT) was delayed relative to the Polar auroral breakup (Lu et al., 1999). These
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Fig. 4. Geotail magnetic field and plasma data for the period 0710-0740 UT on January 12, 1997. The start time of each Polar UVI auroral image is indicated by a vertical dashed line. observations indicate that the auroral breakup took place after 0723:39 UT (the end time of the auroral image observation starting at 0723:03 UT), and likely not prior to 0724:00 UT (a brightening seems to start just before or during the observations starting at 0724:35 UT). For this event, the tailward plasma flows gradually started at 0723:06 UT (the start time of 12-s plasma sampling) prior to any ground onset time, and the flows were continuing at 0724:00 UT (the fifth panel of Figure 4). The tailward plasma flows are usually observed as flows with a slow speed and northward Bz, and then as flows with a high speed and southward Bz. The slow flows with northward Bz are pre-existing plasmas pushed by subsequent plasma jets, which are created with magnetic reconnection. Hence, this is unambiguous evidence that magnetic reconnection in the magnetotail precedes any substorm signatures on the ground. Observations of plasmoids in the magnetotail are compared with auroral observations observed by the Polar UVI (Ieda et al., 2001), and the plasmoid observations are found to precede auroral brightening by about 2 minutes when Geotail is located near the auroral brightening meridian. A statistical study using 402 auroral breakups from Polar UVI has been carried out (Miyashita et al., 2002), and the results obtained using the auroral breakups are the same as those using Pi2 onsets (Miyashita et al., 1999, 2000, 2001). Hence, the results using Pi2 onsets are
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confirmed. Furthermore, it is well established that auroral activity is closely coupled to tail flow activity in the same meridian (Fairfield et al., 1999; Nakamura et al., 2001a, 2001b; Miyashita et al., 2003; Ieda et al., 2003). For multiple-onset substorms, a tailward plasma flow burst with a negative Bz field is observed for each onset beyond 20 Re (Nagai et al., 1998a; Miyashita et al., 2003). It is difficult to determine whether or not magnetic reconnection in association with onsets continues during the expansion phase of substorms. The expansion phase usually appears to continue for 30-40 minutes on the ground; however, it is difficult to precisely determine the end of the expansion phase. Plasma flows in the midtail seem to end prior to the end of the expansion phase. Earthward plasma flows appear in the recovery phase of substorms in the midtail. Past studies have suggested that the magnetic reconnection site retreats toward the distant magnetotail, causing earthward plasma flows in the midtail (e.g., Hones, 1977). However, there has been no direct observation of the retreat of the magnetic reconnection site. Geotail frequently continues to stay near the equatorial plane of the plasma sheet in the expansion and recovery phases of substorms. Cold dense plasmas often appear around the end of the expansion phase. These plasmas are thought to be transported from the duskside plasma sheet or the dawnside plasma sheet (Nagai et al., 1998d). Earthward flows often appear intermittently in the recovery phase. Hence, it is possible that a simple tailward retreat of the magnetic reconnection site does not take place. A possible scenario is that magnetic reconnection ceases in the midtail, and then new magnetic reconnection takes place in the distant tail. This scenario should be tested with future observations. Tailward plasma flows evolve into plasmoids in their tailward propagation (Machida et al., 1994b, 2000; Nakamura et al., 1995; Murata et al., 1995; Nagai et al., 1997a; Ieda et al., 1998). Timings of plasmoid observations relative to substorm onsets are progressively delayed in the magnetotail, indicating the continuous tailward propagation (Nagai et al., 1994, 1997a; Nagai and Machida, 1998). Quasi-stagnant plasmoids are also observed (Kawano et al., 1994, 1996). The duration of plasmoids is typically less than a few minutes (e.g., Ieda et al., 1998). However, there are tailward plasma flows that continue more than 2 hours after the passage of a plasmoid (Nagai et al., 1997a). The long-duration tailward flows seem to continue even after the recovery phase of substorms on the ground. It is believed that these tailward flows are associated with a prolong southward IMF Bz condition. Unfortunately, the IMF data are not available for these Geotail observations. This new class of plasma flows should be studied with future observations. New features are also observed in tail lobe ions. Cold ions are significantly heated shortly before and after the passage of a plasmoid (Mukai et al., 1994), and this process can be reproduced by a simple computer model (Tsubouchi and Terasawa, 1995). After the passage of a plasmoid in the distant tail, an increase in density is frequently observed in the tail lobes (Shirai et al., 2001). This process is thought to be caused by the shrinking of the magnetotail. Earthward plasma flows are studied in the near-Earth tail. However, the existence of high-energy particles may affect precise observations of low-energy plasmas, especially correct measurements of MHD parameters (density, temperature, and plasma velocity). High-speed plasma flows with a speed of >2000 km/s can approach even at radial distances of 10 RE in association with substorms (Fairfield et al., 1998). It is well known that a magnetic substorm signature near geosynchronous orbit is a dipolarization in the field configuration (e.g., Nagai, 1982). Dipolarization in the field is a major signature of substorms, even near radial distances of 10 RE (Fairfield et al., 1998, 1999; Nakamura et al., 1999a; Nagai et al., 2000; Nakai and Kamide, 2000), and dipolarization is usually preceded by earthward plasma flows. It is noted that the magnetic field shows significant fluctuations in the course of dipolarization, and these fluctuations may be related to instability at the head of reconnection jets (Nakamura et al., 2002). In magnetic reconnection, the magnetic field energy stored in the tail lobes in the solar wind/magnetosphere interaction is consumed. The enhancement of the tail lobe energy is observed as an enhancement of the plasma sheet pressure in the growth phase (Nagai et al., 1997b). The enhancement of the plasma pressure in the plasma sheet is caused mainly by an increase in plasma density, and the plasma temperature decreases. The stored energy is thought to be consumed during substorms; however, the energy in the magnetotail in the recovery phase can become smaller than that prior to the growth phase (Nakamura et al., 1999b). The energy storage in the magnetotail is thought to start with a southward turn of the IMF Bz. However, even in the prolonged southward IMF Bz period, substorms take place quasi-periodically, and the increase-decrease sequence in the tail lobe magnetic field is observed for each substorm (Nagai et al., 2003). Hence, the start time of the energy storage can start without any change in the IMF Bz condition. Furthermore, for the prolonged southward IMF Bz, the energy in the magnetotail becomes smaller than that prior to energy storage. These results suggest that magnetic reconnection for substorms is not simply a process that consumes excess energy in the magnetotail.
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CONCLUSIONS All the crucial aspects of magnetic reconnection are confirmed in the in situ Geotail observations at the kinetic level as well as the MHD level. Magnetic reconnection is an indisputable physical process for plasma transport, plasma acceleration and heating, and changes in the magnetic field topology in the magnetotail. Magnetic reconnection is an essential constituent of substorms. Our observations have revealed the ion dynamics and ionelectron decoupling process. Future studies need to focus on resolving electron dynamics in order to expand our understanding of magnetic reconnection.
ACKNOWLEDGMENTS TN thanks I. Shinohara, M. Fujimoto, and R. Nakamura for informative discussions in preparing this paper. The full Polar UVI images were provided by J. Spann (NASA/MSFC), and the CANOPUS magnetic field and photometer data were provided by G. Rostoker (University of Alberta). Y. Miyashita and A. Ieda provided unpublished results on the substorm studies. REFERENCES Bieber, J. W., E. C. Stone, E. W. Hones, Jr., D. N. Baker, S. J. Bame, and R. P. Lepping, Microstructure of magnetic reconnection in Earth's magnetotail, J. Geophys. Res., 89, 6705-6716, 1984. Brittnacher, M., M. Fillingim, G. Parks, G. Germany, and J. Spann, Polar cap area and boundary motion during substorms,/ Geophys. Res., 104, 12251-12262, 1999. Fairfield, D. H., T. Mukai, A. T. Y. Lui, C. A. Cattell, G. D. Reeves, T. Nagai, G. Rostoker, H. J. Singer, M. L. Kaiser, S. Kokubun, A. J. Lazarus, R. P. Lepping, M. Nakamura, J. Y. Steinberg, K. Tsuruda, D. J. Williams, and T. Yamamoto, Geotail observations of substorm onset in the inner magnetotail, J. Geophys. Res., 103, 103-117,1998. Fairfield, D. H., T. Mukai, M. Brittnacher, G. D. Reeves, S. Kokubun, G. K. Parks, T. Nagai, H. Matsumoto, K. Hashimoto, and D. A. Gurnett, Earthward flow bursts in the inner magnetotail and their relation to auroral brightenings, AKR intensifications, geosynchronous particle injections and magnetic activity, J. Geophys. Res., 104, 355-370, 1999. Frank, L. A., J. B. Sigwarth, and W. R. Paterson, High-resolution global images of Earth's auroras during substorms, in International Conference on Substorms-4, edited by S. Kokubun and Y. Kamide, pp. 3-8, Terra Scientific Publishing Company, Tokyo, 1998. Fujimoto, M., M. S. Nakamura, T. Nagai, T. Mukai, T. Yamamoto, and S. Kokubun, New kinetic evidence for the near-Earth reconnection, Geophys. Res. Lett., 23, 2533-2536, 1996. Fujimoto, M., M. S. Nakamura, I. Shinohara, T. Nagai, T. Mukai, Y. Saito, T. Yamamoto, and S. Kokubun, Observations of earthward streaming electrons at the trailing boundary of a plasmoid, Geophys. Res. Lett., 24, 2893-2896, 1997. Fujimoto, M., T. Nagai, N. Yokogawa, Y. Yamade, T. Mukai, Y. Saito, and S. Kokubun, Tailward electrons at the lobe-plasma sheet interface detected upon dipolarizations, J. Geophys. Res., 106, 21255-21262, 2001. Hayashi, T., and T. Sato, Magnetic reconnection: Acceleration, heating, and shock formation, J. Geophys. Res., 83, 217-220, 1978. Hesse, M., and D. Winske, Electron dissipation in collisionless reconnection in current sheets,/ Geophys. Res., 99, 11177-11192, 1994. Hesse, M., and D. Winske, Electron dissipation in collisionless magnetic reconnection, J. Geophys. Res., 103, 26479-26486, 1998. Hesse, M., K. Schindler, J. Bim, and M. Kuznetsova, The diffusion region in collisionless magnetic reconnection, Phys. Plasmas, 6, 1781-1795, 1999. Hesse, M., J. Birn, and M. Kuznetsova, Collisionless magnetic reconnection: Electron processes and transport modeling,/. Geophys. Res., 106, 3721-3735, 2001a Hesse, M., M. Kuznetsova, and J. Birn, Particle-in-cell simulations of three-dimensional collisionless magnetic reconnection,/. Geophys. Res., 106, 29831-29841, 2001b.
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Hones, E. W., Jr., Substorm processes in the magnetotail: Comments on 'On hot tenuous plasmas, fireballs, and boundary layers in the Earth's magnetotail' by L. A. Frank, K. L. Ackerson, and R. P. Lepping, J. Geophys. .Res, 82, 5633-5643, 1977. Hones, E. W., Jr., Transient phenomena in the magnetotail and their relation to substorms, Space Sci. Rev., 23, 393410, 1979. Hone, E. W., Jr., Plasma flow in the magnetotail and its implications for substorm theories, in Dynamics of the Magnetosphere, edited by S.-I. Akasofu, pp. 545-562, D. Reidel, Norwell, Mass., 1980. Hones, E. W., Jr., D. N. Baker, S. J. Bame, W. C. Feldman, J. T. Gosling, D. J. McComas, R. D. Zwickl, J. A. Slavin, E. J. Smith, and B. T. Tsurutani, Structure of the magnetotail at 220 RE and its response to geomagnetic activity, Geophys. Res. Lett., 11, 5-7, 1984. Horiuchi, R., and T. Sato, Particle simulation study of collisionless driven reconnection in a sheared magnetic field, Phys. Plasmas, 4, 277-289, 1997. Hoshino, M., Kinetic ion behavior in magnetic reconnection region, in New Perspectives on the Earth's Magnetotail, Geophys. Monogr. Ser., vol. 105, edited by A. Nishida, D. N. Baker, and S. W. H. Cowley, pp. 153-166, AGU, Washington, D. C , 1998. Hoshino, M., Y. Saito, T. Mukai, A. Nishida, S. Kokubun, and T. Yamamoto, Origin of hot and high speed plasmas in plasma sheet: Plasma acceleration and heating due to slow shocks, Adv. Space Res., 20, 973-982, 1997. Hoshino, M., T. Mukai, T. Yamamoto, and S. Kokubun, Ion dynamics in magnetic reconnection: Comparison between numerical simulation and Geotail observations,/ Geophys. Res., 103, 4509-4530, 1998. Hoshino, M., T. Mukai, I. Shinohara, Y. Saito, S. Kokubun, Slow shock downstream structure in the magnetotail, J. Geophys. Res., 105, 337-347, 2000. Hoshino, M., T. Mukai, T. Terasawa, and I. Shinohara, Suprathermal electron acceleration in magnetic reconnection,/. Geophys. Res., 106, 25979-25997, 2001. Ieda, A., S. Machida, T. Mukai, Y. Saito, T. Yamamoto, A. Nishida, T. Terasawa, and S. Kokubun, Statistical analysis of the plasmoid evolution with Geotail observations, J. Geophys. Res., 103, 4453-4465, 1998. Ieda, A., D. H. Fairfield, T. Mukai, Y. Saito, S. Kokubun, K. Liou, C.-I. Meng, G. K. Parks, and M. J. Brittnacher, Plasmoid ejection and auroral brightenings, J. Geophys. Res., 106, 3845-3857, 2001. Ieda, A., J.-H. Shue, K. Liou, S.-I. Ohtani, C.-I. Meng, D. H. Fairfield, T. Mukai, Y. Saito, S. Machida, T. Nagai, and G. K. Parks, Quiet time magnetotail plasma flow: Coordinated Polar ultraviolet images and Geotail observations, J. Geophys. Res., in press, 2003. Kawano, H., T. Yamamoto, S. Kokubun, K. Tsuruda, A. T. Y. Lui, D. J. Williams, K. Yumoto, H. Hayakawa, M. Nakamura, T. Okada, A. Matsuoka, and K. Shiokawa, A flux rope followed by recurrent encounters with traveling compression regions: Geotail observations, Geophys. Res. Lett., 21, 2891-2894, 1994. Kawano, H., A. Nishida, M. Fujimoto, T. Mukai, S. Kokubun, T. Yamamoto, T. Terasawa, M. Hirahara, Y. Saito, S. Machida, K. Yumoto, H. Matsumoto, and T. Murata, A quasi-stagnant plasmoid observed with Geotail on October 15, 1993, J. Geomag. Geoelectr., 48, 525-539, 1996. Krauss-Varban, D., and N. Omidi, Large-scale hybrid simulations of the magnetotail during reconnection, Geophys. Res., Lett, 22, 3271-3274, 1995. Lin, Y., and D. W. Swift, A two-dimensional hybrid simulation of the magnetotail reconnection layer, J. Geophys. Res., 101, 19859-19870, 1996. Lu, G., N. Tsyganenko, A. T. Y. Lui, H. J. Singer, T. Nagai, and S. Kokubun, Modeling of time-evolving magnetic fields during substorms, J. Geophys. Res., 104, 12327-12337, 1999. Machida, S., T. Mukai, Y. Saito, T. Obara, T. Yamamoto, A. Nishida, M. Hirahara, T. Terasawa, and S. Kokubun, Geotail low energy particle and magnetic field observations of a plasmoid at XGSM = —142 RE, Geophys. Res. Lett., 21, 2995-2998, 1994a. Machida, S., T. Mukai, Y. Saito, M. Hirahara, T. Obara, A. Nishida, T. Terasawa, and K. Maezawa, Plasma distribution functions in the Earth's magnetotail (XGSM ~ -42 RE) at the time of a magnetospheric substorm: GEOTAIL/LEP observation, Geophys. Res. Lett., 21, 1027-1030, 1994b. Machida, S., Y. Miyashita, A. Ieda, A. Nishida, T. Mukai, Y. Saito, and S. Kokubun, Geotail observations of flow velocity and north-south magnetic field variations in the near and middistant tail associated with substorm onsets, Geophys. Res. Lett., 26, 635-638, 1999. Machida, S., A. Ieda, T. Mukai, Y. Saito, and A. Nishida, Statistical visualization of Earth's magnetotail during substorms by means of multidimensional superposed epoch analysis with Geotail data, J. Geophys. Res., 105, 25291-25303,2000.
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E-mail address [email protected]
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BIFURCATED THIN CURRENT SHEETS IN THE EARTH'S MAGNETOSPHERE: COMPARISON OF MODEL AND "INSITU" OBSERVATIONS L. M. Zelenyi1, H. V. Malova2, V. Yu. Popov3, D. C. Delcourt4, and A. S. Sharma5 ' Space Research Institute, RAS, 117810, Profsoyusnaya street 84/32, Moscow, Russia 2 Nuclear Physics Institute, Moscow State University, 119992, Moscow, Russia •'Faculty of Physics, Moscow State University, Vorobyevy gory, 119899, Moscow, Russia 4Centre a"etudes des Environnements Terrestres et Planetaires-CNRS, Saint-Maur des Fausses, France 'Department of Astronomy, University of Maryland, College Park, MD 20742, USA
ABSTRACT A self-consistent analytical theory of thin current sheets (TCSs) is used to investigate their fine structure in the Earth's magnetotail at various stages of temporal evolution. The model is based on the solution of Grad-Shafranov type equations under quasi-adiabatic (QA) approximation. Quasi-adiabaticity allows the construction of an average equilibrium, assuming the conservation of QA invariant of motion 72, and then the investigation of its slower evolution due to Iz - diffusion. This diffusion leads to a gradual trapping of transient and unbounded (or the so called Speiser orbit) particles into orbits trapped in the vicinity of the sheet midplane. It is found that the cross-tail current of such newly trapped population is opposite to the one from transient Speiser orbits and eventually flattens the profile of the magnetic field B near the midplane. As a result profile Bx(z) acquires a complex concave shape instead of a simple linear one, a characteristic of Harris equilibrium. The corresponding TCS cross-tail current profile attains a "double humped" shape, which can be a typical characteristic of TCS during a major part of its "life cycle". This process of current sheet deterioration by quasi-trapped plasma may finally lead to TCS disruption. The results of numerical modeling are compared with Geotail and Cluster observations of double-humped (also referred to as "bifurcated") current sheets in the Earth's magnetotail. The observables predicted by our QA model and the conditions under which they are expected to be observed by Geotail, Cluster and other spacecraft are discussed.
INTRODUCTION During the last decade extremely thin current sheets (TCSs) with thicknesses of about an ion Larmor radius or less have been detected by "in situ" measurements in the Earth's magnetosphere (Mitchell et al., 1990; Pulkkinen et al., 1993, 1994; Sergeev et al., 1993, 1998, 2003; Runov et al., 2003). TCSs have been observed in the near-Earth and distant parts of the magnetotail (Pulkkinen et al., 1993), as well as at the magnetopause (Dunlop et al., 2002). Usually it is assumed that due to the enhanced plasma convection during substorm growth phase the magnetotail current sheet becomes thinner and more stretched. The maximally stretched TCS stores the free electromagnetic energy, which is released during the explosive phase of substorms when some hitherto unidentified plasma instabilities (or other processes) may lead to TCS disruption. The TCSs thus play the key role as the reservoirs of magnetic energy and as the sites of energy transformations during substorms. The general and principal property of TCSs in the magnetotail is their metastability (Galeev and Zelenyi, 1975). Quasi-equilibrium behavior is characteristic for TCS during the growth phase, but subsequently (usually of the order of 1-2 hours) the spacecraft may detect the disruption of TCS (Sergeev et al., 1998, Lui et al., 1992). This may be followed by processes of magnetic field dipolarization and plasma acceleration away from the disruption region. Various plasma instabilities have been proposed as the cause of TCS disruption (Coppi et al., 1966,
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Lembege and Pellat, 1982, Coroniti, 1980, Yoon and Lui, 2001, Silin et al., 2002, Daughton, 1998, Sitnov et al., 1998, 2002) but their roles are currently under debate. Among these one should mention the tearing mode as a "veteran" of these studies (Coppi et al., 1966, Lembege and Pellat, 1982, Sitnov et al., 2002). We would like to look at this problem from a very different perspective and suggest a special "evolutionary" mechanism of TCS disruption not related with any instability. The experimental "hint" is the characteristic "bifurcated", or "double peaked" (or in other words "double humped") structure of TCS current profiles recently found in Geotail and Cluster observations (Hoshino et al., 1996, Runov et al., 2002, Sergeev et al., 2002). This structure is very different from the well-known and mathematically elegant Harris sheet (Harris, 1962), which has a current profile of cosh"2(z/L/) and is most widely used in the interpretation of magnetotail data. Recent findings shed a new light on the observations of a fine structure of TCSs in the magnetotail. Properties of "double humped" current sheets are very different from those of the familiar Harris sheet, but it is not clear what mechanism governs this "double humped" structure. Recently a 3D numerical model was used by Zeiler et al. (2002) to study the stability and structure of double Harris current sheet but neither electron shear flow instability nor kink-like instability was found in these simulations. However, as one might expect for such simulations, the lower hybrid drift instability could develop near the regions of current peaks, although the role of this instability in large-scale reconnection and global current disruption are not clear. From these simulations the authors find that only 2D reconnection is possible in this 3D system. PROPOSED SCENARIO We present a model of a bifurcated current sheet based on a two stage development. Firstly, to describe the plasma equilibrium (or quasi-equilibrium) we need to follow the particle trajectories from their source to the current sheet (CS), and then using Liouville theorem a set of self-consistent equations of Grad-Shafranov type can be derived. To integrate the system of equations controlling particle motion we need to assume the conservation (at least approximately) of the so called CS invariant (Speiser, 1965).
Iz =m(2ny[§vzdz,
(1)
which in fact is an adiabatic invariant arising from the fast motion along the main direction of inhomogeneity (i.e., the z-motion across the CS plasma) (Buchner and Zelenyi, 1986, 1989). Such an approach really enables a description of very thin CS with the thickness of the order of ion Larmor radius or smaller. The separation on a fast and slow motion (and corresponding treatment of Iz as an invariant) is possible when the so called quasiadiabaticity parameter (Buchner and Zelenyi, 1989) K =^/i?c/pmax « 1
(2)
is small (Rc is the minimal curvature radius and pmax is the maximum ion Larmor radius in the current sheet). When the parameter K « 1 , the particle motion is quasi-adiabatic, i.e., the jumps of the adiabatic invariant are smaller than the value of the invariant itself. The general theory of jumps of quasi-adiabatic (QA) invariant has been discussed, for example, by Neishtadt (1987). Later Zelenyi et al. (1990) have shown that the diffusion coefficient D of quasiadiabatic scattering depends on quasi-adiabaticity parameter (2) as D ~ K 2 . The second logical step is to take into account the slower process of the weak but finite jumps of/ z in each separatrix crossing, which results in a gradual particle trapping in the CS. This can significantly modify the structure of current sheets originally formed by transient Speiser orbit ions (Zelenyi et al., 2000). The general effect of the accumulation of quasi-trapped plasma near TCS is two fold: it could appreciably reduce the paramagnetic current density (carried by Speiser ions) near the tail midplane and partially compensate the diamagnetic currents of transient particles at the edges of current sheet (Zelenyi et al., 2000, 2002a). The accumulation of quasi-trapped ions, which we refer to as "TCS aging" (the detailed description of this process is given in papers by Zelenyi et al. (2002a, b), leads also to the flattening of the magnetic field profile in the center of the sheet without a significant change of the TCS thickness. As a result the effective parameter K governing the quasiadiabatic particle dynamics increases, and the rate of ion scattering increases, thus developing a positive feedback loop between the ion scattering and the fine structure of TCS magnetic field. The initially gradual evolution then acquires a catastrophic character and when the resulting quasi-trapped domain becomes highly populated, the current sheet quasiequilibrium decays. This effect provides a new evolutionary mechanism which might explain the disruption of CS structure on a time scale comparable with the duration of the substorm growth phase. Our model, of course, is capable of describing only the approach to this disruption, while fast (or explosive) dynamics or disruption requires much more detailed analysis. -101-
The aim of this article is to consider the self-consistent TCS model (Zelenyi et al., 2002a) taking into account the processes of quasi-adiabatic scattering, investigate the corresponding evolution of the current sheet "deterioration" by quasi-trapped plasma particles and describe the important features of a "middle aged" current sheet which is most likely to be observed by spacecraft crossing the magnetotail mid-plane. BASICS OF THE MODEL Our model of extremely thin current sheets has its roots in a pioneering studies by Speiser (1965) and Eastwood (1972) where the idea of thin anisotropic current sheet supported by ions with strongly "non-adiabatic" trajectories were proposed. Self-consistent model of ion-dominated TCS have been developed later by (Kropotkin and Domrin, 1996, Sitnov et al. 2000, Zelenyi et al., 2000). In this paper we combine the analytical model of thin current sheet (Sitnov et al., 2000, Zelenyi et al, 2000) with the quasi-adiabatic diffusion approach of Zelenyi et al. (1990) to take into account a gradual build-up of the scattered quasi-trapped population of ions in the current sheet. The scheme of the model is shown in Fig.l.
Figure 1. The general scheme characterizing the Figure 2. The profiles of the distribution transient orbits and quasi-trapping near the neutral plane, function of the transient and quasi-trapped plasma within the thin current sheet. The basic assumptions made in developing this QA approach are: 1) current sheet (CS) homogeneity along the Earth-Sun (X-axis) and "dawn-dusk" (7-axis) directions; 2) cross-tail current mainly supported by ions (estimates of the electron contribution have been given by Sitnov et al. (2000) and found to be small); 3) ion population could be subdivided on (a) transient particles having Speiser's orbits, (b) quasi-trapped particles having "cucumber-like" orbits originating from the transient ones due to non-adiabatic scattering. The corresponding quasi-adiabaticity parameter K =A//t/pmax is less than unity for realistic ion parameters in the magnetotail. More detailed description of the corresponding self-consistent Vlasov-Maxwell system of TCS equilibrium in the absence of the scattering process was presented in our earlier publications (Sitnov et al., 2000, Zelenyi et al., 2000). The basic system of equations has the Grad-Shafranov-like form:
db_
4s \VD
2
V)^)+f(-)(C)
(3)
2/2 V
dt n (A l+ «/(s-l) where b(Q=B/Bo is the dimensionless self-consistent magnetic field; Bo is the field at the boundaries of current sheet; C, is the normalized z-coordinate C, = z(£>Q/s4llvD; C00 -eB^/mc; v=V/(E2nVD)is the dimensionless velocity, and the dimensionless parameter s=vj/vD characterizes the particle streaming along the field lines. The right hand side of Eq.(3) signifies the dimensionless ion currents F±(Q, with the signs (+) and (-) corresponding to currents in positive and negative Y-directions. The initial distribution function has the form of a shifted Maxwellian with the normalized plasma density no, average flow velocity VD«VT, where Vj is the thermal velocity taken in the finite interval (0
The detailed discussion of the governing equation, estimates of/* and properties of a resulting distribution function are given by Zelenyi et al. (2002b). Self-consistent system (3) depends on a single parameter z=vjJvD which is assumed to be larger than unity for the calculations below. This is also the case if plasma mantle is the source of ions maintaining TCS. We suppose that the process of quasi-adiabatic relaxation of the system due to scattering processes is much slower than the relaxation of the CS to the quasi-equilibrium states at each instant of CS evolution, therefore one might consider this evolution as a temporal sequence of instantaneous «snapshots» of a set of quasi-equilibria in time. The solutions of Eq. (3) with variable populations of "trapped" ions can be derived from the equations for the temporal evolution of the distribution function of the scattered plasma y(I, t) (where / is the dimensionless adiabatic invariant Iz, t=Q.nz is the dimensionless time, Q.n=eB,Jm,c is the ion gyrofrequency in the normal magnetic field)
^L=AD(I1)^dt di V'ndi
(4) ()
This is a diffusion equation and the diffusion occurs in the interval (lh I2), i.e. across the boundary between Speiser and "cucumber" (quasi-trapped) populations. /; is the boundary between transient and quasi-trapped orbits (/; = v^, where v0 is the total velocity) The boundary 12 (lf=av\, orbits is impenetrable for diffusion:
a=(L/p)m>l)
between quasi-trapped and integrable
d\\i /dt — 0. An approximate expression for quasi-adiabatic diffusion
coefficient found by Biichner and Zelenyi (1989) has the simple form Z>(/,,) = - ^ K
2
(l-/-)
(5)
Here TAB is dimensionless period of large-scale "cucumber" periodic motion (Btichner and Zelenyi, 1989). The profile of the scattered distribution function may change in time due to the accumulation near TCS of particles at quasi-trapped orbits. The corresponding "snapshots" of the scattered distribution function at different times, which were found as solutions of the diffusion Eq. (4), are also shown in Fig.2. NUMERICAL RESULTS After computing the distribution function \[i(I, t) from Eq. (4)-(5) it is straightforward to calculate the
Fig. 3. The accumulation of quasi-trapped (QT) plasma inside TCS. Dimensionless plasma density nQT is a function of dimensionless Z-coordinate Q. At the initial moment of time the QT plasma is supposed to be absent.
Fig. 4. The dimensionless current density profile inside TCS as a function of dimensionless z-coordinate Q for different moments of time.
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dimensionless
plasma
density
j (C,,t ) = \w {f'r""(I(w))+\vi(I(w))\iiw,
nQT(C,
,tt) = (\\i(I(w))cl3w
and
the
current
density
where/™" and i|/, are distribution functions of transient and quasi-
trapped plasma components at time ?,, w - dimensionless particle velocities. The consequent "snapshots" of dimensionless plasma densities nQT(^,ti ) = nQTJn0,
(<; =a»0/e4'Vi), t=Q.nx) as a function of dimensionless z-
coordinate and time t are shown in Fig. 3. Fig. 4 characterizes the accumulation of quasi-trapped plasma within TCS at different times. This picture demonstrates that after some time (t>104) the dimensionless current density jy(C>,t)-J)!(C>,t)/noevTs^'2 in the center of the sheet becomes negative, and the self-consistent solution decays. Particles trapped at "cucumber" orbits undergo a different type of nonlinear meandering motion than transient Speiser orbit particles, as illustrated in Fig. 5. One could see that "trapped" particles in the vicinity of z=0 plane carry currents in a negative Y direction, while at the CS edges their current is positive, i.e. opposite to diamagnetic currents of Speiser particles. The physical process here is similar to the mechanism of current limitation in intense relativistic electron beams in which an increase of the beam intensity or the current leads to a pinch effect and the resulting self-consistent magnetic field affects the electron trajectories, limiting the maximum current that such a beam can carry to the Alfven limit (Benford and Book, 1974).
Fig.5. Elements of quasiadiabatic "cucumber" (solid line) and Speiser (dashed line) orbit in the center of current sheet. The corresponding local current density Jy(z) are shown on the right. The important new feature of the corresponding profiles of magnetic field bfc.jj is a flattening of the profile near z=0 shown in Fig.6. This occurs because the derivative dB(t)/dC, tends to zero as time.advances. Fig. 7a shows the corresponding experimental results by Hoshino et al. (1996) obtained from the analysis of Geotail crossings of the neutral sheet at distances of about 100 RE. Comparing Fig. 4 with Fig.7b one can see the similarity of experimental
Fig. 6. Several profiles of TCS showing the process of the contamination (or "aging") of the sheet by QT plasma. The derivative of the dimensionless magnetic field b(^,t) profile is decreasing with the time near the center plane.
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Fig. 7. The magnetic field (a) and current density (b) profiles as a function of dimensionless zcoordinate, described by Hoshino et al. (1996)
and theoretical profiles of the magnetic fields. It is clear that the "bifurcation" of cross-tail current to "double humped" structure discussed above is well supported by observations. We also investigated contour maps of ion distribution functions in the phase plane {vx,vy} of the "deteriorated" or "old" TCS with a moderate anisotropy s=l. The results are shown in Figure 8, in which six plots, (a)-(f), show the v2=0 cross-sections of the distribution function at distances z/L=].O (edge of TCS); 0.75; 0.5; 0.25; 0.05, and 0.0 (neutral plane), where L is the thickness of the CS. The distribution function in our model is symmetric relative to vx - coordinate; so only the half vx>0 of {vx,vy} plane is shown At the edges of TCS (Fig.8a) the distribution function is a sum of "pure" shifted Maxwellian function (the source distribution) and small additions of quasi-trapped (scattered) plasma reaching the CS edges. One can see how this distribution function becomes more and more distorted when the observer moves towards the CS center. One can see also that the characteristic sharp ridges of the distribution of the scattered plasma and pronounced peaks along the negative direction of v^-axis. The comparison of the details of the model predictions with Cluster and Geotail observations of distribution functions is underway and will be reported in the near future.
Fig.8. Contour plot of cross-section vz=O of ion distribution functions in strongly "deteriorated" TCS (t=100). Cases (a)-(f) correspond to z/L=1.0, 0.75, 0. 5, 0.25, 0.05, and 0.0, respectively. DISCUSSION The present day space instrumentation (multipoint measurements, high temporal resolution, good resolution in velocity space) allows detailed exploration of the fine structure of current sheets in the magnetotail. One of the remarkable findings made even in pre-Cluster era (Sergeev et al., 1993, Hoshino et al., 1996) is that the thickness of CS could be quite small and reach 1-2 RE. In some cases observations of CS with thickness of the order of 0.1-0.2 RE were reported (Asano, 2001 and references therein). The Cluster observations also reveal few cases of very thin CS (Runov et al., 2003), the main finding being that the magnetic field profile in the current sheet could be quite complicated. Contrary to the classical "Harris" like structure (with linear gradients of Bx near current sheet midplane, they could have complicated two-humped shape of cross-tail current, which results in a significant flattening of CS magnetic field profile near neutral plane. In some publications these two-humped current structures are referred to as bifurcated or split (Hoshino et al., 1996, Runov et al., 2003, Sergeev et al., 2003). Hoshino et al., (1996) revealed this feature in statistical analysis of Geotail CS crossing, while Runov et al. (2003) and Sergeev et al. (2003) found this bifurcation in case studies of Cluster crossings. The morphology of these events is not very clear yet. Sergeev et al. (1993) indicated that thin current sheets are more typical for the early stages of substorm growth phase while double CS have been observed closer to substorm onset. There is no generally accepted model of double-peaked current sheet formation. Hoshino et al. (1996) and Asano (2001) elaborated the model where the double sheet is formed due to magnetic reconnection. In their model the electrons accelerated by the induced electric field at the periphery of the X-line are the essential current carriers, and they form the double peaked current distribution. We propose here another model of bifurcated current sheet where the ions are the main current carriers (Zelenyi et al. 2002a). Our self-consistent kinetic model explains the double-peak structure as a result of
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unavoidable quasi-adiabatic scattering of transient plasma particles near the neutral plane due to the jumps of the adiabatic invariant of motion. The scattering process results in the appearance of ions with large "magnetic moments", which may reduce the current at the center of the sheet. This process may be compared with the "aging" of a newly formed current sheet due to its "pollution" by quasi-trapped plasma. When the resulting quasi-trapped domain becomes highly populated and the corresponding double peaked structure becomes strongly pronounced. This effect, according to estimates in Zelenyi et al. (2002a, b), has a time scale comparable with the duration of the substorm growth phase, i.e., of the order of 30-90 min. Therefore we believe that the double-peaked current profiles and corresponding features in ion distribution functions may be considered as a signature of a certain (intermediate) stage of the temporal dynamical evolution of the TCS. In short, our model predicts that spacecraft in the magnetotail should typically expect to encounter a "middle age" current sheet with a characteristic dropout of current density in the center of thin current sheet. Indirectly this is in excellent agreement with the findings by Hoshino et al.(1996) that such double CS structure is statistically significant, i.e., it was detected during a major fraction of the observations. Contrary to the model of Hoshino et al. (1996) and Asano (2001) our model predicts that effects could be observable in ion distribution function which should have pronounced characteristic features in certain domains of a phase space (similar to the so called "anti-loss cone distribution"). Our estimates of the electron contribution (to be published in a separate publication) indicate that it really could be significant either for unrealistic Te/Ti > 1 ratios or for exceedingly small Bn values. So we expect that there is a threshold value Bn where our results for ID current sheet could become compatible with X-line vicinity model of Asano (2001). ACKNOWLEDGMENTS The authors are grateful to A. Nishida, M. Hoshino, V. Sergeev, W. Baumjohann and A.Runov for interesting and productive discussions. This work was supported by the Russian Foundation of Basic Research grants 02-0216003-a, 02-05-64184-a, 01-02-16367, Science School grant HIII-1739.2003.2, RFBR NCNI grant 00-02-220001, PICS APIC0090. The research at the University of Maryland was supported by the NASA grant NAG510298. REFERENCES Asano Y., Configuration of the thin current sheet in substorms, Ph.D. thesis, Univ. Tokyo, 2001. Benford, G., and D. L. Book, Relativistic beam equilibria, in Advances in Plasma Physics ,eds. A. Simon and W. B. Thompson, Vol. 4, pp. 125-172, 1974. Buchner, J., and J.-P. Kuska, Sausage mode instability of thin current sheets as a cause of magnetospheric substorms, Ann. Geophysicae, 17, 604, 1999 Buchner, J., and L. M. Zelenyi, Deterministic chaos in the dynamics of charged particles near a magnetic field reversal, Phys. Lett. A, 118, 395, 1986. Buchner, J., and L.M.Zelenyi, Regular and chaotic charged particle motion in magnetotaillike field reversals: 1. basic theory of trapped motion, J. Geophys. Res., 94, 11821-11842, 1989. Coppi, B., G.Laval, and R. Pellat, Dynamics of the geomagnetic tail, Phys. Rev. Letters, v. 16, N 26, 1207-1210, 1966. Coroniti, F.V., On the tearing mode in quasi-neutral sheets, J.Geophys.Res., 85, N A12, 6719, 1980. Daughton, W., Kinetic theory of the drift kink instability in a current sheet, J.Geophys.Res., 103, 29429, 1998. Dunlop M.W., A. Balogh, P. Cargill, R.C. Elphic, K.-H. Fomacon, E. Georgesku, F. Sedgemore-Schulthess, and FGM team, Cluster observes the Earth's magnetopause: coordinated four-point magnetic field measurement, Annales Geophysicae, 19, 1449, 2001 Eastwood, J. W., Consistency of fields and particle motion in the 'Speiser' model of the current sheet, Planet. Space Sci., 20, 1555,1972. Galeev A. A., and Zelenyi L. M , Metastable states in the diffusion neutral layer and explosive phase of substorm, ZhETP Lett, 22, 170-172, 1975. Harris E. G., On a Plasma Sheath Separating Regions of Oppositely Directed Magnetic Fields, Nuovo Chimento, 23,115,1962. Hesse M., J. Birrn, and M. Kuznetsova, Collisionless magnetic reconnection: Electron processes and transport modeling, J. Geophys. Res., 106, 3721, 2001 Hoshino M., A. Nishida, T. Mukai, Y. Saito, and T. Yamamoto, Structure of plasma sheet in magnetotail: doublepeaked electric current sheet, J. Geophys. Res., 101, 24775, 1996.
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Kropotkin, A.P., and V.I. Domrin, Theory of a thin one-dimensional current sheet in collisionless space plasma, J.Geophys. Res., 101, 19893, 1996. Lembege B., and R. Pellat, Stability of a thick two-dimensional quasineutral sheet, Phys. Fluids, 25, 1995, 1982. Lui A. T. Y., R. E. Lopez, B. J. Anderson, K. Takahashi, L.Z. Zanetti, R.W. McEntire, T.A. Potemra, D. M. Klumpar, E.M. Greene, and R. Strangeway, Current disruptions in the near-Earth neutral sheet region, J. Geophys. Res., 97, 1461, 1992. Mitchell, D. G., G. J. Williams, C. Y. Huang, L. A. Frank, and C. T. Russell, Current carriers in the near-Earth cross-tail current sheet during substorm growth phase, Geophys. Res. Lett., 17, 583, 1990. Neishtadt A.I., About the change of adiabatic invariant during the separatrix crossing in two-dimensional freedom systems, Applied Mathematics and Mechanics, 51, 750, 1987. (In Russian) Petchek H.E., The concept of rapid magnetic field reconnection: a retrospective view, Geophys. Monogr. Ser., v.90, Physics of the magnetopause, ed.by P.Song, B.U.O. Sonnerup, M.F.Thomsen, pp.21-28,1995. Pulkkinen, T. I., D. N. Baker, C. J. Owen, J. T. Gosling, and N. Murthy, Thin current sheets in the Deep Geomagnetotail, Geophys. Res. Lett., 20, 2427 (1993). Pulkkinen T. I., D. N. Baker, D. G. Mitchell, R. L. McPherron, C. Y. Huang, and L. A. Frank, Thin current sheets in the magnetotail during substorms: CD AW 6 revisited, J. Geophys. Res., 99, 5793, 1994. Runov A., Nakamura R., Baumjohann W., Zhang T.I., Volverk M., Cluster observation of a bifurcated current sheet, Geophys.Res.Lett.,30, 1036,doi:10.1029/2002GL016136, 8-1 - 8-4, 2003. Sergeev V. A., D. G. Mitchell, C. T. Russell, and D. J. Williams, Structure of the tail plasma/current sheet at 11 Re and its changes in the course of a substorm, J. Geophys. Res., 98, 17345, 1993. Sergeev V. A., V. Angelopoulos, C. Carlson, and P. Sutcliffe, Current sheet measurements within a flapping plasma sheet, J. Geophys. Res., 103, 9177, 1998. Sergeev V.A., Runov A., Baumjohann W., Nakamura R., Zhang T.L., Volwerk, Balogh A., Reme H., Sauvaud J.A., Andre M., Klecker B., Current sheet flapping motion and structure observed by Cluster, Geophys.Res. Res., 30, 1327, doi:10.1029/2002GL016500, 60-1-60-4, 2003. Silin I., J. Buchner, and L.M. Zelenyi, Instabilities of collisionless current sheets: theory and simulations, Phys. Plasm, 2, 1104,2002. Sitnov M. I., Malova H. V., Sharma A. S., Role of temperature ratio in tearing stability of the quasi-neutral sheet tearing mode, Geophys.Res.Lett., 25, 269-272, 1998 Sitnov, M. I. , L.M. Zelenyi, H.V. Malova, and A.S. Sharma, Thin current sheet embedded within a thicker plasma sheet: Self-consistent kinetic theory, J. Geophys. Res., 105 , 13,029, 2000. Sitnov, M. I. , A.S. Sharma, P. N. Guzdar and P. H. Yoon, Reconnection onset in the tail of Earth's magnetosphere, J. Geophys. Res., 107 (A9), SMP20, 2002. Speiser T. W., Particle trajectories in model current sheets; 1. Analytical solutions, J. Geophys.Res., 70, 4219, 1965. Yoon, P. H., and A. T. Y. Lui, Nonlocal ion-Weibel instability in the geomagnetic tail, J. Geophys.Res., 101, 4899, 1996. Zeiler A., D. Biskamp, J.F. Drake, B.N. Rogers, M.A. Shay, and M. Scholer, Three-dimensional particle simulations of collisionless magnetic reconnection, J.Geophys.Res, 107 (A9), 1230, doi:l 029/2001JA000287, SMP 6-1 - SMP 6-9, 2002 Zelenyi L., A. Galeev, and C. F. Kennel, Ion precipitation from the inner plasma sheet due to stochastic diffusion, J. Geophys.Res., 95, 3871, 1990. Zelenyi L., M.I. Sitnov, H.V. Malova, and A.S. Sharma, Thin and superthin ion current sheets, quasiadiabatic and nonadiabatic models, Nonlinear processes in Geophysics, 7, 127, 2000. Zelenyi L.M., D.C. Delcourt, H.V. Malova, and A.S. Sharma, "Aging" of the magnetotail thin current sheets, Geophys. Res. Lett., 29, 10.1029/2001GL013789, pp. 49-1 - 4 9 - 4 , 2002a. Zelenyi L, M., H. V. Malova, V. Yu. Popov, D. C. Delcourt, and A. S. Sharma, Catastrophic-like evolution of thin current sheets due to non-adiabatic scattering processes, Proc.Int.Conf.Substorm-6, Univ. Washington, Seattle, 25-29 March, 2002, 245-252, 2002b.
E-mail address of L.M. Zelenyi lzelenyi(S)iki. rssi.ru
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STRATIFIED CURRENT SHEET DURING PLASMA SHEET THINNING M. Hoshino 1 1
University of Tokyo, 7S-1 Hongo, Bunkyo, Tokyo 113-0033 Japan
ABSTRACT The current sheet structure during the plasma sheet compression is studied for a one-dimensional Harris-type current sheet. We find that a stratified structure embedded in a thick plasma sheet is formed by taking into account the non-MHD effect. The global (smoothed) current density is enhanced during the plasma sheet compression, but the neutral sheet current does not necessarily increase. The current density can be reduced than the initial state before the plasma sheet compression. This kinetic current sheet structure may play an important role on the onset mechanism of substorms in magnetotail.
INTRODUCTION The onset of substorms in the earth's magnetotail is the most important unresolved issue in space plasma. It is believed that the plasma sheet is one of the key elements for understanding the dynamic phenomena in the earth's magnetosphere. The modern satellite explorations of last decade showed that the plasma sheet has a stratified structure in which a thin current sheet is embedded in a thick plasma sheet. The thickness of the thin current sheet may become as thin as the thermal ion gyro-radius during the active magnetosphere [e.g. Fairfield, 1984^ McComas et al., 1986; Mitcheel et al., 1990]. Mukai et al. [1998] found a super thin current sheet during the substorm growth phase, which thickness is estimated to be less than the ion gyroradius. It is now recognized that the formation of the thin current sheet (TCS) is occurring in a wide region of magnetotail [e.g., Sergeev et la., 1993; Sanny et al., 1994; Pulkkinen et al., 1994; Asano et al. 2002]. In such a thin current sheet plasmas begin to behave in a nonadiabatic motion resulting in non-MHD effects, and both ion and electron dynamics become important. Since the configuration of TCS has a large amount of the free energy, it is favorable for various kinetic plasma instabilities. The onset of magnetic energy release in magnetotail is believed to be strongly coupled with the formation of TCS. It is expected that TCS seems to be formed when the solar wind energy is stored into the magnetosphere and the lobe magnetic field is increasing. In order to study the relationship between the lobe magnetic field and the central current sheet density, it is necessary to investigate the plasma diamagnetic current self-consistently coupled to the -108-
magnetic field profile. Furthermore, since the thickness of TCS becomes of the order of the ion gyro-radius, the non-MHD effects such as the inertia effect, a finite gyro-radius motion etc. should be taken into account. We study the nonlinear time evolution of the plasma sheet by using a particle-in-cell simulation code. EVOLUTION OF CURRENT SHEET: SIMULATION STUDY We consider the general problem of the plasma sheet evolution by compressing the plasma sheet. For simplicity, a one-dimensional, Harris current sheet configuration is assumed as an initial state, and the lobe magnetic field is allowed to increase at the outer boundary of the system to be time dependent. The magnetic field profiled is B x =B 0 tanh(ZA) at t=0. The induced dawn-dusk electric field from the outer boundary starts to penetrate into the plasma sheet with the magnetosonic speed, which results Fig 1. Evolution of the magnetic field in the formation of a thin current sheet. and electric current in a plasma sheet. We study the time evolution of the Note that the electric current increases plasma sheet by using a full with increasing the lobe magnetic field at particle-in-cell simulation, and we the outer boundary. discuss the plasma dynamics on the formation of TCS. The simulation box size is 1200 girds corresponding to 10A., and 106 particles for both ions and electron are used. In the present case, the mass ratio of ion to electron is assumed to be M/m=100, and the ion to electron temperature ratio is taken to be TION/TELE=4.2 in the initial state. The ion diamagnetic drift velocity is set to be 0.2 of the ion thermal velocity. Figure 1 shows the profiles of the electric current Jz and the magnetic field Bx at t=0 and 45, where the time is normalized by the Alfven transit time, TA. One can find that the "global" electric current is intensified with increasing the lobe magnetic field, but the current sheet shows a specific stratified structure in the vicinity of the neutral sheet at Z=0. Let us first discuss the relationship between the "global" (smoothed) electric current and the lobe magnetic field. In an MHD framework where the plasma is frozen into the magnetic field, the magnitude of Jz at the neutral sheet can be expressed as a function of the lobe magnetic field Bi at the outer boundary (Drake et al., 1981), Jz(t)cx(Bi(t)/Bo)4/Y«:(Bi(t)/Bo)2, -109-
where we assumed the equation of state T cc N1 ' with y-2 (i.e., two -dimensional heating perpendicular to the magnetic field). In order to compare the simulation result with the above MHD theory, we first estimate the current intensity in our simulation result by fitting the curve of sech2(Z/z*), which is depicted in the dashed lines in Figure 1. The fitting curve differs near the neutral sheet, but the curve shows the "global" current sheet profile quite well. In Figure 2, the open circles are the Jz(z=0) estimated from the simulation result as a function of the lobe magnetic field Bi, while the dashed curve is the least-square fit for a power-law function of Jz(t) oc (Bi(t)/Bo)p. We obtained p=1.85, which is close to the adiabatic MHD theory with p=2. This is suggestive of the quasi-adiabatic compression for the global current sheet structure. To further confirm the quasi-adiabatic Fig 2. Relationship between the neutral compression, we examine the density sheet current JNS(0 and the lobe magnetic and temperature relation at the field at the outer boundary Bi(t). neutral sheet. Figure 3 shows ion and electron temperatures as a function of the plasma density. Both ion and electron temperatures are normalized by each initial temperatures. The open circles and open squares are the simulation data of ion and electron, respectively, while the solid curves show a power-law function fitting of T ccNY J . We obtained y=2.29 for ion and Y=3.66 for electron. These results suggest that the ion heating is almost two-dimensional compression, while the non-adiabatic electron heating is occurring in the current sheet. Since the initial electron temperature is assumed to be 4.2 times smaller than the ion one, the electron contribution to the total temperature is still Fig 3. Relationship between the plasma negligible. Therefore, the temperature TNS(0 and the plasma temperature and density relation is density N(t) at the neutral sheet. Note almost consistent with the Jz^Bi that the initial ratio of the ion to electron relation in Figure 2. temperature TION/TELE=4.2. -110-
Our essential aim it to understand the current sheet structure embedded in the thick plasma sheet. Figure 4 shows the ion and electron current densities at t=0 (blue) and t=45 (red). The ion and electron currents are shown in positive and negative velocities, respectively. At t=45, we can clearly find that both ion and electron current profiles near the neutral sheet depart from the global current profile described by the sech2(Z/z*) curve fit. The absolute value of the electron current increases with approaching toward the neutral sheet, and its peak current becomes 10 times larger than the initial electron current. However, the current quickly decreases inside TJX < 0.2, and the current shows a positive value. For the ion current, the spatial profile is not so different from the initial current, but one can find the reduction of the current density near the neutral sheet as well. The above fine structure of the electric current may be understood by virtue of the polarization electric field produced by the inertia effect of electrons and ions. The dawn-dusk electric field induced from the outer boundary transports both ions and electrons toward the neutral sheet, but the ions start to be unmagnetized inside the ion inertia layer near the neutral sheet, while the magnetized electrons can be still effectively transported Fig 4. Ion and electron current profile in toward the neutral sheet. Then the the nonlinear evolution of the plasma Hall electric field Ey directed toward sheet. Double current sheet can be found the neutral sheet is induced in the ion for both ions and electrons. inertia layer. Therefore, the Ey x Bx drift motion for electrons can enhance the Jy current, while the effect of the Ey x Bx motion for ions makes the ion Jy current weaken in the ion meandering width. The ions, however, are almost unmagnetized in the layer and this effect is smaller than that for electrons. Inside the electron inertia length, the electric field Ey induced by the electron pressure seems to play an important role. Note that we discussed that a strong electron heating near the plasma sheet in Figure 2, and we found that various plasma waves are excited in the plasma sheet driven by the plasma compression (not shown here). The driven force from the outer boundary seems to be the source of the plasma waves that contribute to the electron heating. By assuming the Boltman-Maxwell state in the thin layer of the electron inertia scale, we can expect the outward Ey from the neutral sheet in a quasi-steady state, which in turn decelerates the meandering electron in the vicinity of the neutral sheet. The thickness of the electron meandering width can be estimated (mTELE/MTiON)1/4(^i X)112 ~ 0.1A. in our simulation parameter. We can find that the thickness of the electron thin current layer in Figure 4 is of the order of 0. IX. -111-
CONCLUSIONS We discussed how a thin current sheet can be formed during the plasma sheet compression, i.e., the energy loading stage in magnetotail. We showed that the kinetic plasma effects play a significant role on the thin current sheet formation, and the double current sheet can be form in consequence of the plasma sheet thinning. The electric current at the neutral sheet does not necessarily increase during the plasma sheet thinning. This study assumed one-dimensional, slab geometry of the plasma sheet, and many important plasma modes such as lower-hybrid-drift instability and tearing instability are not included. Our objective is to investigate the current sheet profile just before the onset of the magnetic reconnection or the current disruption etc., where the effects of such plasma instabilities do not emerge. However, we believe that this paper has important implications for understanding the dynamics on the onset of substorms in the earth's magnetotail.
REFERENCES Asano, Y., T. Mukai, M. Hoshino, Y. Saito, H. Hayakawa, and T. Nagai, Evolution of the thin current sheet in a substorm observed by Geotail, submitted to J. Geophys. Res. (2002) Drake, J. E, N. T. Gladd, and J. D. Huba, Magnetic field diffusion and dissipation in reversed-field plasmas, Phys. Fluids 24, 78 (1981) Fairfield, D. H., Magneotail energy storage and the variability of the magnetotail current sheet, Magnetic Reconnection in Space and Laboratory Plasmas, AGU/Geophys. Mono., W. Hones, Jr. Ed., 30, 168 (1984) McComas, D. J., C. T. Russel, R. C. Elphic, and S. J. Bame, The near-earth cross-tail current sheet: Detailed ISEE 1 and 2 case studies, J. Geophys. Res., 91, 4287 (1986) Mitchell, D. G., D. J. Williams, C. Y. Huang, L. A. Frank, and C. T. Russell, Current carriers in the near-earth cross-tail sheet during substorm growth phase, Geophys. Res. Lett., 17, 583 (1990) Mukai, T, M. Hoshino, Y. Saito, I. Shinohara, T. Yamamoto, T. Nagai, and S. Kokubun, Pre-Onset and Onset Signatures for Substroms in the Near-Tail Plasma Sheet: GEOTAIL Observations, Substorms-4, S. Kokubun and Y. Kamide, Eds., Terra Sci. Publ., P p.l31-136(l998). Pulkkinen, T. I., D. N. Baker, D. G. Mitchell, R. L. McPherron, C. Y. Huang, and L. A. Frank, Thin current sheets in the magnetotail during substorms: CDAW 6 revisited, J. Geophys. Res., 99, 5793 (1994) Sergeev, V. A., D. G. Mitchell, C. T. Russell, and D. J. Williams, Structure of the tail plasma/current sheet at ~HRe and its changes in the course of a substorm, J. Geophys. Res., 98, 17345 (1993) Sanny, J., R. L. McPherron, C. T. Russell, D. N. Baker, T. I Pulkkinen, and A. Nishida, Growth-phase thinning of the near-Earth current sheet during the CDAW 6 substorm, J. Geophys. Res., 99, 5805 (1994) -112-
KINETIC INSTABILITIES IN A THIN CURRENT SHEET A. T. Y. Lui Johns Hopkins University/Applied Physics Laboratory, Laurel, MD 20723, USA
ABSTRACT A thin current sheet with thickness comparable to the ion inertial length gives rise to decoupling of ions from electrons. This situation is favorable for a number of dynamic processes leading to particle acceleration and plasma turbulence. In this article, we investigate the cross-field current instability that can be excited in a thin current sheet with a non-zero magnetic field component normal to the current sheet surface. There is strong observational evidence for this instability and substantial theoretical work on it. We first discuss briefly the observational basis, followed by a brief overview of theoretical predictions. We then present particle simulation of this instability for a thin current sheet using a two-dimensional fully electromagnetic particle-incell code. The instability is found to develop rapidly, resulting in particle acceleration, current filamentation, current density reduction at the sheet center, current sheet broadening, and undulations of the current sheet profile. These features revealed from simulation are quite similar to those observed for current disruption. INTRODUCTION There has been considerable work on thin current sheet dynamics. This trend is motivated by observations that very thin current sheets with thickness comparable to the ion inertial length tend to exhibit very dynamic behavior, e.g., particle acceleration and large magnetic and electric fluctuations [Lui et ah, 1988, 1990; Mitchell et al, 1990; Sergeev et al, 1993; Ohtani et al, 1995; Mukai et al, 2000]. A thin current sheet leads to decoupling of ions and electrons, thus departing significantly from the magnetohydrodynamic (MHD) representation of magnetospheric plasma as a single fluid. Consequently, the kinetic approach, which is far more elaborate than the simple fluid representation in the MHD approach, is required to ascertain the physical processes involved in these dynamic phenomena. Most of the work on thin current sheets is aimed on the possibility of associating energetic phenomena with magnetic reconnection. This underlying expectation led to work focusing on thin current sheets with a zero magnetic field component normal to the current sheet so that reconnection of opposing magnetic field lines on the two sides of a current sheet may occur. However, it is clearly demonstrated from observations in the near-Earth magnetotail that significant magnetic field reconfiguration (dipolarization) occurs in a thin current sheet with a significant normal magnetic field component, strong enough to eliminate tearing instability or magnetic reconnection as the cause [Takahashi et al, 1987; Lui et al., 1988; Ohtani et al, 1992]. It is therefore inconceivable to associate current disruption phenomenon directly with magnetic reconnection or tearing instability in the near-Earth magnetotail. Thus, work investigating thin current sheet dynamics with zero normal magnetic field may be appropriate to the mid-tail environment with a small normal magnetic field component but may be irrelevant to account for current disruption in the near-Earth magnetotail where the normal magnetic field is strong [Lui et al., 1991; Zeienyi et al., 1998]. In this paper, we present observations that have provided the foundation for the development of the cross-field current instability as a potential mechanism for current disruption and substorm expansion onset. Theoretical work along this line is briefly reviewed to provide the theoretical expectations. The nonlinear evolution of the instability is then examined with a two-dimensional particle simulation run for a thin current sheet configuration with a particle-in-cell (PIC) code. The results show very encouraging agreements with observations on the current disruption phenomenon in the near-Earth magnetotail and theoretical work relating to the cross-field current instability for current disruption and substorm onset/intensification. -113-
Fig. 1. An example of magnetic fluctuations in current disruption in the field transition region from dipole to tail-like. CURRENT DISRUPTION OBSERVATIONS Current disruption responsible for substorm expansion onset occurs in the transition region between the dipolar and the tail-like magnetic field configurations. Figure 1 illustrates this point. On the left is a diagram of the magnetosphere. The magnetic field transition region lies near the inner edge of the plasma sheet and the nightside cross-tail current system. On the right is an example of current disruption. Before the activity onset, the magnetic field is rather steady. The horizontal magnetic field component B\\ in the VDH coordinate system is depressed substantially below the dipole value, indicating the existence of a strong current density in the vicinity. This development is consistent with the expectation of the substorm growth phase in which the cross-tail current sheet moves inward at this time. At about 1152 UT, all field components show high levels of fluctuations. This turbulent state, referred to as current disruption, lasts for more than 3 min before the field configuration settles down. After the activity, the B]\ component becomes much stronger, reflecting the reduction of the local current density and approaching to a more dipolar field value. This field reconfiguration, called dipolarization, can be viewed as a relaxation of the stretched magnetic field condition before onset. Let us first examine the bulk plasma behavior around the current disruption interval when the satellite was in the neutral sheet throughout a current disruption event. Figure 2 shows the evolution of plasma pressure and plasma beta. Figure 2a and 2b show the temporal developments of plasma pressure components for three particle species, namely, protons, oxygen ions, and electrons. Proton pressure is the dominant one (for both components) and shows the largest increase prior to current disruption onset at -2314 UT on June 1, 1985. The combined total pressure components are given in Figure 2c. Figure 2d shows plasma beta exceeding 20 prior to onset, reaching to about 70 just before onset, and finally decreasing to less than 5 afterwards. Let us next examine the kinetic properties of the particles. Figure 3 shows the ion behavior. The panels in the figure give the magnitude and latitude angle of the magnetic field accompanied by the ion anisotropy of 31 -43 keV energetic ions. The bottom row of panels shows the ion distribution at three time snapshots. Well before the current disruption onset, the ion distribution appears to be isotropic. Just before onset, the ions exhibit a strong duskward drift, consistent with the expectation of ions undergoing Speiser orbits in the thin current sheet just before onset. Afterwards, the ion distribution is anisotropic with the perpendicular temperature higher than the parallel.
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Time Relative to Current Disruption Onset
Fig. 2. Temporal development of (a) perpendicular plasma pressure, (b) parallel plasma pressure, (c) total plasma pressure, and (d) plasma beta around the current disruption onset at -2314 UT on June 1, 1985.
Ion Velocity Distribution
Fig. 3. Ion velocity distribution at the current sheet around the current disruption interval.
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Fig. 4. Electron behavior and their pitch angle distribution around the current disruption interval.
Fig. 5. Wavelet analysis of current disruption reveals low- and high-frequency components.
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Figure 4 shows the accompanying electron behavior. The top panels show the electron pressure, measured pitch angle, and the magnetic field magnitude. The bottom rows are three time snapshots of the electron pitch angle distribution. Before onset, the electrons have a pancake distribution, peaking at 90° pitch angle. During current disruption, the electrons exhibit a rather isotropic distribution. At the end, the electron pitch angle is cigar-shaped with fluxes peaked along the magnetic field. The wave activity associated with the disturbance is reviewed in Figure 5 with wavelet analysis. The magnetic field component B v is shown at the top and the wavelet decomposition of the signal is shown at the bottom. Superposed at the bottom panel is the trace of local ion cyclotron frequency. About 1 min before onset, a low frequency wave is seen. At onset, a band of wave activity is seen, with frequency reaching well above the ion cyclotron frequency. Later, waves at intermediate frequencies are seen. Intermittent bursts of wave activity after onset are seen at later times as well. This wavelet analysis indicates multiple frequency components in the observed magnetic fluctuations, suggesting the onset of multiple plasma processes. A POSSIBLE ONSET SCENARIO AND THEORETICAL ANALYSIS OF INSTABILITIES The above set of observations suggests a possible scenario for the onset of current disruption and substorm expansion phase. There is probably more than one instability involved. One potential process is the ballooning/interchange instability. This instability has been suggested as a potential substorm onset mechanism by a number of researchers [Liu, 1970; Roux et al., 1991; Samson et al., 1992; Pu et al., 1992; Erickson, 1995; Bhattacharjee et al., 1998; Cheng and Lui, 1998; Horton et al., 1999; Lee, 1999; Wong et al., 2001]. Evidence for its role comes from the buildup of plasma pressure during the substorm growth phase and the development of low frequency magnetic perturbations just prior to current disruption onset. As a result of the nonlinear evolution of the ballooning instability, magnetic field lines stretch tailward with accompanying strong westward ion drift and particle acceleration. A thin current sheet with a strong current density is thus formed. The sudden development of strong current density excites the cross-field current instability (CCI), generating high frequency waves and causing current disruption. Nonlinear development of CCI gives rise to ion heating perpendicular to the magnetic field and electron heating in the parallel direction. Let us first discuss the ballooning instability. Recently, a kinetic form of this instability is examined [Cheng and Lui, 1998] by including finite ion Larmor radius effect to allow for the presence of non-zero parallel electric field and trapped electron dynamics. These kinetic effects are important since they lead to a higher onset threshold than that predicted by the MHD model. The MHD formulation of ballooning instability predicts onset threshold at plasma beta ft ~1 while the kinetic ballooning instability (KBI) predicts it to be at /3-100. The higher critical fi for KBI comes from stabilization of field configuration by field-aligned currents set up by the non-zero parallel electric field. The MHD prediction is inconsistent with observation that shows current sheet being stable even at / 3 » 1 (see Figure 2d). The observed value of the plasma beta at which the current sheet remains stable can be used to test any process proposed for substorm onset. In particular, it can be used to judge the validity of any modification of the MHD ballooning instability as a viable mechanism.
Fig. 6. The quasi-linear calculations show current reduction, acceleration of electrons and ions, and wave growth.
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KBI at its nonlinear stage could lead to a rapid development of strong current densities in the near-Earth current sheet as noted earlier. The strong current density in a thin current sheet created by KBI can then instigate the onset of CCI. Linear and quasi-linear analyses of CCI have been performed based on this scenario [Lui et ah, 1991, 1993; Yoon and Lui, 1993]. From these theoretical analyses, it is shown that CCI could be excited when the relative drift between ions and electrons (which constitutes the current density) approaches a large fraction of the ion thermal speed. The excited waves are oblique whistlers in a broad range of frequency from below the ion cyclotron frequency to well above it. Quasi-linear calculations, e.g. one given in Figure 5, indicate that the current could be reduced by -15-28 %, accompanied by ion heating perpendicular to the magnetic field and electron acceleration along the magnetic field as observed in current disruption phenomenon. The combination of KBI and CCI can account for multiple waves with low and high frequencies being excited during current disruption, enhanced current density just prior to substorm onset [Lui et ah, 1992; Ohtani et ah, 1992], ion heating perpendicular to the magnetic field, and magnetic-field-aligned electron acceleration during current disruption [Lui et ah, 1992]. In addition, it can explain naturally several wellknown observed substorm features [Lui et ah, 1991], such as the skew of substorm onset location toward the evening sector, three different solar wind conditions (including northward turning of the interplanetary magnetic field) for substorm onset, spatial localization of onset regions, and substorm onset location tied to a pre-existing auroral arc [Lui and Murphree, 1998]. Furthermore, since a strong magnetic field normal to the current sheet is favorable for excitation of CCI, the instability therefore can account for substorm onset location residing deep inside the magnetosphere. This is particularly important because substorm onsets during magnetic storm periods occur at very low magnetic latitudes and are on magnetic field lines with equatorial crossing deep inside the inner magnetosphere. This set of favorable evidence for this instability seems to be unmatched by other processes proposed for substorm onsets. TWO DIMENSIONAL PARTICLE SIMULATION In order to see the nonlinear development of the CCI, we have performed a two-dimensional simulation of a thin current sheet using a fully electromagnetic PIC code. In this code, the trajectory of each particle is followed by solving the equation of motion in the presence of the self-consistent electric and magnetic fields, which are determined by solving Maxwell's equations with charge and current densities due to particles collected on a grid. Each simulated particle has a finite-size specified by a shaping function. The force exerting on each particle is calculated by interpolating the fields at its nearby grid points. The contribution of each particle to the charge and current densities is collected at the nearby grid points through interpolation also. Relativistic equation of motion is used in the code. The Maxwell equations are solved through Fast Fourier Transform (FFT) by decomposing the electric field and current into longitudinal and traverse components.
Fig. 7. A schematic diagram to show the setup of the simulation.
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Figure 7 shows the setup for a 2D thin current sheet. The current sheet documented in Lui et al [1995] and Yoon and Lui [1996] is adopted here. This is a modified form of the Harris current sheet and has an exact force balance in the z-direction, i.e., the pressure gradient force balances the Lorentz force. Distinct from the Harris current sheet which has a constant drift velocity all across the sheet, this current sheet has a Lorentz profile of velocity peaking at the sheet center. Velocity shear is found to be essential based on nonlocal instability analysis [Lui et al, 1995]. It is assumed here that such a thin current sheet is formed by the nonlinear development of KBI and does not necessary have a force balance in the x-direction. Furthermore, the time scale for the evolution of the kinetic instability is much shorter than the time scale for equilibrium evolution. Therefore, this force imbalance is not expected to affect significantly the simulation results. The parameters for the simulation run are T[/Te = 8, m[/me = 100, Bn/Bo = 0.2, v e /c = 0.05, and a)p e /Q ce = 6.4. The sheet half-thickness Lz is taken to be 1.3 times the thermal ion gyroradius. A total of -2.5 x 10 6 particles were used over a grid of 256 (z-direction) x 32 (y-direction). The design of this numerical simulation aims to examine excited waves in the wavelength range of 5 < kyLz < 40 expected for the CCI and DKSI [Yoon et al., 2002]. Since there are many previous particle simulations of thin current sheets, it is important to point out that this simulation differs from all previous particle simulations by investigating a current sheet with a strong magnetic field component normal to the current sheet and a Lorentz velocity profile with a peak at the center of the current sheet. These features could be responsible for the difference between previous particle simulation results and the present one. We consider the designed current sheet here to resemble closer to the near-Earth current sheet just prior to onsets of current disruption than a Harris current sheet. The simulation results are given in Figure 8. The initial ion velocity distribution near the center of the current sheet is given in Figure 8a. A definite ion drift in the ^-direction is evident, which contributes to the current density of the current sheet since electrons are initially set to have no net drift. Figure 8b shows the ion velocity distribution in the same region in the simulation box at a later time («Jpe T = 2000). There is clear indication of ion heating from the excitation of the instability. Similar heating of electrons is found. Current densities of the simulated current sheet near its center are shown in Figure 8c for the same time as in Figure 8b. Currents within the sheet at this time are highly inhomogenous, with regions of intense current densities bordered by regions of much reduced current densities. These structures represent the filamentation of the smooth current density profile set up for this simulation initially. A close examination of the current sheet profile indicates that at some j-locations, the current sheet profile develops an even (sausage) undulation of the current sheet boundary while at other ^-locations it has an odd (kink) undulation. The average amount of current density reduction is shown in Figure 8d in which the initial current density profile (given by the dotted curve) is compared with the current density profile seen at this later time (given by the solid curve). These two profiles are current densities averaged over the y-dimension (representing the dawn-dusk dimension). This panel shows clearly that the average current density at the current sheet center is reduced significantly by —40% with an overall broadening of the current density profile. Since there were debates in the literature about the relative importance of sausage (even) and kink (odd) modes in the evolution of a thin current sheet [Zhu and Winglee, 1996; Pritchett et al., 1996; Lapenta and Brackbill, 1997; Yoon et al, 1998; Daughton, 1999; Biichner and Kuska, 1999], we have also decomposed the electric perturbations into even and odd modes. The results are presented in Figure 8e and Figure 8f. These plots indicate that both modes have comparable amplitudes and growth rates, with the odd mode having slightly higher values. This is consistent with the even and odd undulations of the current sheet profile revealed in Figure 8c. Furthermore, there is an indication of nonlinear coupling between these two modes. For example, the reduced current density around the region defined by y = 5-10 and z = 130-140 in Figure 8c represents the even mode perturbation (Figure 8e) reinforcing the odd mode perturbation (Figure 8f). The frequency and wavelength of the instability responsible for the evolution of the thin current sheet have been examined. The excited waves are found to have a relatively broad frequency between the ion gyrofrequency and the lower hybrid frequency. The wavelengths for these modes lie in the range of 14 < kyLz < 23, coinciding with the theoretical expectation as discussed in Yoon et al. [2002]. These excited modes are further verified by examining the time evolution on the amplitudes of electric perturbations, which indicates the excited waves to have a broad frequency between the ion gyrofrequency and the lower hybrid frequency. We have also examined plasma parameters at two different locations in the simulation box, one located near the center of the current sheet and the other located near the boundary of the current sheet. This exercise is to find out the temporal development of plasma parameters at the center and at the edge of the current sheet. We have found that the magnetic field components, number and current densities all show large variations during the simulation run. The number density decreases in the current sheet center while it increases near the current sheet boundary. These changes are consistent with a decrease in the peak current density and current sheet broadening noted in Figure 8d.
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Ion Velocity Distribution Near Sheet Center
Fig. 8. Simulation results: (a) The ion velocity distribution f,(v) of the current sheet at initial setup, (b) f|(v) at a later time, (c) current density near current sheet center, (d) current density profile averaged over the ydimension, (e) even mode electric field, (f) odd mode electric field.
DISCUSSION AND CONCLUSIONS We have reviewed observations, theoretical calculations, as well as particle simulation pertinent to the cross-field current instability for a thin current sheet. It is pointed out that there is strong observational basis for considering this instability as the physical process responsible for current disruption and substorm -120-
expansion onset. This instability can account for a large number of observational features associated with substorm onset. An important aspect of this instability is that it can be excited in a high current density region with a strong magnetic field normal to the current sheet. This feature allows it to account for substorms which occur deep within the inner magnetosphere during magnetic storm intervals. This instability is thus a very promising potential physical process for substorm onset. The identification of this instability in the numerical simulation is based on the consideration of the characteristics of the excited mode. Previous theoretical (local and nonlocal) analyses of CCI indicate the excited mode to have a broad frequency band between the ion gyrofrequency and the lower hybrid frequency and that the wavelength to be short compared with those associated with drift-kink/sausage modes. Furthermore, the excited waves are highly electromagnetic and concentrate near the current sheet center, distinct from the lower hybrid drift instability that excites mainly electrostatic waves located mostly at the edge of the current sheet. All these characteristics expected from CCI are found in the numerical simulation. For these reasons, we tentatively attribute the instability responsible for current filament in the numerical simulation to be the CCI. However, there are a number of issues requiring further work. For example, kinetic nonlocal instability theory for a thin current sheet with a strong magnetic field normal to the current sheet surface is still lacking. The work of Yoon et al. [2001] examines a current sheet without magnetic field normal to the current sheet. Therefore, detailed comparison of simulation results with theoretical predictions has yet to be made. Future studies of this potential scenario for current disruption and substorm onset/intensification would be to evaluate nonlinear evolution of the kinetic ballooning instability for thin current sheet generation, to develop non-local kinetic theory of the cross-field current instability, and to perform three-dimensional numerical simulations for the dynamic evolution of a thin current sheet for particle acceleration in current disruption and for the generation of the substorm current wedge. ACKNOWLEDGMENTS This work was supported by NASA Grant NAG5-10475 and the Atmospheric Sciences Division of NSF Grant ATM-0135667 to the Johns Hopkins University Applied Physics Laboratory. The author is grateful to Drs. Y. Kamide and T. Ogino for collaboration and access to the Nagoya Supercomputer facility, to Dr. P. H. Yoon for useful discussions and collaboration on current sheet instabilities, and to Dr. Sifeng Ma for providing an early version of the particle code. REFERENCES Bhattacharjee, A., Z. W. Ma, and X. Wang, Ballooning instability of a thin current sheet in the high-Lundquist number magnetotail, Geophys. Res. Lett., 25, 861, 1998. Biichner, J., and J.-P. Kuska, Sausage mode instability of thin current sheets as a cause of magnetospheric substorms, Ann. Geophys., 17, 604, 1999. Cheng, C. Z. and A. T. Y. Lui, Kinetic ballooning instability for substorm onset and current disruption observed by AMPTE/CCE, Geophys. Res. Lett., 25, 4091, 1998. Daughton, W., The unstable eigenmodes of a neutral sheet, Phys. Plasmas, 6, 1329, 1999. Erickson, G. M., Substorm theories: United they stand, divided they fall, Rev. Geophys., 33, 685, 1995. Horton, W., H. V. Wong, and J. W. VanDam, Substorm trigger conditions, J. Geophys. Res., 104, 22745, 1999. Lapenta, G. and J. U. Brackbill, A kinetic theory for the drift-kink instability, J. Geophys. Res., 102, 27099, 1997. Lee, D.-Y., Effect of plasma compression on plasma sheet stability, Geophys. Res. Lett., 26, 2705, 1999. Liu, C. S., Low-frequency drift instabilities of the ring current belt, J. Geophys. Res., 75, 3789, 1970. Lui, A. T. Y. and J. S. Murphree, A substorm model with onset location tied to an auroral arc, Geophys. Res. Lett.,IS, 1269-1272, 1998. Lui, A. T. Y. and A.-H. Najmi, Time-frequency decomposition of signals in a current disruption event, Geophys. Res. Lett, 24, 3157, 1997. Lui, A. T. Y., A. Mankofsky, C. L. Chang, K. Papadopoulos, C. S. Wu, A current disruption mechanism in the neutral sheet: a possible trigger for substorm expansions, Geophys. Res. Lett., 17, 745, 1990. Lui, A. T. Y., C.-L. Chang, A. Mankofsky, H.-K. Wong, and D. Winske, A cross-field current instability for substorm expansions, J. Geophys. Res., 96, 11389, 1991.
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Lui, A. T. Y., R. E. Lopez, B. J. Anderson, K. Takahashi, L. J. Zanetti, R. W. McEntire, T. A. Potemra, D. M. Klumpar, E. M. Greene, and R. Strangeway, Current disruptions in the near-Earth neutral sheet region, J. Geophys. Res., 97, 1461, 1992. Lui, A. T. Y., P. H. Yoon, and C.-L. Chang, Quasi-linear analysis of ion Weibel instability in the Earth's neutral sheet, J. Geophys. Res., 98, 153, 1993. Lui, A. T. Y., C.-L. Chang, and P. H. Yoon, Preliminary nonlocal analysis of cross-field current instability for substorm expansion onset, J. Geophys. Res., 100, 19147, 1995. Mitchell, D. G., D. J. Williams, C. Y. Huang, L. A. Frank, C. T. Russell, Current carriers in the near-Earth cross-tail current sheet during substorm growth phase, Geophys. Res. Lett., 17, 583, 1990. Mukai, T., T. Nagai, M. Hoshino, Y. Saito, I. Shinohara, T. Yamamoto, S. Kokubun, Geotail Observations of magnetic reconnection in the near-Earth magnetotail, Adv. Space Res., 25, 1679-1683, 2000. Pu, Z. Y., A. Korth, and G. Kremser, Plasma and magnetic field parameters at substorm onsets derived from GEOS 2 observations, J. Geophys. Res., 97, 19341, 1992. Sergeev, V. A., D. G. Mitchell, C. T. Russell, and D. J. Williams, Structure of the tail plasma/current sheet at ~11 R E and its changes in the course of a substorm, J. Geophys. Res., 98, 17345, 1993. Ohtani, S., S. K. Takahashi, L. J. Zanetti, T. A. Potemra, R. W. McEntire, and T. Iijima, Initial signatures of magnetic field and energetic particle fluxes at tail reconfiguration: Explosive growth phase, J. Geophys. Res., 97, 19311, 1992. Ohtani, S., T. Higuchi, A. T. Y. Lui, and K. Takahashi, Magnetic fluctuations associated with tail current disruption: fractal analysis, / . Geophys. Res., 100, 19135, 1995. Pritchett, P. L., F. V. Coroniti, and V. K. Decyk, Three-dimensional stability of thin quasi-neutral current sheets, J. Geophys. Res., 101, 27413, 1996. Roux, A., S. Perraut, P. Robert, A. Morane, A. Pedersen, A. Korth, G. Kremser, B. Aparicio, D. Rodgers, R. Pellinen, Plasma sheet instability related to the westward traveling surge, J. Geophys. Res., 96, 17697, 1991. Samson, J. C , D. D. Wallis, T. J. Hughes, F. Creutzberg, J. M. Ruohoniemi, and R. A. Greenwald, Substorm intensifications and field line resonances in the nightside magnetosphere, J. Geophys. Res., 97, 8495, 1992. Takahashi, K., L. J. Zanetti, R. E. Lopez, R. W. McEntire, T. A. Potemra, and K. Yumoto, Disruption of the magnetotail current sheet observed by AMPTE/CCE, Geophys. Res. Lett, 14, 1019, 1987. Wong, H. V., W. Horton, J. W. VanDam, and C. Crabtree, Low frequency stability of geotail plasma, Phys. Plasmas, 8, 2415-2424, 2001. Yoon, P. H. and A. T. Y. Lui, Nonlinear analysis of generalized cross-field current instability, Phys. Fluids, B5 (3), 836, 1993. Yoon, P. H. and A. T. Y. Lui, Nonlocal ion-Weibel instability in the geomagnetic tail, J. Geophys. Res., 101, 4899, 1996. Yoon, P.-H., A. T. Y. Lui, and H. K. Wong, Two-fluid theory of drift-kink instability in one-dimensional neutral sheet, J. Geophys. Res., 103, 11875, 1998. Yoon, P. H., A. T. Y. Lui, and M. I. Sitnov, Generalized lower-hybrid drift instabilities in current-sheet equilibrium, Phys. Plasmas, 9, 1526, 2002. Zelenyi, L. M., A. A. Petrukovich, E. Y. Budnick, S. A. Romanov, V. A. Sergeev, T. Mukai, T. Yamamoto, S. Kokubun, K. Shiokawa, C. S. Deehr, J. Buchner, and I. Sandahl, Substorm onset models and observations, Substorms-4, ed. by S. Kokubun and Y. Kamide, Terra Scientific Publishing Co., Tokyo, Japan, 327, 1998. Zhu, Z., and R. M. Winglee, Tearing instability, flux rope, and the kinetic current sheet kink instability in the Earth's magnetotail: A three-dimensional perspective from particle simulations, J. Geophys. Res., 101, 4885, 1996. Email address of A.T.Y. Lui [email protected]
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FORMATION OF THIN ELECTRON CURRENT LAYER ASSOCIATED WITH LOWER HYBRID DRIFT INSTABILITY AND ITS RELATION TO QUICK RECONNECTION TRIGGERING I. Shinohara1, and M. Fujimoto2 'Japan Aerospace Exploration Agency, Institute of Space and Astronautical Science, 3-1-1 Yoshinodai, Sagamihara, Kanagawa 229-8510, Japan 2 Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8551, Japan
ABSTRACT We have recently found that a quick triggering of magnetic reconnection in an ion-scale current sheet is possible. For the quick triggering of magnetic reconnection, the lower hybrid drift waves excited at the edges of the current sheet is indispensable. This wave excitation brings about formation of a thin magnetic neutral layer sustained by accelerated electrons, and this thin layer is subject to the quick reconnection. We found that the electron acceleration process is strongly coupled with the non-linear evolution of the lower hybrid drift instability. The inductive electric field, which is generated through the change of the current profile, can efficiently accelerate meandering electrons around the magnetic neutral layer. As a result, electric current in the thin layer is mostly carried by non-adiabatic electrons. The production of non-adiabatic electrons is playing a crucial role in making the quick triggering available.
INTRODUCTION Investigations of the Earth's magnetotail by Geotail and the preceding satellites reveal that magnetic reconnection in the magnetotail really plays an essential role to the magnetospheric energy release phenomena. In collisionless plasma, such as the magnetospheric plasma, an anomalous dissipation process is necessary for driving magnetic reconnection. So far, a number of theories on the anomalous dissipation process have been proposed, and among these theories, the anomalous resistivity generated by the lower hybrid drift (LHD) instability has been thought as a plausible candidate. However, using the Geotail wave data, Shinohara et al. (1998) have shown that the observed LHD waves cannot provide sufficient resistivity for quick triggering of magnetic reconnection in the magnetotail. Many theorists believe that the turbulence in an electron scale current sheet or strongly enhanced cross-tail current is important to provide sufficient dissipation to trigger fast magnetic reconnection in the magnetotail (for example, Drake et al. 1994, Biskamp et al. 1996). While there is observational evidence that the magnetotail current sheet thickness becomes as thin as comparable to the relevant ion inertia length prior to the substorm onset (Sergeev et al. 1990), whether it thins further down to electron scale to initiate reconnection is an open question. It is therefore an interesting question to ask if some dissipation mechanisms set in already at an ion scale current sheet. An aim of our study is to search dissipation processes associated with cross-field current instabilities in an ion scale current sheet which can lead to the fast onset of magnetic reconnection. It is an important point to understand the explosive nature of magnetic reconnection which is still poorly understood (Drake et al. 2001). Recently, a new idea for the quick triggering of magnetic reconnection is proposed (Shinohara and Fujimoto in preparation) based on a 3D full particle simulation. As has already been reported by many authors (Horiuchi and Sato 1999, Lapenta and Brackbill 2002, Daughton, 2002, Silin, I. and J. Biichner, 2003, Scholer et al., 2003), during the non-linear phase of the LHD instability, a thin electron current layer is formed at around the neutral sheet. However, the previous papers did not fully discuss the mechanism of the thin current formation. In this paper,
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we show that it is a very essence of the quick triggering of magnetic reconnection. Actually, the result of a 3D full particle simulation clarifies that the formation of the thin electron current layer enables triggering of magnetic reconnection with large growth rate. SIMULATIONS First of all, 3D simulation results are presented to show that reconnection configuration can really grow within a few ion gyro-periods without any initial perturbation, even in an ion-scale current sheet. We carried out a 3D electromagnetic full particle simulation of an ion-scale current sheet. (For details of our simulation scheme, see Hoshino (1987).) As the initial condition, we make use of the Harris current sheet for simplicity; Sv(z)=80tanh(z/Z>), Hcs(z)=AWcosh2(z/£>), Jy(z)=enCs(Ul+Ue), and Ui/Ue=-Tj£sJTe,cs- D represents the thickness of the current sheet, which we set D=0.5Xj (Xj=c/(>}pi, the ion inertia length). The mass ration is taken to be mi/me=400, electron thermal velocity is ve/c=l/3 (c is the light velocity), and Wpe/Q^l.O. (wpc and Q e are the plasma frequency and the electron cyclotron frequency, respectively.) The simulation box size in the x direction, Fig. 1 (a) Magnetic field lines obtained at Q, t=l.\. Color represents Lx, is taken as LX=\2D that corresponds to the ion flow velocity, Vx. (b) Time history of the maximum value of the wavelength of the fastest growing |5Z|. Thick and thin lines correspond to the 3D and 2D results, mode of the tearing instability, and Ly respectively. almost corresponds to four wavelength of the LHD wave. Periodic boundary conditions are imposed in the x andy directions, while the conducting walls are set at the z boundaries, z=±4D. The number of grids is AfvxAr>.xA/z=384x82x256 (corresponding to x 9 LrxL)xL2=6X,xl.28X,, 4X,;), and ~1.0x]0 particles for both electrons and ions are used in the calculation. It should be noted that no initial field perturbation is added to the initial condition. In the present case, the ion-electron temperature ratio is taken to be Tij£s/Tees=% so that the ions carry most of the cross-field current. Stationary cold background plasma is distributed outside the current sheet; «BK(Z)=A'BK tanh2(z/D), AWAts=0.1, and Te,Bn=TiiBK-Te,c&- If the cold background is uniformly introduced, the ion-ion kink (IIK) instability can grow. (Daughton 1998) To avoid the growth of the IIK instability, we remove the background component from around the neutral sheet. Although there appears a weak pressure imbalance with this background profile, it quickly relaxes without any significant modification of the current sheet structure. In the simulation setup, at least two modes are possibly excited. One mode is the tearing instability, and the other is the LHD instability. Figure la displays a snapshot obtained at Q,t=lA. Magnetic field lines that are almost located on the plane, y=0.64Xi are shown, and these are colored with the ion flow velocity, Vx. Figure la clearly shows magnetic reconnection has already been well developed at Q,t=l.\. At this time, the maximum ion flow velocity exceeds 0.2 VA. (VA is the Alfven speed, VA=(Bo/4itmiNcs)U2) This quick emergence of reconnection configuration is very surprising because it could not be expected from any 2D theories. To illustrate how fast reconnection configuration evolves in the 3D simulation, the time history of the maximum value of \BZ\ is shown in Figure lb comparing with the 2D simulation result (in the x-z plane). Figure lb clearly shows that the significant growth of \BZ\ starts at around £2j t~A in the 3D result, while no reconnection is recognized in the 2D result. -124-
To see what is changed in the 3D from the 2D case, time histories of various parameters are presented in Figure 2. (a) The electric field component, Ey, in the X line (averaged over in the y direction), that is, the reconnection electric field, is drawn in the top panel. The reconnection electric field appears from Q, t ~4, and it is consistent with the increase of the reconnected magnetic flux, (b) To indicate the activity of the LHD instability at the boundary of the current sheet, the maximum value of \Ey\ at z=D is shown. The observed growth rate and the wavelength of the LHD instability are almost consistent with those predicted from the linear dispersion. The LHD instability saturates at Q, t ~4, and the LHD waves are weakened after reconnection takes place, (c) The bottom panel represents the electric current density component, Jy, in the X line (averaged over in the y direction). The thick solid line presents the total current density, Jy, and the thin solid and dotted lines are the ion and electron current density, Jyi, Jye respectively. The ion current density at the neutral sheet decreases as the LHD instability grows, since the plasma in the current sheet diffuses outward by the effect of the LHD instability, while the electron current density is enhanced. The gain of the electron current density exceeds the reduction of ion current density so that the total current density at the neutral sheet Fig. 2 Time histories of various parameters in the 3D result are increases up to twice of the initial value. As shown, (a) Ey at the X-line averaged over the y direction, (b) the shown in the previous simulation results maximum value of |£y| at z=D, and (c) the current density, Jy. concerning the LHD instability, the Thick solid line presents the total current density, Jy, and thin solid enhancement of the current density occurs and dotted lines are the ion and electron current density, Jyi, Jye during the non-linear phase of the LHD respectively. instability (Horiuchi and Sato 1999, Lapenta and Brackbill 2002, Daughton, 2002, Silin, I. and J. Biichner, 2003, Scholer et al., 2003). It should be noted that the timing of the reconnection onset almost corresponds to the enhancement of the electron current density at the neutral sheet. Thus, the 3D simulation results strongly imply that the formation of the electron current layer is a key process to enable the quick reconnection triggering. Hereafter, to address what is the important factor to boost up the reconnection triggering, we concentrate our attention on the formation mechanism of the electron current layer associated with the LHD instability. In order to study the formation of the electron current layer, the 2D simulation runs of the LHD instability in the j - z plane are sufficient and useful. All the setups and the parameters used in the 2D simulation in the y-z plane are the same as those of a 3D simulation noted above (except for the simulation box size). The simulation box size is A(yX.Arz=48x512 (Z,.xZ,2=0.75A,-x8A,,-), and ~1.3*107 particles for both electrons and ions are used. Figure 3 displays a snapshot at the late saturation phase of the LHD instability, 7=5 Q;"1 (=100Q L H~') Figure 3a shows a color contour of the Bx component, and the electric field vectors in the j - z plane are written over. It is clear that waves are quiet around the neutral sheet, while the strong LHD activity is found at the edge of the current sheet. Figure 3b shows a color contour of the Ey component which is dominant component of the LHD instability, and the electron velocity vectors in the y-z plane are drawn. What is an interesting thing here is that electrons around the neutral sheet are accelerated although no strong electric field is found there. At this time, the electron current density at the neutral sheet reaches more than 10 times of the initial value. The electron velocity distribution
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Fig. 3 A snapshot at the saturation phase of the LHD instability, T =3 £2,~'(= 60 £ \ H '), (a) color contour of Bx with the electric field vectors, and (b) color contour of Ey with the electron flow vectors, (c) An electron distribution function observed at the neutral sheet, T= 5 £2,"'(= 100 OLH ') function observed around the neutral sheet is presented in Figure 3c. It shows a slice of the v,-vz plane at vt=0. The magnetic field is almost aligned to the x axis, and By, and 5 2 components are negligible. It is important that the observed distribution function is not gyrotropic, and that the enhanced electron current around the neutral sheet is carried by the non-adiabatic part of the electron velocity distribution. Similar electron velocity distribution functions are also found at the neutral sheet in the 3D simulation result prior to the growth of the tearing instability. We should note that the non-adiabatic electron distribution functions are found everywhere around the neutral sheet (no dependence in the x direction). It implies that the acceleration process is dominated by the instability in the y-z plane. The production of non-adiabatic electrons can be coupled with the tearing mode in the x-z plane since the growth rate of the collisionless tearing instability essentially depends on the electric current density carried by the non-adiabatic particles. DISCUSSION To address the electron acceleration process, we examine some trajectories of the accelerate electron. Two electrons are sampled from the center of the non-adiabatic part of the electron distribution function shown in Figure 3c. Figure 4 shows two particle trajectories both in the real space and the velocity space. The two electron trajectories are drawn by the blue and red lines. The trajectories correspond to the time interval between T=0 and 5 Q,"1. Figure 4 clearly shows that both two particles are meandering around the neutral sheet, and the positions of the particles are confined very close to the neutral sheet, z=0. The result implies that the electron acceleration occurs locally around the neutral sheet. The initial positions and velocities of the accelerated electrons are plotted in Figure 5. The accelerated electrons are sampled from the crescent part of the velocity distribution shown in Figure 3 c. The red and blue dots represent the particle positions (a) and velocities (b) at T=0, and 5 Q,"1, respectively. It shows that most of accelerated electrons already have negative v, and distributed around the neutral sheet at the initial. Actually, most of electrons in the non-adiabatic part are in meandering orbits from the beginning of the run. Figure 6 shows the time history of (a) |vc|, (b) E-v, and (c) Ey that a particle felt. The representation of Fig. 4 Two trajectories of the accelerated electrons, (a) in the blue and red lines is the same as in Figure 4. As the y-z plane, and (b) in the vy-vz plane.
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Fig. 5 The initial positions and velocities of the accelerated electrons are plotted. The accelerated electrons are sampled from the crescent part of the velocity distribution shown in Figure 3 c. The red and blue dots represent the particle positions (a) and velocities (b) 7=0 and 5 Of1, respectively. shown in Figure 6a, the energy of both two particles gradually increases during about 7=1.5-3.5 O;"'. This time interval almost corresponds to the non-linear phase of the LHD instability. Although some strong electric field spikes are found in the interval in Figure 6b, 6c, these spikes are unlikely to contribute to the gradual acceleration of electrons. Since the accelerated electrons are meandering around the neutral sheet, the electrons can be easily accelerated along the y axis even when the electric field is very weak. The electric current distribution is drastically changed during the non-linear phase of the LHD instability so that the non-linear evolution of the LHD instability generates the inductive electric field around the neutral sheet. To check this, the electric current profile and the magnetic field profile in the z direction at 7=3.5 Q,"'are plotted in Figure 7. (The profiles are averaged over in the y direction.) Figure 7 clearly shows that the electric current at the edge of the current sheet is reduced and that the magnetic field profile is changed. The reduction of the electric current at the edge of the current sheet is due to the LHD instability, and the strong current reduction has been confirmed in the previous simulations on the LHD instability in the Harris type current sheet (for example, Brackbill et al. 1984). From the Faraday's equation, AE}JAz= -ABJAt, and the simulation result, Az -0.1 X,, ABX ~0ABn, (at the meandering width), and At - 2 Q,"1, we can easily estimate the order of the induction electric field, Ey, as AEylBo -0.005 X, Q,/c =0.005vA/c =0.00025. The estimated value is two orders smaller than the electric field of the LHD waves. Assuming the meandering electrons are accelerated along the straight orbit by the inductive electric field, m,dvjdt = -eEy, the amount of the acceleration can be estimated with the inductive electric field as Avjc - (QcAt) EyJBa - 0.2 (Ave -0.6 ve). Thus, the evaluation is quite consistent with the energy gain shown in Figure 6. Fig. 6 Time histories of (a) |vj, (b) E-v, and (c) Ey that a (Since the strength of the inductive electric field particle felt, for the particles shown in Fig. 3.
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Fig. 7 (a) The electric current^, profile and (b) the magnetic field Bx profile at r=3.5 Q,"1. is very small and highly fluctuates, it is difficult to directly recognize the DC inductive electric field in Figure 6.) The inductive field, Eyini, profile is shown in Figure 8. Eyilii is averaged over the period between r=2.5-3.5 Qf1 to filter out the large fluctuation of the inductive electric field. Ey^ /.So around the neutral sheet is —1-0.005, and the value is consistent with the estimation noted above. Therefore, we conclude that the meandering electron acceleration is due to the inductive electric field associated with the global evolution of the LHD instability in the current sheet. CONCLUSION In order to fully examine the mechanism for the thin electron current layer formation, we carried out a 2D full particle simulation on the LHD instability. The possible story for the formation of the electron thin current can be summarized as follows. (1) The electric current at the edge of the current layer is reduced by the LHD instability. (2) Associated with the current reduction, the magnetic field penetrates toward the neutral sheet. (3) The inductive electric field generated by the change of the magnetic field profile accelerates meandering electrons around the neutral sheet. (4) The electric current at the neutral sheet is enhanced by the accelerated meandering electrons, and the electron current layer is formed. (5) The reduced current at the edge of the current sheet is compensated by the enhanced electron current around the neutral sheet, and the total electric current in the system is conserved.
Fig. 8 The inductive field EyiBli profile averaged over the period between 7=2.5-3.5 Qf1
Through the acceleration process, the non-adiabatic electrons are efficiently produced. Since the growth rate of the collisionless tearing instability essentially depends on the electric current density carried by the non-adiabatic particles, the increase of the accelerated meandering electrons strongly helps the evolution of the tearing mode. Thus, the formation process of the electron current layer must be responsible for the quick reconnection triggering observed in the 3D simulation. With the above scenario, the quick triggering of magnetic reconnection is possible.
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We think that the proposed scenario contains a key element in a chain of processes triggering magnetic reeonnection in a thicker current sheet. ACKNOWLEDGMENTS I. S. thanks Dr. M. Scholer for his fruitful discussions. This work has been supported by STE Laboratory of Nagoya University via a collaborative research program. Numerical simulations are performed on Fujitsu VPP at IS AS and Nagoya University.
REFERENCES Biskamp, D., et al., Two-dimensional electron magnetohydrodynamic turbulence, Phys. Rev. Lett., 76, 1264, 1996. Brackbill, J. U., et al., Nonlinear evolution of the lower-hybrid drift instability, Phys. Fluids, 27, 2682, 1984. Daughton, W., Nonlinear dynamics of thin current sheets, Phys. Plasmas, 9, 3668, 2002. Daughton, W., Kinetic theory of the drift kink instability in a current sheet, J. Geophys. Res., 103, 29429, 1998. Drake, J. F., et al., Structure of thin current layers: Implications for magnetic reeonnection, Phys. Rev. Lett. 73, 1251, 1994. Drake, J. F., Magnetic explosions in space, Nature (London) 410, 525, 2001. Horiuchi, R. and T. Sato, Three-dimensional particle simulation of plasma instabilities and collisionless reeonnection in a current sheet, Phys. Plasmas, 6, 4565, 1999. Hoshino, M., The electrostatic effect for the collisionless tearing mode, J. Geophys. Res., 92, 7368, 1987. Lapenta, G. and J. U. Brackbill, Nonlinear evolution of the lower hybrid drift instability: Current sheet thinning and kinking, Phys. Plasmas, 9, 1544, 2002. Scholer, M., et al., Onset of collisionless magnetic reeonnection in thin current sheets: Three-dimensional particle simulations, Phys. Plasmas, 10, 3521, 2003. Silin, I. and J. Bilchner, Nonlinear instability of thin current sheets in antiparallel and guided magnetic fields, Phys. Plasmas, 10, 3561,2003. Sergeev, V. A., et al., Current sheet thickness in the near-earth plasma sheet during substorm growth phase, J. Geophys. Res., 95, 3819,1990. Shinohara, I., and M. Fujimoto, Quick triggering of magnetic reeonnection in an ion-scale current sheet, submitted to Phys. Rev. Lett., 2003. Shinohara, I., et al., Low-frequency electromagnetic turbulence observed near the substorm onset site, J. Geophys.
Res., 103,20365, 1998. E-mail address of I. Shinohara [email protected]
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EFFECTS OF GUIDE FIELD IN THREE-DIMENSIONAL MAGNETIC RECONNECTION C. Hashimoto 1, R. TanDokoro1, and M. Fujimoto1 l
Dept. Earth Planet. Sci., Tokyo Inst. Tech., Meguro, Tokyo 152-8551, JAPAN ABSTRACT
We have studied by MHD simulations the effects of the guide field in three-dimensional magnetic reconnection. The guide field is introduced by adding a constant By component Byo to the anti-parallel reconnecting component Bx = tanh(z). An ad-hoc anomalous resistive region, which facilitates the study of reconnection in MHD, is assumed to have a finite extent in the y direction and this gives rise to a three-dimensional situation. It is shown that the guide field makes the U-shaped reconnected field lines as well as the reconnection jet to be inclined from the z axis. The guide field also acts as an obstacle to the jet and this effect produces a pair of helical streamlines in the jet leading part. Time series data from a virtual spacecraft that encounters or skims the jet leading part are found to show good agreement with the observed PTE signatures on either side of the dayside magnetopause.
INTRODUCTION It is needless to say that magnetic reconnection plays an important role in magnetospheric dynamics. Reconnection between IMF and the Earth's field is crucial for the magnetospheric convection in which most of the magnetospheric activities originate. Reconnection in the near-Earth magnetotail is one of the most prominent ways to release the energy stored during silent convection. As such, it is not surprising that numerous simulation studies on magnetic reconnection have been done. The simplest situation to consider is to have anti-parallel magnetic fields and a symmetric density profile across the current layer. With this initial condition, many two-dimensional (2D) as well as three-dimensional (3D) studies with various methods (MHD, hybrid, full-particle) have been done (e.g., Birn et al., 2001). While this may be a good model for the near-Earth reconnection, there is a variety of field line geometry and density profile in the dayside magnetopause. In this study, we deal with the magnetic field that has the guide field in addition to the anti-parallel field. Such cases have been studied extensively in 2D (e.g., Nakamura and Scholer, 2000) but much less in 3D. In this study, before going to the dayside reconnection issue, we will investigate the general large-scale nature of 3D magnetic reconnection with the guide field. 3D MHD simulations are performed in which the three-dimensionality is induced by assuming an anomalous resistive region to have a finite extent in the current-wise direction. Regarding the dayside reconnection, one of the models for FTEs at the dayside magnetopause claims that its peculiar feature is due to the spacecraft's encounter with flux tubes that has been subject to patchy reconnection (e.g., Elphic, 1995; Scholer, 1995). In agreement with this conjecture, data from a virtual spacecraft flying within our simulation box are found to show signatures quite similar to the observed FTEs. MODEL We solve MHD equations with anomalous resistivity. The current sheet of concern has the structure denoted by Bx = Botanh(z), By = Byo{constant). Hereafter the variables will be normalized by the quantities away from the current sheet z=0. The unit for the spatial scale and velocity are the initial halfthickness of the current sheet and the Alfven velocity, respectively. Equal density is assigned on either side of the current sheet and temperature is initially uniform for simplicity. The plasma beta away from the -130-
Fig. 1. Vx component for Byo = 0 (left) and 1 (right). Non-zero Byo makes the plasma jet to be inclined from the z axis.
current sheet is set to 0.5. The anomalous resisitivity region has the spatial form r] = rjo exp[—(x/Rx)2 — {y/Ryf - (z/Rzf] (e.g., Ugai and Shimizu, 1996). Here, Rx = 4, Ry = 10, and Rz = 2. Finite Ry gives rise to a three-dimensional situation.
SIMULATION RESULTS Figure 1 shows the color contours of the Vx component, which represents the plasma jet, for the Byo = 0 (left) and the Byo = 1.0 (right) cases obtained when the reconnection has fully developed (T=70). The top two panels are on the y=0 meridian (xz) plane and those at the bottom are on the cross-section (yz) plane where Bz is maximized, respectively. On the cross-section plots, the in-plane flow components are shown as well. Comparing the bottom two panels, we find the non-zero Byo makes the plasma jet to be inclined from the z axis. By varying Byo, we find more inclination and smaller maximum value of Vx with increasing Byo. The latter nature of the guide field as an obstacle to the jet becomes clear in Figure 2. On the cross-section plot for Byo = 1.0, we can also see a hint of helical vortex flows that will be shown later.
Fig. 2. Field lines for Byo = 1. (a) and (b) show the projections of the field lines onto the x-y plane z = 0 and the y-z plane, respectively. Red, orange, and pink lines show reconnected field lines, while the blue line shows a field line embedded in the current sheet ahead of the jet. The dashed parts are those below the z=0 plane.
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Figure 2 shows the field lines for Byo = 1.0. Panel (a) and (b) show the projections of the field lines onto the x-y plane z = 0 and the y-z plane, respectively. The dashed line parts are those below the current sheet plane z=0. The Red field line shows reconnection between those field lines both of which are initially outside the current sheet. This U-shaped field line is inclined with respect to the z-axis. The orange and pink lines show reconnection between one field line that is initially within the current sheet and the other initially outside. These field lines are reconnected prior to the red lines. Since the field lines prior to reconnection are on the same side of the current sheet, the resultant field line does not cross the z=0 plane. The blue line has been embedded in the current sheet in front of the jet and is not directly processed by reconnection, but is pushed forward by the jet. The curvature of the blue line clearly depicts the tendency of the guide field to slow down the jet. A helical flow pattern is found to emerge at the jet leading part for non-zero By®. By following two typical streamlines, Figure 3 shows the relation between the helical flow and the jet. The pattern for Byo = 1.0 are shown. Panel (a) and (b) are the projections of the two helical streamlines onto the x-y plane z = 0 and the y-z plane x = 0, respectively. The color contours of Vx show the shape of the jet. We find these helical streamlines to propagate to the sides as well as ahead of the jet, more or less wrapping around the U-shaped reconnected field lines (red in Figure 2).
Fig. 3. A pair of helical vortex flows, lauched from the jet leading part, wrapping around the reconnected field lines, are illustrated.
FTE SIGNATURES Patchy reconnection of non anti-parallel field lines can take place at the equatorial dayside magnetopause when the IMF is not due south. Indeed, one of the models for FTE events proposes that the signature in the data is due to encounter with the flux tubes processed in patchy reconnection. Here we see if a virtual spacecraft in the simulation box will obtain the FTE signatures upon encountering or skimming the reconnected flux tubes. To model the dayside situation better, we set the density on the magnetosheath (z > 0) side to be 4 times the magnetospheric value (z < 0). The magnetic field is the same as before, Byo — 1. That is, we model a subsolar equatorial situation when IMF is due east. The transformation from the xyz coordinates to those of LMN is shown in Figure 4. In the LMN coordinates, the magnetic field is (0, —\/2, 0) in the magnetosheath and (\/2, 0, 0) in the magnetosphere. Reconnection proceeds essentially in the same way as the density symmetric case except that the asymmetry makes the jet to be faster on the magnetospheric side. We show in Figure 4 data from a virtual spacecraft sampled within (b and b') or in the outer skirt (a and a') of the reconnected flux tubes. Data are obtained as the spacecraft traverses the flux tube along the magnetopause surface perpendicular to the jet direction -132-
Fig. 4. Magnetic field, plasma density and bulk flow data from a virtual spacecraft that encounters reconnected field lines. The horizontal labels are the y coordinates. The time should run from (left) to right (left) on the magnetospheric (magnetosheath) side. The four trajectories of the spacecraft are depicted in the right-most panel. The shape of the jet is depicted by the green contour lines at | Vx \ = 0.2 on the z = -2.2 ~ 1.4 planes.
(along the y axis) ahead of its leading part (see the Panel on the most right). The horizontal labels are the y coordinates and the time should run from left (right) to right (left) on the magnetospheric (magnetosheath) side. Here we can see the positive and then negative variations of BN in time that characterizes FTE in every case. On the other hand, other variables show different features depending on which side or how deep inside the flux tube the spacecraft is. Comparing the profiles in Figure 4 and those obtained by observations (Paschmann et al., 1982), we find good agreements if we think case a in Figure 4 to correspond to case B in Paschmann et al., b to C, a' to B', and, b' to C\ respectively. Indeed it is only the BM component profiles in b and C that show discrepancy. We suggest that FTEs are consistent with a patchy reconnection picture and the classification of the events arise due to the different penetrating depth of the spacecraft into the reconnected flux tubes.
SUMMARY In this paper, we have investigated by MHD simulations the effects of the guide field in three-dimensional magnetic reconnection. The U-shaped reconnected field lines as well as the jet are inclined from the current sheet normal direction. Furthermore the guide field acts as an obstacle to the jet and this effect produces helical vortex flows as well as a complicated field line structure. The leading part of the jet is highly distorted from an anti-parallel case and this may have significant implication on particle acceleration in this site. Three-dimensional reconnection with a guide field is expected at the dayside magnetopaiise and the flux -133-
transfer events (FTEs) are proposed to be signatures of patchy reconnection. Categorization of the FTE signatures has been done (Paschmann et al. 1982) and we suggest, by showing quite similar data from a virtual spacecraft flying in the simulation box, the classification reflects the penetration depth of the spacecraft into reconnected flux tubes. We anticipate Cluster-II or future multi-spacecraft missions to confirm this conclusion. REFERENCES Birn, J. et al., GEM magnetic reconnection challenge, J. Geophys. Res., 106, 3715, 2001. Elphic, R. C , Observations of flux transfer events: A review, in Physics of the Magnetopause, AGU Monogr. Ser. Vol. 90, AGU, Washington, D. C , 1995. Nakamura, M. S., and M. Scholer, Structure of the magnetopause reconnection layer and of flux transfer events: Ion kinetic effects J. Geophys. Res., 105, 23,179, 2000. Paschmann, G., G. Haerendel, I. Papamastorakis, N. Sckopke, S. J. Bame, J. T. Gosling, and C. T. Russell, Plasma and magnetic field characteristics of magnetic flux transfer events, J. Geophys. Res., 87, 2159, 1982. Scholer, M., Models of flux transfer events, in Physics of the Magnetopause, AGU Monogr. Ser. Vol. 90, AGU, Washington, D. C., 1995. Ugai, M. and T. Shimizu, Computer studies on the spontaneous fast reconnection mechanism in three dimensions, Phys. Plasma, 3, 853, 1996.
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COMPUTER SIMULATIONS ON THE SPONTANEOUS FAST RECONNECTION EVOLUTION IN THREE DIMENSIONS K. Kondoh, M. Ugai, and T. Shimizu Department of Computer Science, Faculty of Engineering, Ehime University, Matsuyama 790-8577, Japan ABSTRACT The dynamics of large-scale magnetic loop in three dimensions is studied by MHD simulations. The spontaneous fast reconnection model is used in this study. In this model, the fast reconnection evolves by the positive feedback between microscopic anomalous resistivities and macroscopic reconnection flows. It is demonstrated that even in the general three-dimensional situation, once a current-driven anomalous resistivity is ignited in a local region in the current sheet, the fast reconnection mechanism spontaneously evolves explosively by such positive feedback. It is remarkable that a fast shock builds up ahead of the magnetic loop in a very limited extent, since the three dimension fast reconnection jet strongly converges toward the magnetic loop top as in the two-dimension case.
INTRODUCTION Theoretical models on the fast reconnection mechanism involving slow shocks have been studied mainly on the basis of magnetohydrodynamics (MHD). We have proposed the spontaneous fast reconnection model, which describes a new-type nonlinear instability in a long current sheet system (Priest and Forbes, 1986). The basic idea lies in the following positive feedback: the self-consistent coupling between (microscopic) currentdriven anomalous resistivity and (macroscopic) global reconnection flow gives rise to simultaneous growth of localized anomalous resistivity and fast reconnection flow by enhancing each other. Recent satellite observations have shown that dynamics of large-scale magnetic loop heating is fundamental for solar flares (Ugai and Tsuda, 1977); also, in the geomagnetic tail, distinct magnetic field dipolarization occurs in accordance with onset of substorms. The earliest version of the magnetic loop simulation was carried out in a closed system, and it was first demonstrated that the fast reconnection jet becomes supersonic and collides with the magnetic loop, leading to a fast shock ahead of the loop top. More sophisticated 2-D MHD simulations on the loop dynamics have extensively been performed for a variety of parameter values in general situations, and the detailed 2D structure of large-scale magnetic loop is demonstrated; in particular, it is demonstrated that a fast shock builds up and impulsively stands in front of the magnetic loop top. Although the 2-D dynamics of a large-scale magnetic loop is well understood, actual systems are of course three-dimensional. It is fundamental to study the large-scale magnetic loop dynamics in three dimensions. The main theme of the present paper is to study this problem by extending the previous 2-D simulation model, limited in the (x, y) plane, to the 3-D model in the z direction. As already shown in a simplified 3-D situation, the fast reconnection mechanism may build up and proceed in a finite extent in the z direction. The fast reconnection jet, associated with the resulting slow shocks, is also confined in the z direction. -135-
Therefore, we are interested in how the spontaneous fast reconnection model works to bring about a 3-D magnetic loop sturcture in the present specific situation. SIMULATION MODELING In the present simulation model, the previous two-dimensional MHD simulations of large-scale magnetic loop dynamics in the (x: y) plane are extended to three dimensions in the z direction on the basis of the spontaneous fast reconnection model. Initiated by a three-dimensional disturbance, magnetic reconnection occurs in a finite extent in the z direction, so that a strong (Alfvenic) 3-D fast reconnection jet may be caused in the limited region. Hence, as the reconnected magnetic field lines are piled up. a large-scale 3-D magnetic loop structure should be formed in a finite extent in the z direction. In understanding precisely the distinct magnetic loop dynamics, numerical computations with sufficiently high-numerical resolution are needed as already seen in the 2-D simulations. With this in mind, the mesh sizes in the x and y directions are taken to be the same as those taken in our previous 2-D model which has numerical resolution sufficiently high to describe definite slow and fast shocks. As an initial configuration, one-dimensional antiparallel magnetic field B = [Bx(y),Q, 0] is assumed as: Bx(y) = sin(7rt//2) for 0 < y < 1; Bx = 1 for 1 < y < Yi; Bx = cos[(i/-Yi)7r/1.2] for Yj < y < Ym(= Yx+0.6); Bx = 0 for y > Ym; also, Bx(y) = —Bx(—y) for y < 0. The plasma pressure P(y) initially satisfies the pressure-balance condition. P + Bx = 1 + /?o, where /3Q is the ratio of the plasma pressure to the magnetic pressure in the ambient magnetic field region 1 < y < Y], so that P(y = 0) = 1 + 0Q initially (in the present study. (3Q = 0.1 is taken); also, constant density p(y) = f is assumed with fluid velocity u = (0, 0,0). The normalization of quantities, based on the initial quantities, is self-evident: Distances are normalized by the half-width of the current sheet do, B by the field strength in the magnetic field region BXQ, and P by ^o/(2/-*o); also, u by VAxo(= £zO/%//•*<>Po); t i m e * by do/VAxo, current density J by J o = BxO/(fiodo), and so forth. Here, the conventional symmetry boundary conditions are assumed on the (x.y), (y,z) and (2,2) planes. Hence, the computational region can be restricted to the first quadrant only and taken to be a rectangTilar box, 0 < x < Lx, 0 < y < Ly. and 0 < z < Lz; also, for simplicity, the conventional symmetry boundary condition is assumed on the outer boundary plane x = Lx. whereas on the other boundary planes (y = Ly and z = Lz) the free boundary conditions are assumed. As in the 2-D model, a current-driven anomalous resistivity model is assumed in the form: T){r,t) = kR[Vd(r,t) - VC] torVd>Vc, = 0 for Vd < Vc, where Vd(r.t) = | J(r, i)/p(r, t) | is the relative electron-ion drift velocity, and Vc may be a threshold for microinstabilities. Here, kj{ = 0.003 and Vc = 4 are taken. Then, as in the 2-D model, in order to disturb the initial static configuration, a localized resistivity is assumed at the origin in the 3-D form, r)(r) = r]oexp[-((x - Lx)/kx)2 - {y/kyf1 — (z/kz)3]. Here, we take kx = kv — 0.8 and T;O = 0.02 in the manner similar to the previous 2-D simulations; also, kz provides the three-dimenisonal effects, and for the smaller kz the three-dimensional effect becomes larger. The basic physics of the 3-D spontaneous fast reconnection evolution is already studied for different values of kz. and it is demonstrated that the fast reconnection mechanism can be fully realized for kz > 4. With this in mind, in the present study kz = 5 is taken. The disturbance is imposed only in the initial time range 0 < t < 4, and the anomalous resistivity model is assumed for t > 4. Hence, the fast reconnection mechanism may be triggered at x = Lx in this model. Because of the symmetry boundary conditions, the computational region can be restricted to the first quadrant only and taken to be a rectangular box. 0 < x < Lx, 0 < y < Ly. and 0 < z < Lz; also, it should be noted that plasma cannot flow across the (y, z) plane at x = 0, where the mirror boundary is placed. It should be noted that sufficiently small mesh sizes are required for precise computations of the spontaneous fast reconnection evolution, so that we assume A.x = 0.04, Ay = 0.015, and Az = 0.1 (Ax and Ay are taken to be quite the same as those in the previous 2-D simulations). Also, we take the magnetic field region size Y\ = 4, and the whole computational region size is assumed to be Lx = 10, Ly = 9.6 and Lz = 9.8.
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RESULTS Figure 1 shows the temporal variation of the electric field Ez measured at the X reconnection point (located at x = 10 and y = z = 0). which may indicate the effective reconnection rate. for the anomalous resistivity models. As in the 2-D model, the fast reconnection mechanism is found to effectively evolve for the anomalous resistivity model in the following manner. After the initial disturbance is removed, magnetic reconnection does not take place, since there is no effective resistivity anywhere. Then, initiated by the disturbance, global plasma flows grow so as to cause notable current sheet thinning to occur at x = 10 in the finite extent | z |< 5, eventually giving rise to ignition of the anomalous resistivity. Once the anomalous resistivity is ignited (at t ~ 10), the reconnection rate drastically grows in the time range 18 < t < 26, when both the current density and the anomalous resistivity simultaneously grow in the X reconnection region because of the positive feedback Fig 1. Temporal variations of the electric between the global (3-D) reconnection flow and the anomalous field Ez at the X reconnection point for the anomalous resistivity model. resistivity; in fact, in the evolutionary stage (18 < t < 26) the global reconnection flows tend to convect the ambient magnetic field into the X reconnection region in both y and 7, directions so as to enhance the anomalous resistivity at the X region. Figure 1 indicates that the fast reconnection rate is sustained for time t > 30. when the (3-D) fast reconnection mechanism, involving slow shocks, is found to build up and be fully established in a finite extent in the z direction even in the present 3-D situation. Hence, in the resulting fast reconnection configuration, slow shocks are extending in the negative x direction from the localized reconnection (diffusion) region (at x = 10), and a large-scale plasmoid propagates toward the mirror plane boundary placed at x = 0 [(y. z) plane], across which plasma cannot flow. When the plasmoid collides with the mirror boundary (for t > 36), the reconnected field lines are piled up there, and a large-scale magnetic loop is formed. Also, the plasmoid collision with the mirror boundary causes a drastic change in the magnetic field configuration near the boundary (at x = 0), which may propagate back toward the reconnection region (at x = 10) at the AlfVen velocity (~ 1), leading to significant influence on the reconnection region. In fact, Figure 1 indicates that the reconnection rate is impulsively enhanced at t ~ 46, which may hence reflect the distinct effects of the plasmoid collision with the mirror boundary with the propagation time delay (~ 10).
Fig 2. Magnetic field (left) and plasma flow (right) configurations at time t — 36 The magnetic loop dynamics should be determined by the self-consistent interaction between the magnetic -137-
loop and the (3-D) fast recormection jet. In order to see the formation and development of the 3-D magnetic loop (or magnetic field dipolarization), Figure 2 shows the resulting magnetic field and plasma flow configurations at time t = 36, when the plasmoid has just arrived at the left (mirror) boundary and the magnetic loop begins to be formed. In this initial phase of the loop formation, the fast reconnection jet, directed to the negative x direction, tends to diverge in the z direction, since the plasma pressure is significantly enhanced near the left boundary. As the fast reconnection mechanism proceeds, slow shocks stand and extend toward the magnetic loop from the X reconnection region (at x = 10). Figure 3 shows the magnetic field configuration at t = 45, when the 3-D large-scale magnetic loop is fully set up. It should be noted that the magnetic field configuration in the loop top has the W-shaped structure, which is quite similar to the one obtained from the 2-D model [see Figure of Ugai(1999)], where such a W-shaped structure results from the supersonic fast reconnection jet in a narrow region between a pair of slow (switch-off) shocks. The W-shaped structure of the resulting 3-D magnetic loop may suggest that a fast shock stands in front of the magnetic loop as in the 2-D case. Hence, it is important to examine in the present 3-D situation how the fast shock can be realized. Fig 3. Magnetic field configuration at time t = 45. In fact, we find that the fast reconnection jet becomes supersonic (or superfast), and a definite fast shock stands in front of the magnetic loop in the time range 40 < t < 50. For instance, Figure 4 shows the profiles of quantities along the x axis at time t = 45, which indicates that a definite fast shock stands at x ~ 2.2; in fact, the fast reconnection jet ux becomes larger than the sound speed cs = y/P/p and hence supersonic ahead of the fast shock (x > 2.2); also, across the fast shock the fast reconnection jet is reduced to the subsonic flow, whereas the plasma density p and pressure P increase almost discontinuously. The fast shock profile along the x axis, which crosses the center of the 3-D fast shock, is in good agreement with the one already shown in the 2-D model. In this respect, it should be noted in Figure 3 that the fast reconnection jet is significantly confined in the z direction and notably converges toward the magnetic loop top, so that the extent of the resulting fast shock is distinctly limited Fig 4. Profiles of quantities along the x axis both in the z and y directions. {y = z = 0) at time t = 45 REFERENCES Priest, E. R., and T. G. Forbes. New models for fast steady state magnetic reconnection, J. Geophys. Res. 91, 5579 (1986). Ugai, M.. Computer Studies on the spontaneous fast reconnection model as a nonlinear instability. Phys. Plasmas 6, 1522 (1999). Ugai, M., and T. Tsuda, Magnetic field-line reconnection by localized enhancement of resistivity, J. Plasma Phys. 17, 337 (1977).
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SUPERSONIC AND SUBSONIC EXPANSION ACCELERATION MECHANISMS IN FAST MAGNETIC RECONNECTION T. Shimizu and M. Ugai
Department of Computer Science, Ehime University, Matsuyama City, 790-8577 Japan
ABSTRACT The thermodynamic expansion acceleration mechanism associated with the fast magnetic reconnection is studied by two-dimensional magnetohydrodynamic (MHD) simulations and the Rankine Hugoniot analysis. In the acceleration mechanism caused in the reconnection jet region, the jet generated by a pair of slow shocks can be further accelerated by the expansion and propagation of the plasmoid formed by the reconnection process. It is remarkable that the resulting jet can exceed the Alfven speed measured in the upstream magnetic field region. This acceleration mechanism is classified into supersonic and subsonic cases. Either of those cases occurs depending on the upstream field conditions of the jet. The former is similar to the Parker's solar wind acceleration mechanism, which includes two-dimensional expansion process, but the latter is basically one-dimensional expansion process. This paper consists of two topics. Firstly, the upstream field conditions required to cause the supersonic case is shown. Secondly, it is shown that the plasma expansion acceleration in the supersonic case is stronger than the subsonic case because the supersonic case has the circulating enhancement mechanism of the expansion process while the subsonic case does not. INTRODUCTION Fast magnetic reconnection provides a physical mechanism, by which magnetic energy is explosively converted into plasma kinetic and thermal energies. Petschek predicted that when the magnetic reconnection process includes slow shocks extending from the magnetic diffusion region, the reconnection process becomes significantly active, giving rise to high-speed plasma jets (H.E. Petschek, 1964). In his model, the plasma jet can reach the Alfven speed measured in the upstream magnetic field region. However, according to recent numerical MHD study (T.Shimizu and M.Ugai, 2000), the reconnection jet can steadily exceed the Alfven velocity. The plasma acceleration is caused by the adiabatic expansion acceleration mechanism. Es-139-
pecially, when the jet is supersonic, the acceleration mechanism is basically the same as the Parker's solar wind acceleration mechanism, excepting the gravity force. In other words, when the reconnection jet generated by slow shocks is supersonic and expands in the direction normal to the jet, the reconnection jet can be adiabatically accelerated in the downstream region of the slow shock. The resulting jet can exceed the Alfven speed. In this acceleration mechanism, the expansion of the jet is naturally caused due to the swelling of the plasmoid formed in the downstream of the jet. On the other hand, also when the jet is subsonic, the adiabatic expansion acceleration works. Then, the jet can slightly exceed the Alfven speed. First of all, it is important that whether the reconnection jet generated by slow shocks becomes supersonic or subsonic. It depends on the plasma, density and beta value in the upstream magnetic field region of the jet. In this paper, the dependence is shown, including the asymmetric reconnection models. Subsequently, the difference between the supersonic and subsonic expansion acceleration cases is discussed. RANKINE HUGONIOT ANALYSIS OF THE SLOW SHOCK ACCELERATION On the basis of the Petschek model, let us show the required condition to generate supersonic jets. For the simplicity of the analysis, consider when the separatrix angle of the reconnection jet is extremely narrow. At the time, the slow shock becomes a switch off shock in symmetric reconnection case in which both side upstream magnetic field regions have the same plasma conditions and field intensity. The sound Mach number Ma2 of the reconnection jet is estimated by the Rankine Hugoniot relation, as follows. M22 = u22/C2s2 = 2/( 7 /?! + 7 - 1),
(1)
where u2 is the downstream jet speed, Cs2 is the sound speed, 8\ is the beta value in the upstream region. Now, assume the specific heat ratio 7 = 2. When ,81 is less than 1/2, the jet is supersonic, i.e. superfast. Note that the sound Mach number is close to the fast wave Mach number because the reconnection jet is generally high betaConsidering asymmetric reconnection models in the similar manner, the Mach number is given as a set of the following equations.
M22 = (-a + AxlIAx2)2Ax2l((h + 1 - a2), Ax2 = (2ft + 1 - a)/(2fr + 2), A r l =(2/? 1 +2-a-a 2 )/(2 / 3 1 +2).
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(2)
(3)
(1 - aa)2/pla/(2(lla (l-ab)2/plb/(2f3lb ua/ab
+2-aa-a2a) = 2 + 2-ab-a b).
= y/(pla
+ l)/(/3lb + 1).
(4)
The index "a" and "b" respectively correspond to the upper and lower side regions of the current sheet. When f3\a, /3ib and pia/pib are given, aa and ab are obtained with Eq.(4). Note that either index "a" or "b" must be given to every value in Eqs.(2) and (3). When index "a" is selected, MS2a is obtained. Similarly, set index "b" for MS2b- As a result, there are four possible cases for the reconnection jets; case-A: M s 2 a > 1 and MS2b > 1; case-B (and C): M32a < (>)1 and Ms2b > (<)1 and case-D: Ms2a < 1 and Ms2b < 1. For simplicity, let us assume pia — pu,. Fig.l shows the contour map of the sound Mach number Ms26- The oblique dash line corresponds to the symmetric case (flia = /?i&), so that Eq.(l) is applicable. The contour line of Ms2b = 1 starts at (3\a = 0 and fl\b = 1, and then, asymptotically becomes close to flu, = 0 a t {5\a = +oo. This map is consistent with MHD simulations in which cases 1 to 4 in Fig.l were examined.
Figure 1: The Mach number MS2b of the asymmetric reconnection jet is determined by the beta values f3\a and f3\b, when p\a = pu, is assumed. M22b for the lower side jet is shown as a contour map. For upper side jet, simply exchange "a" and "b".
THE DIFFERENCE BETWEEN SUPERSONIC AND SUBSONIC EXPANSION ACCELERATIONS Fig.2(a) schematically shows the downstream region of the reconnection process including the supersonic adiabatic expansion acceleration mechanism. From the left side to the right, Petschek reconnection jet region, a supersonic expansion -141-
Figure 2: (a): Fast reconnection including the supersonic expansion acceleration like Case 1 in Fig.l. (b): The subsonic expansion acceleration case like Case 2
acceleration region and a plasmoid are located. Two pairs of slow shocks are respectively formed around the Petschek region and plasmoid. In the first step of the acceleration, a supersonic jet is generated in the Petschek region. Second, the jet is further accelerated by the supersonic adiabatic expansion process in the middle region. Note that the jet expands in the direction normal to the jet, due to the expansion of the plasmoid. Third, the jet encounters a fast shock at the plasmoid (point A). As the fast shock is enhanced by the accelerated jet, the expansion of the plasmoid is enhanced, returning to the enhancement of the adiabatic expansion acceleration of the jet. Hence, a circulating enhancement mechanism is formed. Fig.2(b) shows the subsonic case. The slow shocks cannot generate a supersonic jet. The resulting subsonic jet is further accelerated by the subsonic adiabatic expansion process in the middle region. At the time, the jet converges toward the center of the jet region because this expansion process is driven and maintained by fast expansion waves generated from the propagating plasmoid pushed by the slow shocks around the plasmoid. Since it is just one-dimensional expansion process, there is no circulating enhancement mechanism, unlike the supersonic case. REFERENCES Shimizu,T. and M. Ugai, Adiabatic expansion acceleration mechanism of superfast jet in the spontaneous fast magnetic reconnection model, Phys. Plasmas, 7, 2417 (2000). Petschek H.E., Magnetic field annihilation, AAS-NASA symposium, NASA Special Publ. SP-50, 425 (1964).
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INSTABILITY AT THE LEADING EDGE OF A RECONNECTION JET R. TanDokoro1 and M. Fujimoto1 l
Dept. Earth Planet. Sci., Tokyo Inst. Tech., Meguro, Tokyo 152-8551, JAPAN ABSTRACT
In a transient reconnection process, the leading edge of the reconnection jet pushes and compresses the plasma standing ahead of it. This causes curved magnetic field lines and a strong pressure gradient to develop at the edge. It implies that the leading edge of the jet could be unstable to the ballooning instability. We studied this situation by three-dimensional MHD simulations and clarified that the interface indeed becomes unstable. The leading edge on the current sheet plane is deformed into a wavy shape and then to a mushroom-like pattern subsequently. The growth rate of the instability is controlled by the wavelength in the current-wise direction with a shorter wavelength mode growing faster. The dispersion curve obtained from a series of simulations is given.
Introduction Magnetic reconnection is one of the most important processes in space physics, in which magnetic energy is converted to kinetic energy of plasma (e.g., Priest and Forbes, 2000). In the process, plasma as well as magnetic field outside the current sheet is converted toward the diffusion region located within the current sheet and a pair of hot jets are ejected out of the region along the current sheet. In a transient case where reconnection is initiated at a certain time, there stands ahead of the accelerated jets the pre-existing current sheet plasma at rest. The interaction between the reconnection jet and the standing plasma is one of the key issues in space plasmas, whereby some visible effects of the energetic phenomena in the magnetotail, the solar corona, etc., are mediated (e.g., Haerendel, 1992; Masuda et al., 1994). In the interaction, curved field lines and a significant pressure gradient are produced at the jet front. It is quite likely that the interface is unstable to a kind of interchange instability, or the ballooning instability to be specific, when the third-dimension along the current is taken into account. The saturated state of the two-dimensional tearing instability, in which a similar situation arises, is shown by Dahlburg et al. (1992) to be unstable to three-dimensional disturbances. A full particle simulation of a reconnection jet has shown that a highly-localized strong plasma flows develops due to coupling to an interchange and a kink modes (Pritchett and Coroniti, 1997; Pritchett, 2001), whose wavelengths are order of the ion gyro-radius. Coupling to an interchange mode is reported also in hybrid simulations (Nakamura et al, 2002a, 2002b). In this paper, we show results from three-dimensional MHD simulations of reconnection to demonstrate that the instability indeed exists and show its dispersion relation within the MHD approximation obtained from a series of simulations. Simulation Model In this study, we have conducted three dimensional MHD simulations to examine the behavior of the leading edge of the reconnection jet. The well-known MHD equations with anomalous resistivity fixed in space and time are used. We use the usual coordinates system in magnetospheric physics: The x axis along the tail axis directed positive (negative) earthward (tailward), the y axis directed positive from dawn to dusk, and the z axis positive northward. The initial magnetic field in our simulation is anti-parallel B(z) = tanh(2)e*x. Here normalizations are follows: magnetic field by the the lobe magnetic field Boplasma density by the lobe density no, spatial scale by the half thickness of the current sheet D, velocity -143-
Fig. 1. The Bz plots on the current sheet plane (z=0 plane) at T=60, 90, 120,and 150. The jet front becomes wavy and subsequently grows into a bubble-like pattern.
by the lobe Alfven speed VA< a n d time by the Alfven transit time T = D/VA- The plasma temperature is assumed to be uniform with the plasma beta in the lobe set to 0.5. Magnetic reconnection is initiated by the localized anomalous resistivity in the current sheet, whose spatial form is r\ = Tftexp [— (x/D)2 — (z/D)2]. i]o = 0.03 is chosen. Note that the resistivity does not change in the dawn-dusk (y) direction. The instability is seeded by adding perturbations to the initial equilibrium density and pressure. The perturbations have the form 5Di
Fig. 2. The Bz plots on the current sheet plane (z=0 plane) at T=100. The panels show the results from various runs in which wavelength are changed: (a)Ly=5, (b)Ly=10, (c)Ly=20, and (d)Ly=40.
To measure the growth rates 7 quantitatively, we have measured the quantity ABz(t) = max{Bz^max(x, t) — Bz^min(x,t)}, where Bz^max(x, t) and BZtinin(x,t) are the maximum and the minimum value, respectively, of the Bz component along a given constant x on the z=0 plane at time t. Having found that this ABz(t) to show linear growth in the small amplitude phase for various cases, we have decided to adopt the growth rate of this quantity as the growth rate of the instability. We note that, because the background field is changing in time at a pace comparable to the growth rate of the instability, we would think that this experimental rather than a theoretical determination of the growth rate is a reasonable choice at present. Figure 3 shows the dispersion relation of the instability obtained in this way. It can be seen that the larger wavenumber (shorter wavelength) modes have larger growth rates. The increase of the growth rate is more or less saturated around ky ~ 1 presumably due to the finite width effect of the jet's leading edge. This saturation makes the growth rate range to be limited within a factor of 2 even though the wavenumber is varied for about two-orders of magnitudes. Discussion We have shown that, in a transient reconnection process, the interface between the jet and the current sheet plasma standing ahead of it is unstable to an interchange instability. The shorter wavelength mode growing faster than the longer wavelength mode is a typical feature of interchange instabilities. The geometry of the interface in which curved field lines are collocated with a sharp pressure gradient makes us to conclude that the instability is the ballooning instability. This conclusion is supported by the results from other runs in which the plasma density at the current sheet (z = 0) is reduced while keeping the plasma pressure the same. The growth of the instability is faster when the initial density at z = 0 is lower. The change in the initial density profile changes the density structure at the jet leading edge but the pressure structure stays essentially the same. The density gradient at the jet leading edge develops less for lower initial z = 0 density. This implies that it is the steep pressure gradient but not a density gradient that is crucial for the instability. Unlike the Raleigh-Taylor instability, which is the most probable alternative candidate that requires Vn, only the Vp effect is crucial for the ballooning. Thus the results are in favor of the ballooning mode. Because of the velocity difference between the fore-going part of the mushroom-like structure and those left behind, the jet's leading part is elongated. When the elongation provides enough space for a velocity shear instability, one would expect this secondary instability to develop at the sides of the elongated spines. Our preliminary results show this to be true. The growth rates differ only by a factor of two when the wavenumbers are varied over the two-orders of magnitudes range. This implies that the results may depend heavily on the initial power spectrum shape of the perturbation. We have shown the shortest wavelength mode to dominate when a run is started from a -145-
Fig. 3. The dispersion relation 7 versus ky of the instability obtained from a series of single-mode simulations.
flat spectrum. The results from a run may change significantly when a power-law spectrum is considered for the initial perturbation, in which a long wavelength mode may dominate if the spectrum is steep enough. Understanding how the resultant field line structure would change according to the initial spectrum is important, for the structure should have significant impacts on particle energization processes taking place at the site. We have studied only the cases without a guide field. The guide field effects could be either ways, bringing more complication into the interface structure up to a moderate value but suppressing the instability beyond a certain value. This issue will be studied in a near future.
Acknowledgments R. T. thanks Dr. M. S. Nakamura for fruitful discussion and useful comments. M. F. is a member of the ACT-JST project 12D-1.
REFERENCES Dahlburg, R. B., et al., Secondary instability in three-dimensional magnetic reconnection, Phys. Fluids B, vol. 4, 3902, 1992. Haerendel, G., Disruption, ballooning, or auroral avalanche - on the cause of substroms, Proc. ICS-1, ESA, Kiruna, 417, 1992. Masuda, S., et al., A loop-top hard X-ray source in a compact solar flare, Nature, 371, 495, 1994. Nakamura, M. S., M. Fujimoto, and H. Mastumoto, Instability at interface between reconnection jet and pre-existing plasma sheet, Adv. Space Res., 29/7, 1125, 2002a. Nakamura, M. S., H. Matsumoto, and M. Fujimoto, Interchange instability at the leading part of reconnection jets, Geophys. Res. Lett, 29, 10.1029/2001GL013780, 2002b. Priest, E., and T. Forbes, Magnetic reconnection, Cambridge Univ. Press, Cambridge, 2000. Pritchett, P.L., Collisionless magnetic reconnection in a three-dimensional open system, J. Geophys. Res., 106, 25,961, 2001 Pritchett, P. L., and Coroniti, F. V., Interchange and kink modes in the near-Earth plasma sheet and their associated plasma flows, Geophys. Res. Lett., vol. 24, 1997.
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SOLAR FLARES AND MAGNETIC RECONNECTION T. Yokoyama National Astronomical Observatory, Nobeyama, Minamimaki, Minamisaku, Nagano, 384-1305 Japan
ABSTRACT
Recent studies by the author and his colleagues on solar flares based on the magnetic reconnection model are reviewed. The studies have been done mainly by using the data sets of space-craft observations and supercomputer simulations. The covered topics are an EUV observation of reconnection inflows, MHD simulations of the reconnection model, and a microwave observation of freely-streaming high-energy electrons. The universal scaling law to explain the relation of the temperature and the emission measure of the solar/stellar flares is also shown. INTRODUCTION — THE MAGNETIC RECONNECTION MODEL OF FLARES Magnetic reconnection is now considered to be one of the best mechanisms which explain solar flares. Many pieces of evidence supporting the magnetic reconnection model of flares (Figure 1) have been found by the recent space-craft observations. The model tells that the flare energy is primarily released high in the corona and is transported by heat conduction from the magnetic X-point along the field lines. The propagating conduction fronts thus forms a cusp-shaped hot region. The observed soft X-ray cusp loops are the manifestation of this process (Tsuneta et al. 1992, Tsuneta 1996). The observed soft X-ray ejecta associated with flares (Shibata et al. 1995; Ohyama & Shibata 1997; Tsuneta 1997) may be one of the bi-directional reconnection jets. The other jet goes down and collides with the closed loop at the base, and forms a fast-mode MHD shock. The hard X-ray source above a soft X-ray loop found by Masuda et al. (1994) may correspond to this shocked region. When the released energy is transported by conduction to the upper chromosphere, the dense plasma is suddenly heated up and expands. The induced pressure-gradient force drives the plasma up to the corona. This process is called the chromospheric evaporation. The blueshifted features of X-ray lines observed, e.g., by the Yohkoh Bragg Crystal Spectrometer is the manifestation of this line-of-sight motion of up-going plasma. More recently, Czaykowska et al. (1999) found evidence of spatially resolved signatures of up-going plasma by using the SOHO/CDS data in the gradual phase of a two-ribbon flare. -147-
Figure 1: Schematic illustration of the reconnection model of a flare. Thick solid lines show magnetic field lines. In this paper, recent studies by the author and his colleagues related to the magnetic reconnection model of solar flares are reviewed. For more comprehensive reviews of the recent progress in the solar flare physics, the readers are expected to refer to Trimble & Aschwanden (2000, 2001, 2000, 1999). The second section shows an BXJV observation of reconnection inflows that might be one of the most important findings to support the reconnection model. For quantitative comparison with the observations described above, the model has to be improved. The MHD simulations of a flare including the heat conduction and the chromospheric evaporation effects shown in the third section are results of one of such efforts. A scaling law between the temperature and the emission measure of solar/stellar flares derived on the basis of the simulation results is shown in the fourth section. The acceleration mechanism of high-energy particles is still a big problem of the solar flare physics. The fifth section is for the finding of rapid propagation of non-thermal microwave sources that are emitted from the freely-streaming electrons accelerated at the reconnection site. The final section is for a summary and discussions in which unsolved problems of the solar flare physics are shown. OBSERVATION OF RECONNECTION INFLOWS IN A SOLAR FLARE The observations described in the previous section support the reconnection model but are indirect evidence in the sense that we are not seeing the energy release process itself. So the search for the inflow or reconnection jet is still ongoing. McKenzie & Hudson (1999) found downward plasma motion above an arcade formed in a long-duration flare and coronal mass ejection. They suggest that this motion is a consequence of the retraction of the reconnected field lines, i.e. the outflow from the magnetic reconnection point. More recently, we successfully observed a reconnection inflow in a flare that is another important piece of evidence for the reconnection model (Yokoyama et al. 2001). The flare occurred on the north-east solar limb. According to the Yohkoh soft X-ray observation, it showed a nice cusp-shaped loop and a plasma blob (plasmoid)
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Figure 2: (left) A snap shot of EUV (Fe 195 A) image by SOHO/EIT. The field-of-view size is 500 arcsec x 500 arcsec (« 350.000 km x 350,000 km). (right upper) Time evolution of 1-dimensional distribution of EUV intensity along the thick solid line in the left panel, (right bottom) Soft X-ray light curve by Yohkoh/SXT in arbitrary units. ejection. The EUV observation of the same flare by SOHO/EIT shows a bubble-like void ejection. The core of this EUV void corresponds to the soft X-ray plasmoid and is bright in X-ray but dark in EUV because of its temperature « 4 MK. As this void moves away from the limb, the lower part of the void becomes thinner and thinner. When the void finally detaches from the limb, an X-shaped structure appears at the detaching point. We believe a magnetic reconnection process occurs at this detaching point because of the following reasons: First, it is natural to consider that the void is a bubble of plasma surrounded by magnetic field lines. As the void detaches, the surrounding lines are also pulled up and at the detaching point the anti-parallel field lines meet. Second, as the void continues rising, the pattern of threads (which we assume that the threads are field lines) merge toward the 'X! point. Figure 2 shows the evolution of the slice cut across the 'X:-point. A clear merging pattern toward the 'X' line can be seen. We believe this is the inflow of the magnetic reconnection. Third, as a consequence of this process, reconnected field lines should shrink down below the 'X' point. We observe an evolving cusp-shaped loop on the limb in soft X-ray and EUV which is the result of the piling up of the shrunk magnetic field lines (Tsuneta et al. 1992; Hiei et al. 1993; Forbes & Acton 1996). Through these observations, we conclude that we have finally found the last link demonstrating the reality of ongoing magnetic reconnection process. The measurement of the inflow velocity is done by tracing the movement of the thread-like patterns in the
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EIT movie. We measure the speed of the movement in a slice cut indicated in the left panel of Figure 2. The speed of the incoming pattern is about VpatteTn = (1.0 - 4.7) km sec"1 over the period 4:006:00 TJT (right upper panel of Figure 2). We also estimated the reconnection rate, M\ that is defined as the ratio of the inflow speed to the Alfven velocity. The Alfven velocity is estimated from the soft Xray observations as follows. The emission measure and the temperature are derived from the filter ratio method as EM = 1030 - 1031 cm"5 and T = 2.7 - 4.2 MK, respectively. If we assume the depth (L) along the line of sight 1.5 x 105 km (that is approximately the same size as the loop length), density is n = (0.8-2.6) x 1010 cm"3[L/(1.5 x 105 km)]" 1 / 2 . Thus the total thermal energy is Eth = 3nk^TL3 = (0.3 - 1.5) x 1031 erg. where /CB is the Boltzmann constant. Using the lifetime of this flare r « 3 x 104 sec, we obtain the output rate of the thermal energy as E^/T = (1-5) x 1026 erg sec"1. Balancing the energy output is the released magnetic energy input rate, expressed as 2_B2VjnflowZ/2/(47r) with unknown magnetic field strength B. Equating the increasing rate of thermal energy and the released magnetic energy rate, we obtain the magnetic field strength of B = (2.4 - 12) G. Using these values, the Alfven speed becomes CA = (160 - 970) km sec" 1 . Since the inflow speed is V^nflow = (1.0 - 4.7) km sec"1, the reconnection rate is ^inflow/CA = 0.001 - 0.03. This value is roughly consistent with the Petschek's (1964) reconnection model not with the Sweet-Parker's model (Sweet 1958; Parker 1963), because the latter predicts much less value (MA = .Rm ~ 10"7) than this observation (MA = 0.001 - 0.03) since the magnetic Reynolds number defined by the Spitzer-type resistivity is extremely large (Rm « 1014) in the corona. Note, however, that we found no evidence of MHD slow-mode shocks the formation of which is the key of the Petschek's model. Hence, other mechanisms such as anomalous enhancement of magnetic diffusion above the Spitzer values by, e.g., MHD turbulence and/or plasma micro-instabilities cannot be ruled out.
MHD SIMULATIONS OF A SOLAR FLARE For quantitative comparison with the observations, the model also has to be improved. The MHD simulations of a flare including the heat conduction and the chromospheric evaporation effects shown in this section are results of one of such efforts. We study the reconnection by solving numerically the two-dimensional MHD equations with nonlinear anisotropic heat conduction effect (Yokoyama & Shibata 2001; see also Chen et al. 1999). Figure 3 shows the results. Because of the enhanced resistivity, the point at z = 20 in the current sheet evolves to an X-point of the magnetic reconnection. The reconnected field lines together with the frozen-in plasma are ejected as upward and downward jets from this X-point due to the magnetic tension force. The velocity of the reconnection jets Vjet = 1.5 - 2.2 is approximately the Alfven speed of the inflow region. A pair of inflows take place from positive and negative rr-directions of the X-point. The velocity of the inflows is V{n « O.ICA- Slow-mode MHD shocks are formed at the boundary between the inflows and the outflows. These characteristics, such as jets, inflows, slow-mode MHD shocks are consistent with those predicted by Petschek's (1964) model. In the temperature distribution in Figure 3, a clear cusp shape is seen. The outer edge of this cusp is a pair of the conduction fronts propagating from the X-point. The conduction fronts transport the thermal energy, which is released at the slow-mode shocks in this region, along the field lines and form the cusp shape of hot plasma. Note that the global cusp shape of the hot region behind the conduction front is similar to the soft X-ray solar flare loops in long-duration (LDE)
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Figure 3: Results of the simulation of a flare. Left and right panels show temperature and density distribution, respectively. The arrows represent the velocity, and lines show the magnetic field lines. The unit of length, velocity, time, temperature, and density is 3000 km, 170 km s" 1 , 18 s, 2 x 106 K, and 109 cm" 3 , respectively. flares (Tsuneta et al. 1992). Since the conduction front propagates faster than the inflow of the shock, the temperature jump between the inflow and the outflow is smoothed away. The jumps in the density distribution in Figure 3 are, therefore, isothermal slow-mode MHD shocks. Thus, it can be said that the adiabatic slow-mode MHD shocks are dissociated into conduction fronts and isothermal shocks by the heat conduction effect (Forbes et al. 1989). The reconnected magnetic field which goes down from the X-point finally forms a closed loop. At the top of this loop, the reconnection jet collides with the loop plasma. A fast-mode MHD shock is formed and the jet plasma is compressed behind this shock. The dense blob at (x, z) = (0,11) is a product of this compression. In the density distribution, a plasma mound can be seen. This is a direct consequence of the chromospheric evaporation. The plasma in the upper chromosphere is suddenly heated up by the penetration of the heat conduction front. As a result, the gas pressure in this region is increased to drive a back flow toward the corona, i.e. the evaporation flow. This flow carries the dense plasma into the corona. The velocity of the evaporation flow is 0.2 - 0.3 times the local sound speed. The density of the evaporated plasma is « 10 times the initial coronal density at maximum and « 5 times in average.
SCALING LAW OF THE FLARE TEMPERATURE One of the interesting questions in flare physics is what determines the flare temperature ? We proposed a model to this question based on the simulation results of the magnetic reconnection model described in the previous section (Yokoyama & Shibata 2001) and it is extended to a universal scaling law to explain the correlation of the temperature versus the emission measure of solar and stellar flares by Shibata k. Yokoyama
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Figure 4: The theoretical EM-T relations overlaid on the observed EM-T relation based on recent observations with ASCA (Koyama et al. 1996, Tsuboi et al. 1998. Hamaguchi et al. 2000, Hamaguchi 2001, Ozawa et al. 1999, Ozawa 2000) and Chandra (Imanishi et al. 2001). Four theoretical curves are plotted for different magnetic field strengths B = 5, 15. 50. and 150 G. The flare loop length L = constant curves are also plotted with dash-dotted lines for L = 108,1010,1012 cm. (From Shibata & Yokoyama 2002) (1999; 2002). If we assume that the input of the energy to a loop balances with the conduction cooling rate, the temperature at the loop apex is T « (2QL2/KQ)2!1 where Q is the volumetric heating rate, L is the half-length of the loop, and KQ is the constant part (~ 10~6 in CGS unit) of the conduction coefficient. In our simulations, the heating mechanism is magnetic reconnection so that the heating rate is described as Q =
B 2 /(4TT)
• V\n/L • l/sin#, where V^n is the inflow velocity (~
Petschek's theory, where CA = B/^/4irmno, «0i
m are
O.ICA
from our result and also from
Alfven velocity, density of the background corona,
and the mean particle mass, respectively), and 8 is the angle between the slow-mode MHD shock and the loop and is approximately given by sin# w V{n/C\. By manipulating the equations, we find TJ
B
3
" ) \
(i)
\2nnoy4nmno J In this model, after the reconnection occurs, the released energy is transported to the top of the chromosphere by heat conduction, and heats the chromospheric plasma suddenly. Then its plasma pressure increases enormously, and drives the upward flow back into the (reconnected) coronal magnetic loop, creating hot dense flare loops. As a result of the chromospheric evaporation, the flare loop density increases to n ( » no). This evaporated plasma with density n is the source of X-ray emission; EM ~ n2L3,
(2)
where EM is the emission measure defined as EM = / n2dV and is roughly proportional to the X-ray flux of a flare. Here, we assumed the volume of a flare is approximated by V ~ L3. Since this evaporated plasma has high gas pressure, we have to assume that the magnetic pressure of the reconnected loop must be larger -152-
Figure 5: NoRH images of the microwave emission at 17 GHz at selected times. Solid contours indicate the level of Tb = 5.6, 10., 18. x 10s K at 17 GHz. Note that the first signature of the propagation is seen at 00:56:19.660 UT and it reaches the other loop end at 00:56:20.160. The distance between these sources are 4.5 x 104 km. Thus, the apparent velocity is « 9 x 104 km s" 1 . than the gas pressure of evaporated plasma to maintain stable flare loops. From this, it may be reasonable to assume
2n* B T~|J,
(3)
as a first approximation, (kn is the Boltzmann constant.) In fact, some previous observations of post flare loops showed that the gas pressure of the flare loop is as high as the inferred magnetic pressure (e.g., Tsuneta 1996). Eliminating n and L from above equations, we find
This relation can explain the correlation of the temperature versus the emission measure of solar and stellar flares in Figure 4. It is also interesting to note here that this scaling law of equation (1) is found to be applicable even in the laboratory plasmas in the scale size of 1 meter (M. Brown from Swarthmore College, private communication). RAPID PROPAGATION OF NON-THERMAL SOURCES It is known that some amount of the released energy of the solar flare goes to the acceleration of particles to non-thermal levels. (For the comprehensive review of the recent studies on the particle accleration, see, e.g. Priest & Forbes 2000.) The free streaming of the electrons in the corona has been recognized from observations at meter wavelengths. The so-called type III radio burst (e.g. Bastian et al. 1998) is thought to be caused when the streaming electrons travel along an open magnetic field. The speed of the electrons is estimated to be a fraction of the speed of light from the shifting rate of the emitting frequency which is related to the background plasma density of the corona, based on the model atmosphere where the density is given as a function of a height. There are a few examples of observation which succeeds in detecting the motion of the accelerated electrons as a series of images. Motion at 3000 km s"1 was reported by Bastian
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et al. (1994) using VLA observations at 8.4 and 15 GHz with a time resolution of 0.2 s. White, Janardhan, & Kundu (2000) found the propagation of a non-thermal disturbance at a speed of 26,000 km s"1 at 0.33 GHz. Recently we observed a propagation of non-thermal microwave emission sources inside an extended flaring loop (Yokoyama et al. 2002) with the Nobeyama Radioheliograph (NoRH; Nakajima et al. 1994). The main microwave structure of the flare consists of a point source and an elongated one (Figure 5). From the spatial coalignment with the magnetograms obtained by SOHO/MDI, it is shown that the point source is located near a sunspot. And the opposite polarities at the two ends of the elongated source suggest that this is a magnetic loop. We can interpret the emission from the elongated source to be due to the non-thermal gyrosynchrotron mechanism based on the following diagnostics. The gyroresonance emission is ruled out because the emission comes from the site where the magnetic strength is not strong enough to emit 17 GHz microwave. The thermal free-free mechanism is also ruled out because the brightness temperature Tj is too intense to be explained by such thermal emission (e.g. Dulk 1985). The peak intensity was Tj = 8.7 x 106 K at 00:56:42 UT. Therefore, the most probable emission mechanism is non-thermal gyrosynchrotron emission. We find several propagation features of the non-thermal sources from the south-east end to the north-west end of the elongated loop. They are clearly seen in a movie of the 0.1-second cadence from around 00:56:10 UT near the peak of the event. Figure 5 shows the fastest propagation. The propagation duration is 0.5 second from 00:56:19.66 to 00:56:20.16 UT and the distance is 4.5 x 104 km. Thus, this rough estimation gives a speed of 9.0 xlO4 km s"1 that is 30 % of the light speed. An apparent velocity of 9 x 104 km s" 1 corresponds to the energy of « 23 keV. This seems to be much less than the energy (a 1.3 MeV) necessary for emitting the microwaves (17 GHz) that is estimated as follows. By using the potential magnetic field obtained by extending the magnetogram data on the photosphere by SOHO/MDI and using the numerical code Magpack2 (Sakurai 1982), we obtain the distribution of the magnetic field strength along the loop, which is approximately 200 Gauss in the middle area of the loop. From the power-law index a of the flux density, we derive the electron power-law index 5 = 4 based on Dulk (1985)'s approximation formula. Finally, according to the estimation by Bastian (1999), the energy of the electrons which mainly contribute to the gyrosynchrotron emission at 17 GHz is 1.3 MeV when the background magnetic field strength is 200 Gauss and the power-law index of electrons is S = 4. Thus, the actual velocity of the electrons is nearly the light speed. The observed apparent velocity 9 x 104 km s"1 (« 23 keV) is much less than the light speed. If we suppose that the bulk of energetic electrons was injected into the magnetic loop at some large angle to the field lines, then one can explain the low apparent velocity by the fact that the electrons were rotating around the magnetic field lines and effective path was much longer than the apparent length of the loops. In this case, the pitch angle of the electrons is Qpitch ~ arccos (v/c) ss 70 degree where v is the apparent velocity v w 9 x 104 km s" 1 and c is the speed of light. This may suggest that there is a possibility that the high energy electrons have large (> 70 degree) pitch angles when they are injected. The acceleration site of the high energy electrons is located near the south-east end of the loop where the propagation of the non-thermal signal starts from. The primary energy release site, therefore, is around this area. We consider the cause of this event is the moving parasitic polarity located just east of the compact -154-
microwave source near the south-east end of the loop where a relatively strong magnetic spot is located. From August 27 to the next day, a negative (opposite to the spot) parasitic polarity emerges just in the east of the spot and then moves to the west toward the spot with velocity of 0.2 km s"1. This motion may cause the magnetic reconnection with the pre-existing large scale loops connecting the sunspot positive polarity and the negative polarity near the northwest end of the elongated source. This situation is similar to the configuration suggested by Hanaoka (1999) and Nishio et al. (1997). SUMMARY AND DISCUSSIONS This paper is a review of the recent studies by the author and his colleagues in the solar flare physics based on the magnetic reconnection model. The observation of the reconnection inflows is strong evidence for supporting the magnetic reconnection model. However, the observation is still a unique example, hence an extensive search for other examples is necessary to confirm the universality of the model. Moreover, the MHD slow-mode shocks are another important target to find since they are the key for the Petschek-type fast reconnection model. The quantitative measurement of the physical variables, especially of the velocity, is essential to confirm the flows. These are important targets of the future missions, e.g. Solar-B. The MHD simulations including the effects of the heat conduction and the chromospheric evaporation are shown and a scaling law between the temperature and the emission measure of flares are derived based on their results. Since the simulation model is two-dimensional, extension to three-dimensional simulations is remained for the future work (see e.g. Birn et al. 2000). One of the expected three-dimensional effects would be the reduction or induction of the reconnection rate due to the formation of small-scale magnetic ropes in the neutral diffusion region. The onset of a flare as well as energy build-up of a flare are another issue to be solved that is not answered yet (see e.g. Choe and Cheng 2002). Theoretical and observational studies by using the vector magnetograms are becoming more and more important (Kusano et al. 2002). The quality of the data is expected to be dramatically improved when Solar-B is launched. The theory of the reconnection is, on the other hand, not yet completed. A self-consistent solution for the fast magnetic reconnection is not yet obtained that matches from the global to the microscopic scale under extremely large magnetic Reynolds number like the solar corona. The idea of the turbulent reconnection suggested by several authors ( e.g. Tajima & Shibata 1997) may be promising although the realization of this idea is still under investigation. The mechanism to produce the high-energy particles shown in this paper is still under discussion. The combination of hard X-rays observation by RHESSI and microwave observation by e.g. NoRH may give some indication.
REFERENCES Bastian, T. S., A. O. Benz, & D. E. Gary, Radio Emission from Solar Flares, Ann. Rev. Astron. Astrophys., 36, 131 (1998) Bastian, T. S., N. Nitta, A. L. Kiplinger, & G. A. Dulk. Energy Transport during a Solar Flare: VLA Observations of the Ml.9 Flare of 20 Aug 1992, in New Look at the Sun with Emphasis on Advanced Observations of Coronal Dynamics and Flares, eds. S. Enome, T. Hirayama 199, NRO Report 360,
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SECTION 3: Magneto spheric Substorms
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TYPES OF AURORAL-ZONE DISTURBANCES: CURRENT STATUS ON IDENTIFICATION AND UNDERSTANDING L. R. Lyons Department of Atmospheric Sciences, UCLA, 405 Hilgard Ave., Los Angeles, CA 90095-1565, USA
ABSTRACT It is now known that there are several different types of auroral zone disturbances, each reflecting fundamentally different processes within the magnetosphere-ionosphere system. Four types known to be dynamically important are described here. Substorms follow a > 0.5 hr growth period of enhanced convection. Onset occurs within the near-Earth plasma sheet and is often followed by new magnetic X-line in the mid-tail. Evidence now indicates that substorm onset results from a transition from a stable to unstable state of the nightside magnetosphere that initiates a few minutes prior to onset as the result of a reduction in the strength of convection caused by an appropriate change in the interplanetary magnetic field impacting the magnetosphere. Convection driven auroral enhancements occur at midnight-to-dawn MLTs and result from plasma sheet electrons, which are strongly energized during periods of enhanced convection and magnetic drift around the dawn side. Enhancements in solar wind dynamic pressure cause global disturbances that can be quite dramatic. These disturbances consists of increases in magnetospheric and ionospheric currents and electric fields, enhancements in auroral emissions and broadening of the auroral oval, and significant shrinkage of the polar cap size. Poleward boundary intensifications (PBIs) are auroral enhancements that initiate along the poleward boundary of the auroral oval. They occur during all overall levels of geomagnetic activity and are associated with few minute flow bursts in the tail plasma sheet. They are the most common disturbance. They are often repetitive and may be a manifestation of a large-scale ULF oscillation. INTRODUCTION Transient enhancements of auroral emissions and ionospheric currents often occur within the nightside auroral oval. Until the past few years, the substorm was the only disturbance of the magnetosphere-ionosphere system that was recognized as being important in causing these enhancements. However, recent studies have shown that the enhancements are associated with several different types of disturbances, each being an important and distinct type of disturbance and each reflecting fundamentally different dynamical processes within the magnetosphereionosphere system. Here I discuss four types of disturbances that are now known to be important: substorms, convection driven disturbances, dynamic pressure disturbances, and poleward boundary intensifications (PBIs). Each is described in a way that allows it to be distinguished from the others. The current status of our understanding of the physics of each is summarized and outstanding questions critical to the understanding of each are identified. SUBSTORMS Substorms are by far the most studied and recognized disturbance of the nightside auroral zone. The concept of a global substorm was based on ground-based observations of aurora (Akasofu, 1964) and is illustrated in Figure 1. Substorms are preceded by a growth phase of enhanced convection during which time the energy content of the plasma sheet plasma and lobe magnetic field increases. During the growth phase there may be localized brightenings along the poleward boundary of the nightside auroral oval (PBIs), some of which can extend equatorward through much of the auroral oval. These are unrelated to the substorm expansion phase and are discussed later. The substorm expansion phase initiates with a sudden brightening of an auroral arc within the equatorward portion of the auroral zone near magnetic midnight. After the initial auroral brightening, the region of active aurora expands poleward and azimuthally at speeds on the order of 1 km/s. Expansion-phase active aurora is accompanied by an enhanced westward ionospheric current that causes a large negative change in the ground magnetic field in the vicinity of magnetic midnight (e.g., Akasofu and Meng, 1967; Nishida and Kokubun, 1971; Kisabeth and Rostoker, 1971). This enhanced current, known as the "westward electrojet", moves poleward with the region of active aurora during the expansion phase (Kisabeth and Rostoker, 1971).
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Typically there are one or more east-west oriented auroral arcs within the auroral oval during the substorm growth phase. Using modern optical observations of substorm onsets obtained by the Canadian CANOPUS program, Lyons et al. (2002a) have recently found that auroral break-up at expansion phase onset does not generally occur along a pre-existing growth phase arc, as had previously been thought. Instead it at least often occurs along a thin new arc that forms equatorward of all growth phase arcs a few minutes prior to expansion-phase onset. After becoming discernible ~2-8 min before onset, this breakup arc grows in intensity monotonically until onset. After onset, the arc's intensity grows explosively and its shape becomes distorted by the development of large swirls. Figure 2 gives an example of all-sky-imager (ASI) observations from Lyons et al. (2002a) that show the formation and evolution of the breakup arc for times surrounding a substorm onset at 0459 UT on January 18, 1996. Notice that the breakup arc became discernible at ~0454 UT and then increased in intensity. The arc began to develop swirls and increase in intensity much more rapidly at 0459 UT, the time identified as substorm onset.
Fig. 1. Illustration of substorm expansion phase aurora and relation to auroral oval, open-closed magnetic field line boundary, and PBIs.
Fig. 2. 5577 A emissions from the Gillam ASI for periods surrounding an onset at 0459 UT on January 18, 1996. North is to the top and east is to the right in each image. A grid of magnetic latitude and longitude at 2° intervals is overlaid on the first image, the thicker horizontal and vertical lines indicating 67° latitude and 330° longitude, respectively. Intensity scales have been adjusted to emphasize the relatively low auroral intensities prior to and at onset (based on Lyons et al., 2002a). These observations imply that the processes responsible for expansion-phase onset initiate within the inner plasma sheet a few minutes prior to the time normally identified as onset, and they are consistent with a classical instability that grows linearly prior to onset and then becomes nonlinear (e.g., Voronkov et al., 2000). This implies that a few minutes prior to onset the magnetosphere-ionosphere system undergoes a transition from a stable to an unstable configuration that leads to the substorm expansion phase. For many years, it had been argued by a number of researchers that substorm expansion phase onset was the result of an instability associated with the formation of a magnetic x-line and enhanced reconnection in the midtail (x S -20 RE)- Consistent with this idea, Geotail spacecraft observations have shown significant evidence for x-line formation and enhanced reconnection at x —25 Rg for many substorms (Nagai et al., 1998; Machida et al, 1998, 1999). However, it was established about ten years ago that auroral breakup occurs on field lines that extend to very near the inner edge of the plasma sheet, which lies at x —6-10 R E (e.g., Samson et al., 1992). It was then speculated that reconnection at an X-line formed in the tail plasma sheet at x 5 -20 R E could eject plasma and magnetic flux earthward, leading to the initiation of expansion phase phenomena closer to the Earth (Birn and Hesse, 1996; McPherron and Fairfield, 1998). For this to occur, reconnection would have to initiate ~6-
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10 min prior to onset in order for flows to reach the inner plasma sheet and lead to formation and growth of the breakup arc. Geotail observations have now clearly shown that this does not occur (Lyons, 2000a and references therein). Additionally, flows perpendicular to the magnetic field earthward of x ~ -25 Rg have been found to be far too small to carry significant plasma and magnetic flux to the onset region (Lui et al., 1998). Thus the idea that an instability of the mid-tail leading to enhanced reconnection is the cause of substorm onset must be discarded, though the Geotail observations have indicated that enhanced reconnection at a mid-tail x-line can be an important aspect the dynamics of the substorm expansion phase. Whether or not such reconnection is important for all substorms has yet to be determined. It could, for example, be important for full substorms but not for pseudo breakups. Future research is necessary to answer such questions. We do not yet have a clear understanding of what happens within the plasma sheet that leads to the transition to instability a few minutes before expansion phase onset. We do, however, have significant information on the instability that will be important in the eventual understanding of this instability. First of all, we know that the instability does not initiate in the mid-tail, but instead initiates on field lines of the near-earth plasma sheet. This is clear from the Geotail observations. It is also clear from auroral observations, such as illustrated in Figure 2. These observations show that regions of the oval lying poleward of the initial region of bright expansion phase aurora are unaffected by the expansion phase until expansion phase auroral moves poleward to those regions (Akasofu, 1964; Lyons et al, 2002a). This tells us that expansion-phase dynamical changes within the plasma sheet that are significant enough to discernibly couple to the ionosphere move tailward from the inner plasma sheet with the expansion phase aurora. We furthermore know that the energy content of the plasma sheet plasma and lobe magnetic field increases during the substorm growth phase, a time when convection is enhanced. However, energy does not continue to increase during periods of prolonged enhanced convection, as would be expected if energy accumulation drives the magnetosphere into an unstable configuration that leads to the substorm expansion phase. Instead a balance is reached between tail energy input and loss, so that enhanced convection can persist for several hours without the occurrence of substorms. Periods of stable, enhanced magnetospheric convection are referred to as "convection bays" or "steady magnetospheric convection" periods (SMCs), and the energy content of the nightside magnetosphere appears to remain approximately constant during such periods. This implies that prolonged enhanced convection drives the magnetosphere into a stable equilibrium configuration with enhanced energy content on the nightside. It has also been established that substorm onsets are at least often "triggered" by interplanetary magnetic field (IMF) changes that are expected to lead to a reduction in the strength of convection (Lyons et al., 1997, and references therein; Hsu and McPherron, 2002). It has not been possible to use IMF measurements to determine whether or not all substorms are triggered in this way because of uncertainties in our knowledge of when, and if, IMF changes measured by a spacecraft outside the magnetosphere actually impact the magnetosphere. However, it is possible to directly look at the strength of convection using ionospheric radar measurements of plasma flow on the dayside. SuperDARN radar (Greenwald et al., 1995) measurements of dayside convection have shown that the vast majority of substorms with well-defined onsets are preceded by a reduction in the strength of convection (Lyons et al., 2002b). Limitations on radar echo coverage and measurement time resolution do not allow us to determine whether all substorm onsets, and in particular onsets separated by only a few minutes or less, are preceded by a reduction in the strength of convection. However, since at least most well defined onsets are preceded by such a reduction, it is plausible that most other onsets are preceded by such a reduction. Since the energy input to the plasma sheet particles increases with the strength of convection, it is reasonable to assume that the equilibrium energy content of the nightside magnetosphere increases with the strength of convection. This would imply that a reduction in the strength of convection would result in an equilibrium configuration for the magnetosphere with reduced energy content on the nightside. Thus if the energy accumulated on the nightside during a period of enhanced convection exceeds that of the equilibrium configuration after a reduction in the strength of convection, the excess energy would represent free energy that should be released after the reduction of convection [Atkinson, 1991]. It would then be the convection reduction and the resulting decrease in the equilibrium energy content of the magnetosphere that causes the transition to instability that leads to the substorm expansion phase. If it is the convection reduction that leads to the transition to instability, then the initiation of the convection reduction should occur at about the same time as the formation of the breakup arc. The observations of dayside convection of Lyons et al. (2002b) indicate that this is the case. An example of their observations is given in Figure 3, which shows polar plots of radar line-of-sight flows at times before, during, and after a substorm onset that occurred at 0913 UT on February 19, 1998. Measurements were made during successive two minutes intervals centered on the time indicted in each panel. Arrows give line-of sight flow speeds for each ionospheric location where a usable flow speed was obtained. These plots show that there was good data coverage between
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~75° and 80° latitude near local noon, which is an excellent region for measuring the strength of convection imparted to the magnetosphere from the solar wind. When looking at these plots, it must be remembered that only one component of the total plasma flow velocity is given by each arrow. It is expected, however, that a change in speed that is seen in a large number of vectors indicates an overall change in large-scale convection. A reduction in speed associated with onset can clearly be seen in Figure 3 at nearly all measurement locations, indicating a reduction in the strength of large-scale convection. The beginning of this reduction is first seen at a large number of measurement locations during the 0909 UT measurement interval. This indicates that the reduction initiated ~4-6 min prior to onset, which is the same time prior to onset as when the breakup arc typically initiates.
Fig. 3. Polar plots (MLT versus magnetic latitude) of line-of-sight velocities observed by SuperDARN radars for times before and after a substorm onset at 0913 UT on February 19, 1998. Vectors point toward or away from location of the radar that made the measurements and point away from small dots that give the location of each measurement. Locations of radars are indicated by heavy dots (Py for Pykkvibaer; St for Stokkseyri). Hankasalmi is at a lower latitude than included in the figure; its longitude is indicated by Ha (from Lyons et al., 2002b). That the reduction in convection precedes onset and occurs at about the same time relative to onset as does the formation of the breakup arc suggests that the reduction in the strength of convection causes the magnetosphere to make the transition from stability to instability. The instability that results from a reduction in convection would be expected to occur when the equilibrium energy corresponding to the lowered rate of convection is below that which has accumulated during the previous period of enhanced convection. This accumulated energy would be the equilibrium energy if the period of enhanced convection were sufficiently long to be an SMC. Shorter periods of enhanced convection, without sufficient time for equilibrium to be reached, would constitute the typical substorm growth phase and accumulated nightside energy would likely be below equilibrium. The instability will then reduce the energy stored on the nightside to the equilibrium level appropriate for the lowered strength of convection, and the release in energy would constitute the substorm expansion phase. If this is indeed the case, then the amount of energy released (i.e., substorm size) should increase with the amount of convection reduction for a given amount of energy storage prior to the convection reduction. It is also plausible that if the convection strength were to decrease, and then increase back to near or greater than the initial strength within a short period of time (less than ~10 min?), the instability would be terminated before leading to significant energy loss leading to a pseudo-breakup. If the above inferences are correct, then the critical outstanding question for understanding
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the substorm expansion phase is how does a convection reduction cause the transition to instability that leads to formation of the breakup arc prior to onset and to the substorm current wedge after onset. CONVECTION DRIVEN DISTURBANCES Large enhancements in auroral emissions have been recently identified at midnight to dawn MLTs during prolonged periods of strongly enhanced convection, which can include the growth phase of some substorms, SMCs, and magnetic storms, (Anderson et al, 2000a,b; Shue et al., 2002). Images from the POLAR spacecraft located well above the auroral oval are given in Figure 4 that show such a morningside auroral enhancement. This enhancement developed during a prolonged growth phase period of strongly southward IMF that preceded a large substorm with onset at 2227 UT on November 24, 1996. The morningside enhancement can be seen to have developed before the substorm onset and was distinct from the expansion phase enhancement that initiated near magnetic midnight. These morning auroral enhancements are believed to result from plasma sheet electrons that are strongly energized as they convect earthward during periods of strongly enhanced convection. As the electrons reach the inner region of the plasma sheet, magnetic drift carries them around the dawn side of the Earth, where strong pitch angle diffusion causes their precipitation and the resulting intense diffuse aurora.
Fig. 4. Images from the Visible Imaging System on the POLAR spacecraft (Frank et al., 1995). The convection driven auroral enhancement is at post-midnight to dawn MLTs, and the auroral enhancement associated with the substorm onset at 2227 UT is near midnight MLT (from Lyons et al., 2001).
Fig. 5. Left panel: X-ray data from the POLAR spacecraft for the October 19, 1998 storm. Grey scale gives the number of photons collected in 1 hr MLT bins during 1 min intervals with a 10 s time step. Acjjpc as measured directly along DMSP spacecraft trajectories is given by asterisks and by diamonds as corrected for the MLT of trajectory. Right panel: Horizontal component of ground magnetic field from auroral-zone ground stations located approximately every 3 hr in MLT. For each ground station, magnetic latitude and geographic longitude, respectively, are given (based on Anderson et al., 2000b). Because of the strong energization of these electrons, the disturbances show particularly clearly in POLAR observations of auroral X-rays, which are produced by precipitating electrons having energies >10 keV (Anderson et al., 2000, 2001). An example of the POLAR X-ray observations during a weak magnetic storm on October 19, 1998 is shown in the left panel of Figure 5. Here the number of X-rays counted, integrated over latitude, is shown as a function of MLT and UT. Overlaid on the left panel of Figure 5 is the cross-polar cap potential drop (A<j>pc) obtained from the low-altitude Defense Meteorological satellite Program (DMSP) satellites. The IMF z-
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component was quite steady at ~-15 nT during this period. The morningside enhancement can be seen clearly from ~24-7 MLT throughout most of the 8 hr period shown in the Figure, illustrating that these disturbances can persist for several hours. It can also be seen that the intensity of the morningside emissions correlates quite well with the magnitude of Acfipc, as expected if the precipitating electrons responsible for the emissions are energized by the cross-tail electric field. The distinction between morning side enhancements and substorm enhancements can be seen following a substorm onset that occurred at ~1325 UT and was followed by a large enhancement in X-ray emissions at -18-24 MLT. The right panel of Figure 5 shows the horizontal magnetic field from auroralzone ground stations located approximately every 3 hr in MLT. The magnetograms at SMI, EAG, and MCQ show a few hundred nT magnetic depression on the morning side throughout much of the period of the morningside auroral enhancement. Shorter effects of the 1325 UT substorm can also be seen on the night side. The morningside depression indicates that the morningside auroral enhancements are associated with significantly enhanced ionospheric currents and thus also with important enhancements in the currents coupling the magnetosphere and ionosphere. Not much attention has yet been paid to these morningside enhancements driven by enhanced convection. Thus much needs to be done to determine the simultaneous ionospheric and field-aligned currents and electric fields associated with the enhancements, their overall contribution to magnetospheric dynamics, and a more precise description of conditions for when they are important.
Fig. 6. A: IMF and solar wind data for 09-12 UT on January 10, 1997 from WIND located upstream and to the dawn side of the Earth. B: Ultraviolet images from the POLAR spacecraft located several RE above the auroral zone. Midnight is to the right, and dusk is to the bottom, of each image. Lighter shadings indicate more intense auroral emissions (based on Lyons, 2000b]). DYNAMIC PRESSURE DISTURBANCES Recent studies have shown that solar wind dynamic pressure enhancements have very important effects on auroral zone activity that are very different from those associated with other types of disturbances (Zesta et al., 2000a, Lyons, 2000b; Lyons et al., 2000a; Chua et al., 2001). These effects were seen clearly during a widely studied period of strongly southward IMF that occurred on January 10, 1997. Figure 6A shows the IMF y- and zcomponents and the solar wind radial speed from 9-12 UT on that day as observed by the WIND spacecraft located upstream from the Earth. A pulse of strongly enhanced density was observed between 1030 and 1055 UT. This corresponds to a strong enhancement of solar wind dynamic pressure (the solar wind speed did not change appreciably during the pulse) which, based on simultaneous measurements from Geotail within the magnetosheath, first impacted the magnetosphere at ~1048 UT. Ultraviolet images from the POLAR spacecraft showing the two-dimensional evolution of the aurora during the period of this sharp pressure pulse are displayed in Figure 6B. As can be seen from the images, the strong pressure pulse caused a rapid global enhancement of auroral emissions as well as a significant poleward motion of the poleward boundary of the aurora. These effects initiated at nearly the same time as the dynamic pressure enhancement first impacted the magnetosphere and were seen nearly simultaneously over most of the auroral oval (see 1048 UT image in Figure 6B). The auroral enhancement began to fade away when the dynamic pressure went back down (see 1112 and 1118 UT images). The total poleward displacement of the poleward boundary of the aurora was ~10° in latitude during a ~10 min
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period at some MLTs (Zesta et al., 2000a). This corresponds to an extremely rapid and large closing of open polar-cap magnetic field lines, a broadening of the auroral oval, and a reduction in area of the polar cap. These effects have been seen for other events, and Boudouridis et al. (2002) made an initial study of how the effects depend on the orientation of the IMF during the time of a pressure enhancement. They found that the effects are strongest when the IMF is southward before and after the enhancement, such as was the case for the January 10, 1997 example. For other IMF orientations, most of the same effects were observed, but they were not as strong. Boudouridis et al. (2002) used particle observations from the DMSP low-altitude spacecraft to directly investigate the changes in auroral precipitation associated with dynamic pressure increases. Figure 7 shows plots of integral and differential energy fluxes of precipitating electrons and ions observed by the DMSP F13 spacecraft before (left column) and during the peak (right column) of the pressure pulse on January 10, 1997. A small insert at the top right corner of each panel gives the orbit of the spacecraft, which for F13 is approximately dusk-to-dawn. The second and fourth panels for each orbit show, respectively, the total precipitating electron and ion energy fluxes, and the third and fifth panels show, respectively, the electron and ion precipitation energy fluxes versus particle energy. Vertical lines in the second panels indicate the equatorward and poleward boundaries of the region of auroral particle precipitation, the poleward boundary being at the boundary between the plasma sheet and polar rain, which is the boundary between open and closed magnetic field lines. Before the pressure pulse (left column) F13 observed the separatrix at 65° in the late afternoon and at 68° just past dawn. During the pressure pulse, F13 observed the separatrix at 67° near dusk and 79° near dawn. This corroborates the dramatic reduction of the size of the polar cap region of open flux and the large enhancement of the latitudinal extent of the plasma sheet precipitation region for this event that was deduced by Zesta et al. [2000a] from auroral images. The integral precipitating particle energy fluxes show that the enhancement of solar wind dynamic pressure also strongly increased the energy flux of electrons precipitating into the atmosphere. On the dusk side the increase was associated with an increase of the precipitating flux per unit area, and on both the dawn and dusk sides it was associated with a large increase in the latitudinal width of the region of high precipitating particle energy fluxes. It should be noted that the dawn side particle measurements in the right panel of Figure 7 were made at 11221127 UT, which is several minutes after the solar wind pressure reduced and after significant auroral fading had occurred (see Figure 7). Thus the precipitating energy fluxes per unit area of the dawn side were almost certainly significantly higher during the peak of the dynamic pressure pulse than observed by F13.
Fig. 7. DMSP F13 particle observations and horizontal flow velocities in the direction normal to the satellite trajectory (top panel, positive sunward) from passes before and during the peak of the January, 10, 1997 pressure pulse. The integral (eV/cm2-sr-sec) and differential (eV/cm2-sr-sec-eV) precipitating ion and electron energy fluxes are plotted for each pass, and the latitudinal extent of the regions of auroral particle precipitation is identified by vertical lines and shading for each pass (based on Boudouridis et al., 2002). It is well established that solar wind dynamic pressure strongly affects the strength of the magnetopause and crosstail currents. Recently it has been shown that solar wind dynamic pressure also exerts strong control on the strength of global ionospheric currents (Shue and Kamide, 1998; Lyons et al., 2000a), and in particular the pressure pulse on January 10, 1997 greatly increased the strength of these currents (Zesta et al., 2000a). This
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global current enhancement suggests that the region 1 field-aligned current system, which feeds the ionospheric currents, is also enhanced when the solar wind dynamic pressure increases (Zesta et al., 2000a). An increase in ionospheric currents must be due to an increase in ionospheric conductivities and/or the mapping of magnetospheric electric fields to the ionosphere. The enhancement in auroral intensities indicates that an increase in ionospheric conductivities is at least partially responsible for the current increase within the auroral oval. However, the analysis of Zesta et al (2000a) shows current enhancements at very high latitudes during the January 10, 1997 pressure enhancement, and a detailed analysis of magnetic records from near the magnetic pole (Lukianova and Troshichev, 2002) shows that currents well poleward of the auroral oval have as strong a response to solar wind dynamic pressure variations as they do to the z-component of the IMF. Because these studies were of currents well poleward of the aurora, the currents could not have been affected by conductivity changes associated with the aurora. Thus the results of the studies imply that that the strength of the polar cap electric field depends strongly on the solar wind dynamic pressure. Shown in the top panel of Figure 7 are preliminary new results from Boudouridis et al. (personal communication, 2002) on the speed of convection. This panel shows flow speeds measured in the direction normal to the F13 trajectory. The flows show the traditional two-cell convection pattern, which is sunward within the auroral oval and anti-sunward over the polar cap. Comparing the magnitude of the flow before and during the pressure enhancement shows that the increase in dynamic pressure led to a large (factor of ~2-3 at most locations) increase in convection speed throughout both the auroral zones and the polar cap. We thus see that solar wind dynamic pressure enhancements have large effects on the strength of global convection, as well as on the intensity of auroral precipitation, the size of the auroral oval, the strength of magnetospheric and ionospheric currents, and the size of the polar cap. These effects appear to be as large as the effects of the IMF. However, the importance of these effects has only recently been realized, and much more study is needed to assess their full significance and their variations with the IMF. POLEWARD BOUNDARY INTENSIFICATIONS PBIs occur frequently in the nightside auroral zone and are generally the most intense auroral disturbance at times other than during the expansion phase of substorms or following a strong enhancement in solar wind dynamic pressure. They occur repetitively so that many individual disturbances can occur during time intervals of ~1 hr. They also occur during all levels of overall geomagnetic activity, including very quiet periods, substorm periods, SMCs, and magnetic storms. As illustrated in Figure 1, PBIs have an auroral signature that initiates along the poleward boundary of the auroral oval and often extends equatorward. They are longitudinally localized, and there can simultaneously be several separate intensifications along the auroral oval. Typical ground magnetic perturbations associated with PBIs are a few tens of nT, but perturbations can be as high as ~500 nT during periods of enhanced convection. Individual PBIs are longitudinally localized. They are associated with longitudinally localized regions of enhanced equatorward plasma flow in the ionosphere (de la Beaujardiere et al., 1994) and longitudinally localized bursts of earthward flow in the tail (Lyons et al., 1999; Zesta et al., 2000b). PBIs can extend a considerable distance equatorward of the poleward boundary of the auroral zone leading to auroral structures that are elongated in the north-south direction (e.g., Henderson et al., 1998). It is not unusual for individual events to traverse essentially the entire latitudinal extent of the plasma sheet, which would correspond to flow bursts that extend from the distant tail plasma sheet (~50-100 R E ) all the way to the vicinity of synchronous orbit. The plasma flow bursts in the tail associated with PBIs are known to be an important dynamical feature of the magnetosphere-ionosphere system. They are associated with electric fields that, when mapped to the ionosphere, give the enhanced ionospheric flows reported by de la Beaujardiere et al. (1994). The longitudinal localization of these electric fields results in regions of converging ionospheric Pedersen currents that close with upward fieldaligned currents that are often > 1 x 10"6 A/m2 at the ionosphere. This is sufficiently large to require that electrons be accelerated by field-aligned electric fields, leading to the PBI auroral signature. Consistent with this, magnetic perturbations in the tail associated with flow bursts imply structured currents in the tail having densities > 1 x 10 -1 ° A/m2. These currents would map along field lines to current densities of > 1 x 10~6 A/m2 in the auroral ionosphere, as expected if the flow bursts connect along field lines to the electric fields that lead to the PBIs (Lyons et al., 2000b). Figure 8 shows nightside 5577 A emission intensities versus UT and magnetic latitude from a merging of meridian scanning photometer (MSP) data from Rankin Inlet and Gillam, located along the same magnetic meridian at latitudes of 73.3° and 67.1°, respectively. The MSPs measure the intensity of auroral emissions at various wavelengths and complete one full meridional scan every minute. By assuming that the auroral emissions at one wavelength come from the same altitude in the ionosphere, the intensities measured versus zenith angle have been converted to intensity versus geomagnetic latitude. Intensifications of 5577 A emissions, which result from harder (>1 keV) aurora electron precipitation, indicate the existence of discrete auroral forms. The white line
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in Figure 8 is at the poleward boundary of the 6300 A emissions, which has been shown to be within ~1° in latitude of the magnetic separatrix boundary (Blanchard et al., 1997). The time period shown is 0200-0800 UT on January 7, 1997, a time period analyzed by Zesta et al. (2000b). Figure 8 shows a substorm onset at this longitude at 0305 UT (~2030 MLT), preceded by a growth phase with equatorward motion of the equatorward edge of the aurora. The substorm onset is identified by intensification of the aurora near its equatorial edge and subsequent poleward expansion of the region of intensified aurora. A second intensification occurred at ~0420 UT. Then, from 0440 UT onwards, there is a continuous series of repetitive PBIs as indicated by intensifications near the poleward boundary of the auroral zone. Some of these intensifications extend equatorward, as occurred for a series of PBIs that initiated after 0500 UT and are labeled 1 5 in the figure. These data show that PBIs can be quite intense and are significantly different from substorms.
Fig. 8. 5577 A emission intensities versus UT and magnetic latitude on January 7, 1997 from a merging of MSP data from Rankin Inlet and Gillam. Thin vertical arrows identify five PBIs extending equatorward into the indicated field of view of the Gillam all-sky imager at 5577 A. A thicker black arrow identifies a substorm onset (based on Zesta et al., 2000b).
Fig. 9. Series of 5577 A images from Gillam all-sky imager. White vertical and horizontal grid lines indicate 330° longitude and 67° latitude, geomagnetic, respectively. Black longitude grid lines are 2° apart and black latitude grid lines are 1° apart (based on Zesta et al., 2000b). The equatorward extending PBIs labeled 1-5 in Figure 8 extended into the 5577 A field-of-view of an all-sky auroral imager located at Gillam. Figure 9 shows a series of 5577 A images from this imager. North is at the top of each image and east is to the right. Two different time periods are included in the figure: 0304-0310 UT and 0512-0525 UT. The image at the beginning of each time series includes for reference a grid of geomagnetic coordinates. The top time series shows a clear westward traveling surge that occurred during a substorm expansion phase. The surge appear as a large auroral swirl that entered the imager field-of-view at ~0304 UT. The bottom time period shows two north-south aligned aurora forms that appeared in the northeast edge of the Gillam field of view and then were observed moving westward. The first one is the weakest, appears first in the 0512 UT image, had an east-west tilt, and is most clearly seen in the 0515 UT image. The second north-south structure is first visible in the 0514 UT image. It became very intense by 0518 UT and can be seen to have
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quickly moved westward until 0524 UT. These two north-south aligned forms correspond very well to PBIs 1 and 2 of Figure 4, and they are very different from the substorm onset aurora shown in the 0304-0310 UT time series. Images subsequent to those shown in Figure 9 show that the other three PBIs identified in Figure 8 were also approximately north-south structures. Other examples show that equatorward extending PBIs can sometimes be equatorward moving east-west structures and that PBIs that stay near the polar cap boundary often take the form of beads or swirls along that boundary (Zesta et al., 2002). It is not currently know what causes PBIs and the tail flow bursts. Both can be repetitive and take the appearance of oscillations with periods of ~15 min or ~30 min. It has been suggested that PBIs and tail flow bursts may both be components of large-scale ULF waves on the night side (Sanchez et al., 1997; Lyons et al., 2002). If this is true, then global ULF modes would be critical to both tail dynamics and to magnetosphere-ionosphere coupling along auroral field lines, and they would often play an important role in auroral-zone activity. CONCLUDING STATEMENT In this review and important and how each distinguishing research.
I have considered four different types of auroral-zone disturbances, each of which reflects a distinct type of disturbance of the magnetosphere-ionosphere system. What is currently known about each, can be uniquely identified, has been summarized. I hope that this information will be of value in the different types of disturbances from each other and in identifying important areas for future
AKNOWLEDGEMENTS As indicated by the co-author lists below, much of the work described above was done in collaboration with a large number of other researchers. I am extremely grateful for this talented and generous group of collaborators. Writing of this paper has been supported in part by NSF grant OPP-0136139 and NASA grants NAG5-7962 and NAG5-11858 REFERENCES Akasofu, S.-I., The development of the auroral substorm, Planet. Space Sci.. 12. 273, 1964. Akasofu, S.-I., and C.-I. Meng, Polar magnetic substorm in the evening sector, J. Atmos. Terr. Phys.. 29. 1127, 1967. Anderson, P. C , D. L. McKenzie, L. R. Lyons, and M. Hairston, Global X-ray observations of magnetospheric convection-driven auroral disturbances, Geophys. Res. Lett., 27, 3233, 2000a. Anderson, P. C , D. L. McKenzie, M. L. Brittnacher, M. W, Chen, M. Hairston, and M. F. Thomsen, Global storm time auroral X-ray morphology and timing and comparison with UV measurements, J. Geophys. Res., 15,757, 105, 2000b. Atkinson, G., A magnetosphere wags the tail model of substorms, in Magnetospheric Substorms, eds. J. R. Kan, T. A. Potemra, S. Kokubun, and T. Iijima, p. 191, American Geophysical Union, Washington, 1991. Birn, J., and M. Hesse, Details of current disruption and diversion in simulations of magnetotail dynamics, J. Geophys. Res., 101, 15,345, 1996 Blanchard, G. T., L. R. Lyons, and J. C. Samson, Accuracy of 6300 A auroral emission to identify the separatrix on the night side of the Earth, J. Geophys. Res., 102, 9697, 1997. Boudouridis A., E. Zesta, L. R. Lyons, P. C. Anderson, and D. Lummerzheim, The effect of solar wind pressure pulses on the size and strength of the auroral oval, J. Geophys. Res., 2002 (in press). Chua, D., G. Parks, M. Brittnacher, W. Peria, G. Germany, J. Spann, and C. Carlson, Energy characteristics of auroral electron precipitation: A comparison of substorms and pressure pulse related auroral activity, J. Geophys. Res., 106, 5945, 2001. de la Beaujardiere, O., L. R. Lyons, J. M. Ruohonmiemi, E. Friis-Christensen, C. Danielsen, F. J. Rich, and P. T. Newell, Quiet-time intensifications along the poleward auroral boundary near midnight, J. Geophys. Res., 99, 287, 1994. Frank, L. A., J. B. Sigwarth, J. D. Craven, J. P. Cravens, J. S. Dolan, M. R. Dvorsky, P. K. Hardebeck, J. D. Harvey and D. Muller, The Visible Imaging System (VIS) for the Polar spacecraft, Space Sci. Rev., 71, 297, 1995. Greenwald, R. A., et al., DARN/SuperDARN: A global view of the dynamics of high latitude convection, Space Sci. Rev., 71, 761, 1995. Henderson, M. G., G. D. Reeves, and J. S. Murphree, Are north-south structures an ionospheric manifestation of bursty bulk flows?, Geophys. Res. Lett., 25, 3737, 1998. Hsu, T.-S., and R. L McPherron, An evaluation of the statistical significance of the association between northward turnings of the IMF and substorm expansion onsets, J. Geophys. Res., 2002 (in press). Kisabeth, J. L., and G. Rostoker, Development of the polar electrojet during polar magnetic substorms, J. Geophys. Res., 76, 6815, 1971.
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Lui, A. T. Y., K. Liou, P. T. Newell, C.-I. Meng, S.-I. Ohtani, T. Ogino, S. Kokubun, M. J. Brittnacher, and G. K. Parks, Plasma and magnetic flux transport associated with auroral breakups, Geophys. Res. Lett., 25, 4059, 1998. Lukianova, R. and O. A. Troshichev, Magnetospheric response to the solar wind dynamic pressure inferred from the polar cap index, in Sixth International Conference on Substorms (ICS-6) Eur. Space Agency, Noordwijk, 2002 (in press). Lyons, L. R., Determinations of Relative Timing of Near-Earth Substorm Onset and Tail Reconnection, in Fifth International Conference on Substorms (ICS-5), p. 255, Eur. Space Agency, Noordwijk, 2000a. Lyons, L. R., Geomagnetic disturbances: characteristics of, distinction between types, and relations to interplanetary conditions, J. Atmos. and Solar-Terr. Phys., 62, 1087, 2000b. Lyons, L. R., G. T. Blanchard, J. C. Samson, R. P. Lepping, T. Yamamoto, and T. Moretto, Coordinated observations demonstrating external substorm triggering, J. Geophys. Res., 102, 27,039, 1997. Lyons, L. R., T. Nagai, G. T. Blanchard, J. C. Samson, T. Yamamoto, T. Mukai, A. Nishida, and S. Kokubun, Association Between GEOTAIL Plasma Flows and Auroral Poleward Boundary Intensifications Observed by CANOPUS photometers, J. Geophys. Res., 104, 4485, 1999. Lyons, L .R., E. Zesta, J. C. Samson, and G. D. Reeves, Auroral disturbances during the January 10, 1997 magnetic storm, Geophys. Res. Lett., 27, 3237, 2000a. Lyons, L. R., T. Nagai, J. C. Samson, E. Zesta, T. Yamamoto, T, Mukai, A. Nishida, and S. Kokubun, Structured currents associated with tail bursty flows during turbulent plasma sheet conditions, in Magnetospheric Current Systems, ed. by S. Ohtani and R, L. Lysak, American Geophysical Union, Washington, p. 267, 2000b. Lyons, L .R., R. L. McPherron, E. Zesta, G. D. Reeves, and J. Sigwarth, Timing of substorm signatures during the November 24, 1996 Geospace Environment Modeling Event, /. Geophys. Res., 106, 349, 2001. Lyons, L. R., I. O. Voronkov, E. Donovan, and E. Zesta, Relation of substorm breakup arc to other growth-phase auroral arcs, J. Geophys. Res., 2002a (in press). Lyons, L. R., J. M. Ruohoniemi, S. Liu, S. I. Solovyev, J. C. Samson, Observations of dayside convection reduction leading to substorm onset, J. Geophys. Res., 2002b (submitted). Lyons, L. R., E Zesta, Y. Xu, E. R Sanchez, J. C. Samson, G. D. Reeves, J. M. Ruohoniemi, and J. B. Sigwarth, Auroral Poleward Boundary Intensifications and Tail Bursty Flows: A Manifestation of a Large-Scale ULF Oscillation?, J. Geophys. Res., 2002c (in press). Machida, S., Y. Miyashita, A Ieda, A. Nishida, T. Mukai, Y. Saito, T. Yamamoto, and S. Kokubun, Time evolution of the Earth's magnetotail associated with substorm onset: GEOTAIL observations, in Substorms 4, edited by S. Kokubun and Y. Kamide, Kluwer Academic Publishers, Boston, p. 149, 1998. Machida, S., Y. Miyashita, A Ieda, A. Nishida, T. Mukai, Y. Saito, and S. Kokubun, GEOTAIL observations of flow velocity and north-south magnetic field variations in the near and mid-distant tail associated with substorm onsets, Geophys. Res. Lett., 26, 635, 1999. McPherron, R. L., and D. H. Fairfield, Session 1 Summary: What are the major expansion phase actions seen at various regions, in Substorms-4: International Conference on Substorms-4, Lake Hamana, March 9-13, 1998, edited by S. Kokubun and Y. Kamide, p. 29, Kluwer Acad., Norwell, Mass., 1998. Nagai, T., M. Fujimoto, Y. Saito, S. Machida, T. Terasawa, R. Nakamura, T. Yamamoto. T. Mukai, A. Nishida, and S. Kokubun, Structure and dynamics of magnetic reconnection for substorm onsets with GEOTAIL observations, J. Geophys. Res., 103, 4419, 1998. Nishida, A., and S. Kokubun, New polar magnetic disturbances: SqP, SP, DPC, and DP2, Rev. Geophys., 9, 417, 1971. Samson, J. C , L. R. Lyons, B. Xu, F Creutzberg, and P. Newell, Proton aurora and substorm intensifications, Geophys. Res. Lett., 19, 2167, 1992. Sanchez, E. R., J. D. Kelly, V. Angelopoulos, T. Hughes, and H. Singer, Alfven modulation of the substorm magnetotail transport, Geophys. Res. Lett., 24, 979, 1997 Shue, J.-H., and Y. Kamide, Effects of solar wind density on the westward electrojet, in Substorms-4 edited by S. Kokubun and Y. Kamide, Kluwer Academic, Boston, p. 677, 1998. Shue, J.-H., P. T. Newell, K. Liou, C.-I. Meng, Y. Kamide, and R. P. Lepping, Two-component auroras, Geophys. Res. Lett, 29(10), 10.1020/2002GL014657, 2002. Voronkov, I., E. F. Donovan, B. J. Jackel, and J. C. Samson, Large-scale vortex dynamics in the evening and midnight auroral zone: Observations and simulations J. Geophys. Res., 105, 18,505, 2000. Zesta, E, H. J. Singer, D. Lummerzheim, C. T. Russell, L. R. Lyons, M. J. Brittnacher,., The effect of the January 10, 1997 pressure pulse on the magnetosphere-ionosphere current system in Magnetospheric Current Systems, ed. by S. Ohtani and R, L. Lysak, American Geophysical Union, Washington, p217, 2000a Zesta, E., L. Lyons, and E. Donovan, The auroral signature of earthward flow bursts observed in the magnetotail, Geophys. Res. Lett., 27, 3241, 2000b. Zesta, E., E. Donovan, L. Lyons, G. Enno, J. S. Murphree and L. Cogger The Two-Dimensional Structure of Auroral Poleward Boundary Intensifications (PBI), J. Geophys. Res., 2002 (in press)
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GEOTAIL-POLAR OBSERVATION OF SUB STORM-TIME FIELD INCREASE IN THE TAIL AND THE POLAR MAGNETOSPHERE H. Kawano1, C. T. Russell2, G. Rostoker3, G. Le 4 , G. K. Parks5, Y. Saito6, and T. Mukai6 1
Department of Earth and Planetary Sciences, Kyushu University, Fukuoka City, Fukuoka 812-8581, Japan Institute of Geophysics and Planetary Physics, University of California, Los Angeles, CA 90095-1567, U.S.A. 3 Department ofPhysics, University of Alberta, Edmonton, AB T6G2J1, Canada A Laboratory for Extraterrestrial Physics, NASA/Goddard Space Flight Center, Greenbelt, MD 20771, U.S.A. s Space Sciences Laboratory, University of California, Berkeley, CA 94720, U.S.A. 6 'Institute of Space and Astronautical Science, Sagamihara, Kanagawa 229, Japan
2
ABSTRACT We present a case study of substorm-time magnetic field perturbations in high-altitude polar lobe and in the tail lobe, simultaneously observed by the POLAR and GEOTAIL satellites, together with ground-based CANOPUS observations and WIND solar-wind observations. During the growth phase, the magnetic field strength (5totai) increased both in the polar lobe (observed by POLAR) and in the tail lobe (observed by GEOTAIL); this is ascribed to the pileup of dayside-reconnected field lines over the magnetopause. On the other hand, while Btoa\ at GEOTAIL decreased during the expansion phase, 5totai at POLAR did not start decreasing until ~35 min after the expansion onset. This absence of the field-decrease signature in the polar lobe for ~35 min could be attributed to dipolarization/compression of closed field line region of the nightside inner magnetosphere. This compression could offset the effect of the magnetic field decrease in the polar lobe, and would be absent in the open field line region of the tail lobe allowing GEOTAIL to detect the decrease in -Btotai during this ~35 min period.
INTRODUCTION POLAR, launched in 1996, had its apogee in the high-altitude polar magnetosphere, and thus stayed there for several hours during each rotation around the Earth. This unique long dwell time enables a detailed in-situ study of perturbations in the high-altitude polar magnetosphere during substorms. Kawano et al. (2002) (referred to Paper 1 below) took advantage of such an opportunity and presented a case study of substorm-time magnetic field perturbations in the high-altitude polar magnetosphere, corresponding to the polar lobe, using data from the POLAR spacecraft together with ground-based CANOPUS observations and WIND solar-wind observations. Reported features in the paper include the following: During the growth phase, the magnetic field strength (5totai) increased at POLAR; and after the expansion onset, Stotai continued to increase until ~28 min after the expansion onset, then decreased. The increase in Btotai during the growth phase is consistent with past observations in the tail lobe, where Btotai also tends to increase during the growth phase (e.g., Caan et al., 1978); they are both ascribed to the pileup of magnetic field lines over the polar to tail magnetosphere. On the other hand after the expansion onset, the past observations show that 5 ^ 1 m t n e t a 'l l°be tends to decrease; thus, the increase in 5t0tai in the polar magnetosphere during the expansion phase, as reported in Paper 1, is opposite to what is usually observed in the tail lobe. It is generally agreed that the total magnetic flux content of the tail lobes decrease when reconnection at the near-Earth neutral line (NENL) reaches the open field lines of the lobes. This results in a decrease in the field strength (rarefaction) in the lobes as the shape of the flaring tail changes. This rarefaction effect propagates, via the fast mode wave, to the lobe parts away from NENL, including the polar lobe; as a result, Btotai m the polar lobe should also start
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decreasing. Because the fast mode wave speed is very high (> 5000 km/s) in the lobe, 5t0tai in the polar lobe should start decreasing soon (in less than a minute) after the expansion onset. The expansion-phase increase in Btotai reported in Paper 1 is opposite to this expectation, and thus surprising. To explain this discrepancy, Paper 1 suggested that the inner magnetosphere, including the polar magnetosphere, was compressed during the expansion phase due to the dipolarization (see their Figure 6 for illustration); because this compression effect does not operate in the tail lobe, Btotai in the tail lobe should decrease due to the reconnection at NENL, consistent with observations there. In this way, this explanation appears to fit with both the POLAR Btotai observation of Paper 1 and the typical 5totai feature in the tail lobe. However, since Paper 1 did not have simultaneous observations in the tail, one cannot deny that the tail lobe field somehow increased during the expansion phase of this event, too. The purpose of this paper is, then, to find substorm cases for which both the polar region and the tail region was simultaneously observed by two satellites, and to study B tota | perturbations in the two regions at the same time. We have found such a case, as presented below, in which the polar magnetosphere was observed by the POLAR satellite while the tail was observed by the GEOTAIL satellite. OBSERVATIONS The substorm event studied in this paper took place on January 28, 1997; for this event, Figure 1 shows data from the POLAR (Russell et al., 1995) and GEOTAIL (Kokubun et al., 1994; Mukai et a l , 1994) satellites, along with data from the geosynchronous satellite GOES 8, the solar wind data from the WIND spacecraft (Lepping et al., 1995; Ogilvie et al., 1995), and the key parameter CL as a monitor of the substorm activity; CL is derived using ground magnetometer data from CANOPUS (Rostoker et al., 1995) and indicates the strength of the westward electrojet flowing over the stations of the array. Locations of POLAR and GEOTAIL satellites in GSM coordinates are shown at the bottom of the figure. As shown there, POLAR was located in the polar magnetosphere and GEOTAIL was located in the tail. Location of GOES 8 satellite is also shown at the bottom of the figure in terms of the GSM local time ("LTGSM"), which is defined as (arctan (y/x) /n x 12 + 12), in which x and y refer to the GOES 8 position in GSM coordinate system. We first examine the solar wind data. Their propagation lag is already roughly corrected in Figure 1 by adding —XjVx to the observed time, where Vx and X are the GSM X components of the solar wind ion bulk velocity observed by WIND and the WIND position, respectively. The position of WIND in GSM coordinates was (172.4,-10.5,-23.3) RE at 00:00 UT and (173.1,-15.4,-20.0) RE at 06:00 UT. Panel J of Figure 1 shows that the solar wind dynamic pressure was fairly constant throughout the interval of the figure. On the other hand, Panel I shows that the IMF Bz was negative from ~02:35 UT to ~04:05 UT. Thus, it is expected that the growth phase of this substorm started around ~02:35 UT (vertical line I). Panel E shows the CL parameter. It started to enhance around ~03:10 UT. From this, we estimate the expansion onset time of this substorm to have been around M)3:10 UT (vertical line II). We note that the CANOPUS stations were off the nominal onset position: At the north-south station line through Fort Churchill, located at ~265° geographic longitude and at the eastern end of CANOPUS stations, the local time at ~03:10 UT was ~21 hr. However, simultaneous POLAR auroral images (not shown here; readers are referred to the CDAWeb, http://cdaweb.gsfc.nasa.gov, for the images) show that the onset around ~03:10 UT actually took place near the Fort Churchill line. Thus we think this onset time is more or less correct. Panel F, G, and H show the magnetic field observed by a geosynchronous satellite GOES 8. Panel F shows the magnetic field strength; Panel G shows the field latitude [deg] defined as (arctan (BH/VBV2+BD2) /n x 180J, in which BH, BV, and BD are the magnetic field components in the VDH coordinate system; Panel H shows the field longitude [deg] defined as (arctan (BD/BV) /n x 180). Around the time of line II, the magnetic field strength (Panel F) and the field latitude (Panel G) started to increase; this shows the start of dipolarization, consistent with the onset identification based on the CL parameter. (There could be up to ~ 10 minutes of difference in the CL-based onset time and GOES-based onset time, but it does not much affect the storyline of this paper, because the main point of this paper is the time difference between line II and line III, as we will discuss later.) Panel D shows the pressure inferred from GEOTAIL in the tail. The dotted curve shows the magnetic pressure, and the solid curve shows the total pressure, that is, the magnetic plus thermal pressure. One can see that the two curves are close to each other near the times of the vertical lines; this means that the spacecraft was mainly located in the tail lobe during this substorm. -173-
Fig. 1. Shows, from top, the magnetic field observed by POLAR (Panels A, B and C), the pressure inferred from GEOTAIL data (D), the key parameter CL from the CANOPUS network (E), the magnetic field observed by GOES 8 (Panels F, G and H), and the IMF Bz (in GSM coordinates) (I) and the solar wind dynamic pressure (J) observed by WIND. Panel B shows the magnetic field strength; the solid line shows the observed values, and the dotted line shows the Tsyganenko (1996) model values. Panel C shows Sfitotai. the difference of the two lines in Panel B. Panel A shows 8^, perturbation in the "field flaring angle": Refer to text for its definition. In Panel D, the dotted line shows the magnetic pressure, and the solid lines shows the total pressure, that is, the sum of the magnetic and plasma pressures. Panel F shows the magnetic field strength. Panel G and H show the field latitude and longitude [deg]; refer to text for their definitions. Vertical lines I, II, and III correspond to the start time of the growth phase, expansion phase, and the decrease in 55totai at POLAR; refer to the text for details.
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The time profile of the total pressure at GEOTAIL, in the tail lobe, is such that it increased during the substorm growth phase (between the vertical lines I and II), and decreased during the substorm expansion phase (after the time of the vertical line II). This is consistent with the statistical pattern obtained in the past studies (e.g., Caan et al., 1978). We finally look at the POLAR data in the polar magnetosphere: In Panel B of Figure 1, the solid curve shows the observed 5 t 0 t a |. The curve shows a smoothly decreasing long-term trend, and superposed to it a perturbation starting near the vertical line I, as shown by the bent there. The smoothly decreasing long-term trend in 5,0tai is not a time-dependent change in the field but an apparent effect of the satellite motion; we want to remove it and extract the true time-dependent perturbation in5 t0t ai. For this purpose, we use the Tsyganenko (1996) model as a model expressing the position dependence of the field, and subtract its value at the satellite position from the observed data. For this purpose the model itself should not be time dependent, so we fix the model parameters (Dst=0 nT, Pdyn=2 nPa, IMF By-0 nT, and IMF Bz=O nT). The corresponding model field strength at the satellite position is shown in Panel B as the dashed curve. (We note that we have tried a few other sets of values as the fixed model parameters, and we have found that the observed-minus-model field, that is, &Btotai Del°w> are essentially the same except for constant offsets (not shown). That is, the results shown below are essentially unaffected by the values of the fixed model parameters.) Panel C of Figure 1 shows &Btota|, the observed field strength (solid curve in Panel B) minus the above-explained model field strength (dashed curve in Panel B): As explained above, this 85 t o t a | shows the true time-dependent perturbation in the magnetic field strength in the polar magnetosphere. As expected, the perturbation starting around the vertical line I is clearly identified in 85 tota |. That is, 8fftota| started to increase around that time, and it continued to increase till the substorm expansion onset time (vertical line II). However, unlike the GEOTAIL data (Panel D), 85totai did not show a decreasing trend after the expansion onset time; it kept its high values until ~03:45 UT (vertical line III), ~35 min after the expansion onset (vertical line II), and then decreased. Panel A of Figure 1 shows 6^ = i;obs — ^modei, where i; is the "field flaring angle" of the magnetic field in the GSM XZ plane, which means the flaring angle of the magnetic field line at the point of the satellite. That is, in the northern hemisphere, the field flaring angle is zero if the field vector is parallel to (1,0), where 1 is the x-component and 0 is the z-component; if the field vector is parallel to ( 0 , - 1 ) , the field flaring angle is 90°. This quantity was also used in Paper 1. ^obs is % of the observed field (not shown), and £modei is £ of the above-explained model field (not shown). The reason for subtracting ^modei is, as in the case of 8Btotai above, to remove the position-dependent apparent change in %. As in Paper 1, 8!; decreased while 8#totai increased, and increased while 85totai decreased. DISCUSSION AND SUMMARY With the motivation as stated in the introduction section, we have compared the time profiles of the magnetic field strength/pressure simultaneously observed in the polar lobe (by POLAR) and in the tail lobe (by GEOTAIL) for a substorm case. The result is that, while the pressure in the tail lobe increased (decreased) during the growth (expansion) phase, as in previously reported statistical pattern (e.g., Caan et al., 1978), the magnetic field strength in the polar lobe did not start decreasing until ~35 minutes after the expansion onset, similar to the case of Paper 1. Also similar to the case of Paper 1, the "field flaring angle," which decreased during the growth phase, did not start recovering until ~35 minutes after the expansion onset. This result supports the interpretation of Paper 1 that the compression effect, operative in the inner magnetosphere (including the polar magnetosphere) but not in the tail lobe, exists after the expansion onset. We note again this is not a matter of course, because it is usually thought that the field-decreasing effect of NENL after the expansion onset should propagate in the tail lobe at the very high speed (> 5000 km/s) of fast mode waves. We think this rarefaction effect did exist at POLAR for this event, but the above-stated compression effect overcame the rarefaction effect. We regard the cause of this compression effect to be the dipolarization, as in Paper 1; in support of this, the time when the increased (due to dipolarization) field strength at GOES 8 recovered to the quiet level, ~05:00 UT, was close to the time when the increased 85totai at POLAR recovered to the quiet level (see Figure 1, Panels C and F). A way to summarize the above feature is that the compression effect originates in the closed-field region (i.e., inner magnetosphere), which also affects the adjacent polar lobe, while the rarefaction takes place in the open-field region (i.e., the tail lobe and the polar lobe). The duration of the delay from the expansion onset to the maximum in 85totai at POLAR was ~28 minutes for Paper 1 's event, while it is ~35 minutes for our event; the two are fairly close. On the other hand, in Paper 1 's event,
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the maximum in 8Stotai was close in time to the maximum in \CL\. While in our event, the former was delayed from the latter (at ~03:26 UT) by ~20 minutes. It is suggested that the current disruption causes the dipolarization (e.g., Lui, 1996, and references therein); if the current disruption causes the dipolarization (which compresses the inner magnetosphere) on one hand, and enhances the auroral electrojet on the other hand, the durations of the compression and the electrojet enhancements are expected to be similar, as in the event of Paper 1; this point should be studied on a statistical basis in the future, with wider local-time coverage of the ground substorm activity. In doing so, it is also important to consider the magnetic flux transport rate from the nightside to the dayside: If it is large, the compression effect (in the nightside) would cease soon. For the event of Paper 1, 85totai at POLAR monotonically increased after the expansion onset until the maximum of SBtotai- On the other hand for our event, 55 tota | was more flat after the expansion onset, even though its maximum was ~35-min later. This difference may reflect the event-to-event difference in the ratio of the compression effect to the rarefaction effect: For our event the ratio may have been smaller than the event of Paper 1. To test this possibility with more similar events, and to study what controls the ratio between the two effects, belongs to the future research. ACKNOWLEDGMENTS CANOPUS was constructed and is operated by the Canadian Space Agency. We thank R. Lepping for the data from Wind Magnetic Field Investigation. We also thank K. Ogilvie for the data from WIND Solar Wind Experiment. We also thank H. Singer for the magnetometer data from GOES 8 satellite.
REFERENCES Caan, M. N., R. L. McPherron, and C. T. Russell, The statistical magnetic signature of magnetospheric substorms, Planet. Space Sci., 26, 269-279, 1978. Kawano, H., G. Le, C. T. Russell, G. Rostoker, M. J. Brittnacher, and G. K. Parks, Substorm-time magnetic field perturbations in the polar magnetosphere: POLAR observations, accepted for publication in Earth, Planets, and Space, 2002. Kokubun, S., T. Yamamoto, M. H. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr, 46, 7-21, 1994. Lepping, R. P., M. H. Acuna, L. E. Burlaga, W. M. Farrell, J. A. Slavin, The WIND magnetic field investigation, Space Sci. Rev., 71, 207-229, 1995. Lui, A. T. Y., Current disruption in the Earth's magnetosphere: Observations and models, J. Geophys. Res., 101, 13067-13088, 1996. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The Low Energy Particle (LEP) Experiment onboard the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 669-692, 1994. Ogilvie, K. W., D. J. Chornay, R. J. Fritzenreiter, F. Hunsaker, J. Keller, J. Lobell, G. Miller, J. D. Scudder, E. C. Sittler, Jr., R. B. Torbert, D. Bodet, G. Needell, A. J. Lazarus, J. T. Steinberg, J. H. Tappan, A. Mavretic, and E. Gergin, SWE, a comprehensive plasma instrument for the WIND spacecraft, Space Sci. Rev., 71, 55-77, 1995. Rostoker, G., J. C. Samson, F. Creutzberg, T. J. Hughes, D. R. McDiarmid, A. G. McNamara, A. V. Jones, D. D. Wallis, and L. L. Cogger, CANOPUS - a ground-based instrument array for remote sensing the high latitude ionosphere during the ISTP/GGS program, Space Sci. Rev, 71, 743-760, 1995. Russell, C. T., R. C. Snare, J. D. Means, D. Pierce, D. Dearborn, M. Larson, G. Barr, and G. Le, The GGS/POLAR fields investigation, Space Sci. Rev., 71, 563-582, 1995. Tsyganenko, N. A., Effects of the solar wind conditions on the global magnetospheric configuration as deduced from data-based field models, in Proceedings of the ICS-3 Conference on Substorms, Spec. Publ. SP-389, pp. 181-185, European Space Agency, Paris, France, 1996.
E-mail address of H. Kawano [email protected]
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PLASMA SHEET EXPANSION OBSERVED BY CLUSTER AND GEOTAIL R. Nakamura1, W. Baumjohann1, H. Noda1, K. Torkar1, T. Nagai2, M. Fujimoto2, T. Mukai3, B. Klecker4, G. Paschmann4, P. Puhl-Quinn4, H. Vaith4, J. Bogdanova4, A. Balogh5, H. Reme6, J. A. Sauvaud6, J. Quinn7, R. Torbert7, L. Kistler7 1
Space Research Institute, Austrian Academy of Sciences, Schmiedlstr. 6, 8042 Graz, Austria; rumi. nakamura §oeaw. ac.at 2 Earth and Planetary Sciences, Tokyo Institute of Technology, Tokyo, 152-8551, Japan 3 The Institute of Space and Astronautical Science, 3-1-1 Yoshinodai, Sagamihara, 229-8510, Japan 4 Max-Planck-Institut fur extraterrestrische Physik, Postf. 1312, 85741 Garchinq, Germany ^Imperial College, London, SW7 2BZ, U.K. 6 CESR/CNRS, B.P. 4346, F-31028 Toulouse Cedex 4, France 7 University of New Hampshire, SSC, Morse Hall, Durham, NH 03824, U.S.A. ABSTRACT Between 1000 and 1800 UT on October 10, 2001, when Cluster was approaching the plasma sheet from the northern lobe, Geotail traversed the southern lobe also approaching the plasma sheet during two substorm intensifications identified in the Kakioka magnetogram. In this paper we examine the temporal change of the tail configuration and evolution of the plasma sheet between 11 and 13 UT, when multiple intensifications of the first substorm took place. Changes in the tail configuration toward a dipole-like field were identified during all intensifications, suggesting plasma sheet expansion which were embedded in a gradual northward motion of the tail. We emphasize the importance of simultaneous measurements in both hemispheres in order to identify plasma sheet expansion or dipolarization, because flapping motions or local enhancements of the tail current density may mask and mimic the signatures. During a later intensification, Geotail and Cluster encountered the plasma sheet boundary layer, which is the most direct signature of expansion of the plasma sheet. Using the time difference among the Cluster four spacecraft and Geotail, the possible propagation speed of the dipolarization and spatial scale of the disturbance is discussed. In addition to the usual dawn-to-dusk electric field, significant contribution from a north-south electric field were observed both in the lobe and at the boundary of the plasma sheet associated with the dipolarization. This suggests the importance of the effects from a localized source region for dipolarization and/or reconnection process.
1. INTRODUCTION The dynamics of plasma sheet in the region between 6 and 15 RE during substorm are expected to be controlled both by current wedge signatures near the inner magnetosphere and by processes associated with near-Earth reconnection around 20 R^;. Statistical studies showed that the plasma sheet expands associated with substorm onset in the region Earthward of 15 R^ (Hones et al., 1984; Baumjohann et al., 1992). At expansion onset, the magnetic field returns to a more dipolar configuration within a longitude range that is spanned by the substorm current wedge and the plasma sheet expands. Formation of the reconnection region, on the other hand, also results in dipolarization at the earthward side of the reconnection region via pile up of the magnetic flux (Birn and Hesse, 1996). Between 1000 and 1800 UT on October 10, 2001, when Cluster was approaching the plasma sheet from the -177-
Fig. 1. Geotail and Cluster orbit between 10:00 and 18:00 UT on October 10, 2002 in (a) the X-Y plane and (b) the X-Z plane. Relative location of the four Cluster spacecraft in (c) the X-Y plane, (d) the X-Z plane, and (e) the Y-Z plane.
northern lobe, Geotail traversed the southern lobe and plasma sheet during two substorms with multiple intensifications. In this study we discuss on the temporal change of the tail configuration and evolution between 11 and 13 UT when multiple intensifications of the first substorm took place. We emphasize that: (1) thinning and thickening can only be unambiguously determined based on interhemispheric observations, which is very rare in the previous literature; (2) Cluster four-point observations enable to monitor spatial gradients and motion of the boundary and, together with Geotail observations, to discuss on the scale size of the disturbance; and (3) new Cluster instruments enable to notice the importance of the north-south electric field in the lobe as well as in the plasma sheet boundary layer, which has not been identified in previous studies. 2. OVERVIEW OF THE EVENT Figure 1 shows the location of the orbit of Geotail and Cluster between 10:00 and 18:00 UT on October 10, 2001 in (a) the X-Y plane and (b) the X-Z plane in geomagnetic solar magnetospheric coordinates (GSM), which will be used throughout this study. Geotail traversed the southern lobe and plasma sheet from premidnight to post midnight, while Cluster was located in the northern hemisphere in the premidnight sector approaching the equatorial plane from the north. Closest approach took place around 12 UT. Cluster 1 was the most northern satellite, Cluster 3 the most earthward, and Cluster 4 the most tailward and equatorward one. IMF Bz was mainly southward fluctuating between 0 and -4 nT based on ACE data between 10 and 16 UT. There were mainly two AE/AL substorms taking place with a westward electrojet of about 500 nT, based on the provisional AE and geomagnetic data plots provided from WDC-C2. Figure 2a shows the H and D components of the magnetogram from Kakioka. Positive bay enhancements associated with these two major onsets can be seen with peaks around 1210 UT and 1530 UT. These onsets can be also identified as enhancement in the AH/dt, representing the Pi2 activity, shown in Figure 2b. In this study, we concentrate on the first substorm that took place near the time when Cluster and Geotail were at closest distance. At 13:00 UT, Cluster was located X - -10.1, Y = 7.2, Z = 5.3 RE, while Geotail was located
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Fig. 2. (a) The H and D components of the magnetic field and (b) the dH/dt at Kakioka (geomagnetic latitude 26.6° ; geomagnetic longitude 207.8°) and (c) the H and D components of the magnetic field at Victoria (geomagnetic latitude 54.30°, geomagnetic longitude 296.6°).
X = —7.5, Y = 3.7, Z = —3.3 RE. All the satellites were located therefore in the premidnight sector (22.5-23 LT) with foot point of the field line near the same local time sector as Kakioka station. Figure 2c shows the H and D components of the magnetic field at Victoria, which is located 90° east of Kakioka, at postmidnight during the first substorm interval. Systematic relationships exist between the current wedge and midlatitude positive bay and perturbation in the D component (Nagai, 1987). Namely, positive (negative) D perturbation is expected at western (eastern) part of the current wedge region. By taking into account the H and D perturbation from both Kakioka at premidnight and Victoria at postmidnight, multiple intensifications, when at least one of the stations identified positive bay enhancement, were identified between 11 and 13 UT. These events are shown as vertical lines in Figure 2 indicating 11:18, 11:55, 12:20, 12:38, and 12:55 UT. We also checked the high-latitude magnetograms from the 210 MM and and GIMA chain and identified corresponding enhancements of the westward electrojet with additional 100-300 nT negative bays associated with these intensifications. The D perturbation was always positive in Kakioka and negative in Victoria, indicating that the center of the current wedge is located between these two stations. The D perturbation at 11:55 and 12:38 UT, indicated by solid lines in Figure 2, took both at Kakioka and Victoria simultaneously. Cluster observations from the fluxgate magnetometer (FGM) experiment (Balogh et al., 2001) and from the Composition and Distribution Function Analyser (CODIF) of the Cluster ion spectrometry (CIS) experiment (Reme et al., 2001) between 11 and 13 UT are shown in Figure 3a. Spin-resolution (4s) data of the X, Y, and Z components and the elevation angle, A, of the magnetic field from all the four spacecraft and proton density from SC 1, 3, and 4 are shown from top to bottom. Cluster was in the lobe until the transient enhancement of the density suggesting encounter of the boundary of the plasma sheet associated with the 1238 UT onset. The proton beta first exceeded 0.1 at 1240 and 0.5 at 1255UT. Figure 3b shows the 12-sec averaged data from the Geotail magnetic field (MGF) (Kokubun et al., 1994) and low energy particle (LEP) (Mukai et al., 1994) experiments in the same format as in Figure 3a. Geotail also stayed initially in the lobe and encountered the boundary of the plasma sheet after the 1220 UT onset, about 10 min earlier than Cluster. The ion beta reached 0.1 at 1249 UT. The dashed-dotted lines in the magnetic field plots show the Tsyganenko 96 model value with an offset so that the model field value for 11 UT agrees with Cluster 3 and Geotail data in Figure 3a and 3b, respectively. For Cluster profile the model value agrees well, -179-
Fig. 3. (a) The magnitude, X, Y, and Z components and the elevation angle, A, of the magnetic field and proton density from Cluster. The dashed-dotted line indicate the Tsyganenko 96 model with an offset indicated in the figure, (b) Same from Geotail except for ion density. The magnetic field traces in the Cluster panels are from all four spacecraft, whereas the density traces are from the SC 1, 3, and 4. The vertical lines show the midlatitude geomagnetic activation in Figure 2.
Fig. 4. (a) The X, Y, and Z components of the drift velocity of the 1-keV electrons from Cluster 3. (b) 3-min averages of the electric field perpendicular to the field determined from the electron drift velocity from the same spacecraft. The vertical lines show the midlatitude geomagnetic activation in Figure 2.
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particulary after the plasma sheet encounter. The offset value was large for Geotail (30 nT in magntitude). Yet the model obtains general trend of the current sheet consistent with the data and therefore we believe that these model profile can be used as a reference value. While the model value at Cluster location shows rather constant value, the model at Geotail shows gradual increase in B and Bx, consitent with satellite moving toward the midnight region at location with stronger tail curret. Among the activations the most pronounced onsets were at 1155 UT and 1238 UT, in which both Cluster and Geotail observed similar magnetic field perturbations. The 1155 UT onset was associated with a smooth signature of dipolarization in the lobe. The 1238 UT also was associated with dipolarization but contained short-time scale magnetic fluctuations followed by plasma sheet encounters. The strongest fluctuation was in By for this event, indicating the feature of the field aligned current. In the following we examine the detailed particle and field signatures observed during the lobe intervals and near the boundary of the plasma sheet.
3. LOBE CONVECTION AND RECONFIGURATION OF THE TAIL Both Cluster and Geotail stayed in the lobe until the 1220 UT activation. Gradual decreases in Bx were observed in both the northern and the southern hemispheres relative to the model value. This suggests that the tail was moving northward relative to the spacecraft during this interval. Note that the tail current is not an infinite current sheet in X — Y plane so that Bx is not constant and spatial gradient exists in the lobe. In fact, Cluster four-spacecraft observation in the lobe for this time interval showed a 0.4 nT/m gradient in the Z direction (larger absolute value of Bx away from the equatorial plane). The observed decrease in Bx was not steeper than -4 nT / 10 min. This corresponds to a maximum velocity of 17 km/s in the Z direction, which is a speed comparable to the average vertical plasma motion in the magnetotail (Petrukovich and Yermolaev, 2002). Simultaneous decrease in Bx at Cluster and Geotail, therefore, is very likely due to a tail flapping. Note that for these intervals, if the plasma sheet changes were determined from the Bx traces from one hemisphere only, opposite trend would have been deduced: the Cluster trace looks like expansion, while the Geotail trace looks like thinning. Overlapped with the general trend of decrease in Bx, the 1118 and 1155 UT onsets were associated with some decrease in the absolute value of Bx so that the decrease became more rapid in Cluster and slower in Geotail, which corresponds to a weakening of the taillike field. It should be stressed that whether these changes are due to flapping or a change in the current sheet can only be confirmed from the interhemispheric observations. Clear enhancements in Bz and elevation angle, which are the signature of dipolarization, were observed both at Geotail and Cluster for the 1155, 1220 and 1238 UT onset. The other onsets showed relatively weak signatures of dipolarization except for Bx signatures supporting the decrease of the current discussed in the previous paragraph. Dipolarization was most pronounced for the 1155 UT onset. Cross correlation analysis of the Bz component between 1155 and 1210 UT indicated that there was up to 16 s time difference among the Cluster spacecraft, whereas Geotail leads Cluster by 32 sec on average. Assuming that the dipolarization developed as a planar structure with constant velocity within the Cluster tetrahedron, we obtained a velocity vector of (2.5, 24, 123 km/s), which was mainly in the Z direction, as the expansion speed of the dipolarization at the Cluster location. If we assume that this dipolarization propagates symmetrically as a planar structure with the same speed in the northern and southern hemisphere, we estimated that Geotail, which is located effectively inward, should have observed dipolarization 2 min before Cluster. This is larger than the Geotail-Cluster time difference, which means that the propagation speed is too low to explain the timing difference of the observations. The small time difference between Geotail and Cluster could suggest that the spatial scale of the dipolarization initially involved at least ±4 R^ in Z direction. Figure 4a shows the X, Y, and Z components of the drift velocity of 1-keV electrons with 4 s-resolution obtained from the Electron Drift Instrument (EDI) (Paschmann et al., 2001) for Cluster 3. Since the gradient of the magnetic field is an order of magnitude smaller than the drift due to electric field, the figure mainly shows the E x B drift velocity. Since we are interested mainly in the DC component associated with substorm activation, we averaged the data over 3 min whenever more than 20 data points were available. Figure 4b shows the 3-min averaged electric field from EDI. Since the lobe field is mainly directed in the +X direction and EDI measures the perpendicular field and flow, the profile in Vy is equivalent to Ez, while Vz is equivalent to — Ey. EDI data were not available after 1240 UT when the spacecraft stayed mainly in the
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plasma sheet/plasma sheet boundary layer with large fluctuating magnetic fields. As expected in the northern lobe under the influence of usual dawn-dusk convection electric field, the flow direction in Z was mostly southward. The 1155 and 1220 UT onsets were associated with enhancement in southward flow (or duskward electric field). Large equatorward convection was observed also in previous substorm observation by Geotail in the lobe (Taguchi et al., 1998) and by ISEE 1 and 2 near the boundary of the plasma sheet (Forbes et al., 1981; Pedersen et al., 1985). Interesting to note for Cluster observation is that there were also velocity perturbations in the Y direction (or electric field in Z direction) so that the response to the substorm intensification could be seen in both the Y and Z components. For example, the major signature associated with 1118 UT onset was an enhancement in the duskward flow (or northward electric field). Explaining these these electric field variation therefore needs to take into account both the dawn-dusk and north-south component. One explanation could be that these electric field enhancement is due to dB/dt when dipolarization takes place. In fact the Ez enhancement at 1118 UT corresponds to the time when B j dropped. If we assume that dBx/dt < 0 observed by Cluster was due to the decrease in the tail current dawnward of the spacecraft (where the center of the current wedge was located), we can expect a northward excursion of the inductive electric field at the duskside edge, where Cluster is located. This profile agrees with the observation. On the other hand, the major change for 1155 and 1220 UT was in the Bz direction. Considering the fact that Geotail, which was located more Earthward and closer to the center of the tail, observed larger Bz disturbance than Cluster, Cluster was located at the tailward part of the dipolarization area. Hence we expect a positive Ey excursion, which was the case as shown in Figure 4. Effects of dipolarization in the tail have also been observed as changes in the electric field and magnetic field by EDI and FGM instruments onboard Cluster when the spacecraft altitude were at about 7 RE on field lines mapping to the central polar cap (Quinn et al., 2002). These observations suggest also that effect of the dipolarization can be propagated in the lobe/open flux region. Another explanation of the observed electric field is that the enhanced Ey is associated with the reconnection field, which is most likely located around 20RE during substorm expansion onset (Nagai et al., 1998). A possible interpretation of the enhanced equatorward flow and field disturbance in the lobe is a transition from closed flux reconnection in the plasma sheet to open lobe field reconnection and the associated change in the field configuration (Taguchi et al., 1998). In fact, for the 1220 UT onset, Geotail identified Earthward directed highly energetic ion beams, as will be discussed in detail in the next paragraph, which is supporting evidence of reconnection and an enhanced dawn-to-dusk electric field. The north-south electric field, on the other hand, could be related to the finite extension of the reconnection region.
4. PLASMA SHEET ENCOUNTER At 1230 UT Geotail encountered the plasma sheet and so did Cluster at 1240 UT. Figure 5 shows energytime spectrogram of (a) H + from Cluster 4 and (b) ions from Geotail. Two different components can be seen in the Cluster panels when the spacecraft encountered the plasma sheet. One is a high energy (> 10 keV) component which is mainly streaming Earthward, as indicated by an arrow in the second panel, and the other is a more diffuse background component up to several keV, indicated by an arrow in the third panel. Eventually the distribution becomes a more isotropic plasma sheet type distribution. Geotail spectra also shows the high energy Earthward beam component, as indicated in the second panel, and the low energy component (indicated in the third panel). Yet, there seems to be other components apparent in the Geotail panels, which are possibly due to the effect from different species. In fact, Cluster observed O + ions with up to 25 % of the proton density (0.05/cc) in the plasma sheet boundary layer during this interval (not shown). The plasma sheet encounter was also associated with fluctuating magnetic field signatures and enhanced difference in the fields among different spacecraft, as can be seen in the spread of the magnetic field traces in Figure 3a. Hence the gradient scale of the field has changed. Particularly, the 1238 UT onset was associated with short-time scale magnetic fluctuations accompanied by dipolarization. The strongest fluctuation was in By, possibly indicating field aligned currents. The Geotail perturbation occurred on a larger time scale. The difference may arise from the fact that Cluster was near the boundary of the plasma sheet, passing the local field-aligned current structure, whereas Geotail, located more inside the plasma sheet, observed the integrated effect of the field aligned currents.
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Fig, 5. (a)Energy-time spectrogram of H + in four direction (top to bottom: duskward, tailward, dawnward, sunward looking direction) obtained from Cluster 3 CODIF measurement, (b) Energy-time spectrogram of ions in four direction (same order as for Cluster 3 except that the axis labels indicate direction of the ion motion) obtained from Geotail LEP measurement.
Fig. 6. (a) The .V, V, and Z components and A of the magnetic field and proton density from Cluster 1 and 4. (b) V • B and V x B obtained from Cluster, (c) The X, Y, and Z components of trie velocity perpendicular to the field calculated using proton with energy between 40 eV and 5 keV from Cluster 1 and 4. The vertical lines indicate the first plasma sheet encounter at Cluster 1.
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To further investigate the field-aligned current signature, we examined the four Cluster spacecraft difference. Figure 6a shows the X, Y, and Z components and A of the magnetic field and the proton density from Cluster 1 and 4. The first plasma sheet encounter of Cluster 4 around 1240 UT occurred about 30 s earlier than Cluster 1. The vertical line in the figure shows the plasma sheet encounter of Cluster 1. A bipolar perturbation can be seen in the By trace in Cluster 4, while only positive By can be seen in Cluster 1. After the transient entry of the plasma sheet between 1240-1241:30 UT, both satellites returned to the lobe. A similar bipolar feature in By was again observed at the next plasma sheet encounter at 1242 UT in Cluster 4 and again, only positive By in Cluster 1. The time difference between the two spacecraft for the second plasma sheet encounter was about 4 min. Note that Cluster 1 was the northernmost spacecraft. The time difference therefore was most likely to be associated with a localized pair of field aligned-currents (current flowing mainly in the X direction) moving northward together with the plasma sheet expansion. This can be further identified from the linear estimation of V • B and V x B using the reciprocal vectors of Cluster tetrahedron (Chanteur, 1998) as shown in Figure 6b. The quality of the current density estimate using this approximation can be checked by comparing the V • B and V x 5 (Robert et al., 1998). A positive-then-negative spike in the X component of V x B, significantly larger than V • B, can be seen during the first plasma sheet encounter. This indicates a downward and an upward field-aligned current, with the latter density been about twice the former. Such a pair of field-aligned currents associated with high-energy Earthward ion beams was also observed by Fujimoto et al. (2001) as an effect of the reconnection tailward of the satellite. Note that the large By spike at Geotail in the southern hemisphere also suggests an upward field-aligned current southward or tailward of the satellite and is consistent with the Cluster observations. The observed upward current in the premidnight could be also a part of the substorm current wedge, which was located tailward (or poleward) of Geotail. Figure 6c shows the X, Y, and Z components of the velocity perpendicular to the field calculated using proton flux with energy between 40 eV and 5 keV from Cluster 1 and 4. It is interesting to note that a strong flow shear mainly directed in the Y direction was identified associated with the perturbation in By. This suggests an existence of a strong electric field (about an order larger than the electric field obtained in the lobe) near the boundary. Large flow shear co-located with the By shear near the boundary of the plasma sheet has been predicted in global simulations (Birn and Hesse, 1996) and was identified in a GeotailEquator S conjunction observation (Nakamura et al., 1999). Transient field disturbance during substorm expansion was also observed by ISEE and was interpreted as being due to impulsive reconnection (Sergeev et al., 1987). Detailed structure of particles and fields and the possible mechanism responsible for the spiky large field fluctuation will be studied in a separate paper.
5. SUMMARY Simultaneous observations in the northern and southern lobe and the plasma sheet during substorm intensifications between 11 UT and 13 UT on October, 10, 2001, were presented using Cluster and Geotail field and plasma data. Conjugate observations enabled us to monitor the global reconfiguration of the tail, while the four Cluster spacecraft could monitor the detailed structure and propagation of the disturbances. We summarize the results as following: (1) Magnetic field and electric field disturbances consistent with change toward a dipolar configuration were observed during multiple intensifications of the substorm. (2) It is essential to monitor both hemispheres to identify the dipolarization/plasma sheet expansion, since local enhancement in current density or tail flapping causes similar changes in Bx- (3) The plasma sheet encounter involved inter-hemispheric large-scale field-aligned current, but also contained transient field disturbances. (4) Both in the lobe as well as at the boundary of the plasma sheet, significant contribution from the flow in the Y direction (or electric field in Z direction) were observed. These shear components in the flow (and magnetic field) could be essential to understand the dynamics of the tail. ACKNOWLEDGEMENTS We thank V. A. Sergeev and R. Treumann for helpful discussions and comments. The authors are grateful to H. Eichelberger, G. Laky, G. Leistner, E. Georgescu, C. Mouikis for helping in the Cluster data analysis, and N. Ness and D. J. McComas for providing ACE data. We acknowledge 210MM, GIMA, CDAWeb, GSC, WDC-C2, CSDS, GCDC, ACDC for making available data used in this study. The October 10, 2001, event was selected as Event D in the Substorm Onset session in the Cluster March 2002 workshop. We thank H.
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Laakso and the workshop participants for their helpful comments. REFERENCES Balogh, A., et al., The Cluster magnetic field investigation: overview of in-flight performance and initial results, Ann. Geophys., 19, 1207, 2001. Baumjohann, W., G. Paschmann, and T. Nagai, Thinning and expansion of the substorm plasma sheet, J. Geophys. Res., 97, 1992. Birn, J., and M. Hesse, Details of current disruption and diversion in simulations of magnetotail dynamics, J. Geophys. Res., 101, 15,345, 1996. Chanteur, G., Spatial interpolation for four spacecraft: theory, in Analysis Methods for Multi-Spacecraft Data, ed. G. Paschmann and P. W. Daly, pp. 349, ESA Publications Division, Noordwijk, 1998. Forbes, T. G., E. W. Hones, Jr., S. J. Bame, J. R. Asbridge, G. Paschmann, N. Sckopke, C. T. Russell, Substorm-related plasma sheet motions as determined from differential timing of plasma changes at the Isee Satellites, J. Geophys. Res., 86, 3459, 1981. Fujimoto, M. T. Nagai, N. Yokokawa, Y. Yamade, T. Mukai, Y. Saito, and S. Kokubun, Tailward electrons at the lobe-plasma sheet interface detected upon dipolarizations, J. Geophys. Res., 106, 21255, 2001. Hones, E. W., T. Pytte, and H. I. West, Associations of geomagnetic activity with plasma sheet thinning and expansion: A statistical study, J. Geophys. Res., 89, 1984. Kokubun, S., T. Yamamoto, M. H. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The Geotail magnetic field experiment, J. Geomagn. Geoelectr., 46, 7, 1994. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particle (LEP) experiment onboard the Geotail satellite, J. Geomagn. Geoelectr., 46, 669, 1994. Nagai, T., Field-aligned currents associated with substorms in the vicinity of synchronous orbit. II - GOES 2 and GOES 3 observations, J. Geomagn. Geoelectr., 92, 2432, 1987. Nagai, T., M. Fujimoto, Y. Saito, S. Machida, T. Terasawa, R. Nakamura, T. Yamamoto, T. Mukai, A. Nishida, and S. Kokubun, Structure and dynamics of magnetic reconnection for substorm onsets with Geotail observations, J. Geophys. Res., 103, 4419, 1998. Nakamura, R., G. Haerendel, W. Baumjohann, A. Vaivads, H. Kucharek, B. Klecker, E. Georgescu, J. Birn, L. M. Kistler, T. Mukai, S. Kokubun, P. Eglitis, L. A. Frank, J. B. Sigwarth, Substorm observations in the early morning sector with Equator-S and Geotail, Ann. Geophys., 17, 1602, 1999. Paschmann,G., et al., The electron drift instrument on Cluster: overview of first results, Ann. Geophys., 19, 1273, 2001. Pedersen, A., C. A. Cattell, C.-G. Falthammer, K. Knott, P.-A. Lindqvist, R. H. Manka, and F. S. Mozer, J. Geophys. Res., 90, 1231, 1985. Quinn, G. Paschmann, E. Georgescu, H. Vaith, H. Frey, Convection Transients in the Polar Cap During Substorms, Eos Trans. AGU, in press, Fall Meeting suppl., 2002. Reme, H., et al., First multispacecraft ion measurements in and near the Earth's magnetosphere with the identical Cluster ion spectrometry (CIS) experiment, Ann. Geophys., 19, 1303, 2001. Robert, P., M. W. Dunlop, A. Roux, and G. Chanteur, Accuracy of current density determination, in Analysis Methods for Multi-Spacecraft Data, ed. G. Paschmann and P. W. Daly, pp. 395, ESA Publications Division, Noordwijk, 1998. Sergeev, V. A., V. S. Semenov, and M. V. Sydneva, Impulsive reconnection in the magnetotail during substorm expansion, Planet Space Sci., 35, 1199, 1987. Taguchi, S., M. Kiyohara, T. Mukai, T. Yamamoto, M. Nose, Y. Saito, and S. Kokubun, Geotail observations of north-south plasma velocity enhancements in the lobe near substorm expansion phase onset, Geophys. Res. Lett, 24, 4125, 1998.
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DIFFERENCE BETWEEN EARTHWARD AND TAILWARD FLOWS IN THEIR DEPENDENCES ON GEOMAGNETIC AND IMF CONDITIONS A. Ieda1 2, T. Mukai1, S. Machida3, J.-H. Shue4, S.-I. Ohtani4, T. Nagai5, and Y. Saito1 1
2
Institute of Space and Astronautical Science, Sagamihara, Kanagawa 229-8510, Japan Now at Solar-Terrestrial Environment Laboratory, Nagoya University, Toyokawa, Aichi ^^2-8501, Japan 3 Department of Geophysics, Kyoto University, Kyoto 606-01, Japan 4 The Johns Hopkins University Applied Physics Laboratory, Laurel, MD 20723-6099, USA 5 Earth and Planetary Sciences, Tokyo Institute of Technology, Tokyo 152-8551, Japan
ABSTRACT Earthward and tailward perpendicular fast flows in the plasma sheet were studied and their differences in response to geomagnetic conditions and to interplanetary magnetic field (IMF) conditions are discussed. We first identified the plasma sheet from 3.5 years of Geotail plasma and magnetic filed observations between 8 and 32 RE down the tail. We then studied occurrence rates of fast flows during geomagnetically quiet and active intervals as identified by Kp and ASY indices, and during northward and southward IMF intervals. As a result, both earthward and tailward flows were observed more often during active or southward IMF intervals than during quiet or northward IMF intervals, as expected. On the other hand, we found that there is a difference between earthward and tailward flows: Dependences on geomagnetic conditions are more evident in tailward flows than in earthward flows. We further discuss that tailward flows indicate the substorm expansion phase better than earthward flows do.
Introduction Reconnection in the magnetotail yields a pair of earthward flow with northward magnetic field and tailward flow with southward field. Thus earthward or tailward flows are often thought to indicate that a satellite is located earthward or tailward of a neutral line. Near the times of the substorm onsets identified by ground magnetic field, Nagai et al. (1998) statistically found that the earthward flow with northward magnetic field is seen earthward of 20-30 RE and the tailward flow with southward field is seen beyond this region. This result can be interpreted as indicating that reconnection occurs somewhere between 20-30 RE at the onsets, on the assumption that most fast flows at onsets are caused by near-Earth reconnection. On average, however, many more earthward flows than tailward flows are observed inside ~ 30RE (e.g., Paterson et al., 1998). Thus at least some fractions of earthward flows are not likely to be substorm expansion phase signatures. These earthward flows may come from the distant neutral line that is activated during southward IMF intervals. Due to these additional earthward flows, earthward flows could be statistically more associated with IMF conditions than tailward flows are. In other words, tailward flows may better mark the substorm expansion phase than earthward flows do. Thus geomagnetic and IMF conditions may have different significance between earthward and tailward flows, although geomagnetic and IMF conditions are basically correlated with each other. The purpose of this study is to find this difference by comparing earthward and tailward flows under different geomagnetic and IMF conditions.
Identification of IMF and Geomagnetic Conditions We identified IMF north (Bz > 1) and south (Bz < —1) intervals by OMNI 1-hour averages of IMFBz data in the GSM coordinates, while we used Kp and ASY indices to identify geomagnetically quiet
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Table 1. Plasma Sheet Samples (number ( xlO 3 ) fraction (%)) \XAGSM\(RE)
Average \X\(RE) All intervals All IMF IMF-BZ > 1 IMF-BZ < - 1 Active Quiet
I3-13 10.5 164 (100) 154 (94) (29) 45 (33) 50 (14) 23 (26) 43
13-20 16 5 172 (100) 168 (98) 52 (31) (33) 56 (16) 28 (31) 53
20-27 23.7 177 (100) (96) 169 52 (31) 55 (33) (12) 22 (28) 49
27-32 29.1 188 (100) 184 (98) (34) 63 53 (29) 22 (11) (24) 45
and active conditions in the polar ionosphere, respectively. Longitudinally asymmetric disturbance (ASY) indices provide longitudinal asymmetries in the H and D components of the ground magnetic field among 6 mid-latitude stations with 1 min time resolution (Iyemori and Rao, 1996). The stations (including Dst stations) are located roughly every 60 degrees in longitude. We assumed that the ASY indices were sensitive to a formation of a major substorm current wedge, while the ASY indices are not very sensitive to localized geomagnetic activities at high-latitudes. The Kp index indicates geomagnetic variations in the sub-auroral regions. Kp is likely to be better than ASY index in identifications of quiet intervals in the polar ionosphere, because the Kp stations are in higher latitudes. On the other hand, Kp is not a suitable index to identify the substorm expansion phase that has much shorter (10-30 min) time scales than the temporal resolution (3 hours) of Kp. However, the long time resolution of Kp is not a crucial problem in identification of quiet intervals. For the reasons above, we used Kp < 1 (Kp = 0, 0+, and 1-) to identify quiet intervals. To exclude fast flows in the tail that occurred just before the end of a Kp quiet interval and were associated with an enhancement in the next Kp, the last 30 min was excluded when the next Kp was relatively high (Kp > 1). We assumed that major substorm expansion phases were excluded in these "quiet intervals" but that weak or localized disturbance was sometimes included. To define geomagnetically active intervals we first calculated the maximum variations in ASY-H (dH) and ASY-D (dD) during 30 min (± 15 min) for every 1-min time and defined dASY as higher one of dH or dD. We defined active intervals by dASY > 10 nT. We assumed that our "active intervals" include major substorm expansion phases with time resolution of 15-30 min and are less sensitive to the substorm recovery phase when geomagnetic variations are slower.
Identification of the Plasma Sheet The plasma sheet was identified with 12-s plasma (Mukai et al, 1994) and magnetic field (Kokubun et al, 1994) data from the Geotail spacecraft for 3.5 years from October 30, 1994 through April 30, 1998. (Geotail plasma data after April 30, 1998 are still under calibration). Geotail orbited in the solar-ecliptic plane and had perigee of ~ 9RE and apogee of ~ 31RE- We used the aberrated GSM coordinates with an aberration of 4 degrees. Tailward-moving plasmoids are known to be observed most often around YAGSM = 3RE (Ieda et al., 1998). Around the substorm onsets, fast flows also tend to be observed in the pre-midnight region (Nagai et al, 1998). Thus we concentrated on studying the plasma sheet in —4 < YAGSM < WRE- We further divided the plasma sheet samples into the four regions of 8 < \X\ < 13RE, 13 < \X\ < 20RE, 20 < \X\ < 27RE, and 27 < \X\ < 32RE, SO as the sampling numbers are close to each other. The plasma sheet was identified by 0 > 0.1, where 0 is the ratio of the plasma pressure to the magnetic pressure. We assumed the ratio of ion to electron temperature as five (Slavin et al., 1985). As a result, we found 701 thousand plasma sheet samples (equivalent to 97 days) as summarized in Table 1, where fractions (%) in IMF northward and southward intervals were normalized by all IMF intervals, and other fractions were normalized by all intervals.
Statistical Results Bz
Figure 1 shows the occurrence rates (%) of earthward perpendicular fast flows (VpCrP:X > 300 km/s and > 0) on the left, and tailward perpendicular fast flows (VpCrPjX < —300 km/s and Bz < 0) on the
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right against XAGSM- ^porp,x is the X component of the velocity perpendicular to the magnetic field. Each panel shows fast flows for the all intervals (thin line with filled circles), for the northward IMF (thick line with squares) and the southward IMF (thick line with circles) intervals, and for the quiet (dashed line with squares) and the active intervals (dashed line with circles). For all intervals (thin lines), the occurrence rates of both earthward and tailward flows increased with radial distance. This is consistent with the results by Paterson et al. (1998). Earthward flows were observed 10 times more than tailward flows between 20 and 30 RE, where the near-Earth reconnection is supposed to occur. We further studied their dependences on geomagnetic and IMF conditions. Both earthward and tailward flows were more often observed during active intervals than during quiet intervals, which have been previously shown inside ~ 22Rg (e.g., Angelopoulos et al., 1994). Both flows were also more often found during southward IMF than during northward IMF intervals, as expected. However, there are differences in the fast flow dependences between geomagnetic and IMF conditions. Tailward flows appear much (three times around 30 RE) more different between quiet and active intervals than between IMF northward and southward intervals. On the other hand, the dependences on IMF and geomagnetic conditions are close for earthward flows, or even earthward flows between 13 and 20 RE appear rather dependent of IMF than geomagnetic conditions. Thus tailward flows are relatively (three times around 30 RE) more sensitive to geomagnetic conditions than earthward flows are.
Fig. 1. The occurrence rates (%) of fast flows against XAGSM in —4 < VAGSM < IORE of the plasma sheet, which was identified by
0 > 0.1. Earthward perpendicular fast flows (VpOrP:X > 300 km/s, Bz > 0) are shown on the left and tailward perpendicular fast flows (Vpcrp.x < -300 km/s, Bz < 0) are shown on the right. Both panels show the results for all, northward and southward IMF, quiet and active intervals as marked.
We also studied slower flows (100200 km/s and 200-300 km/s, not shown). Their occurrence rates had similar but weaker characteristics as fast flows havein Figure 1, indicating that the selection criterion of 300 km/s for fast flows is not qualitatively critical to our results. Note that each bin typically has 20 x 103 to 60 x 103 samples, as shown in Table 1. Thus the occurrence rates of 0.1% contain 20-60 fast flow 12-s samples. If the fast flow events typically continue 1 min, they correspond to 4-12 events. Thus occurrence rates below ~ 0.1% are not very reliable but just indicate that fast flows are rare.
Tailward flows appear more dependent of geomagnetic conditions than of IMF conditions in the right panel of Figure 1. This result is partly because less active intervals were identified than southward IMF intervals with our criteria. Relative importance of IMF to geomagnetic conditions is dependent of the selection criteria of conditions. However, criteria-dependent factors in the relative importance of IMF to geomagnetic conditions can be mostly canceled when comparing earthward and tailward flows. We emphasize that we have compared earthward and tailward flows and they were not separately discussed in this study.
Discussion There are two possible reasons why fast flows inside ~ 3O.RE are more often observed during southward than northward IMF intervals. One is the large-scale convection due to enhanced dayside reconnection and subsequent reconnection in the distant neutral line, presumably around lOOJR^. This only increases the number of earthward flows, especially around the tail axis. The other is near-Earth reconnection around 20-30 RE, which is supposed to be associated with substorms that occur more often during southward IMF
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intervals. This increases the number of both earthward and tailward flows in the substorm expansion phase, and predominantly increases that of earthward flows in the recovery phase because neutral lines retreat beyond the Geotail locations. Thus most tailward flows in the near tail are supposed to be found during the substorm expansion phase. This expectation is consistent with our result that tailward flows are more dependent of geomagnetic conditions than IMF conditions, when compared to earthward flows. The frequent observation of tailward flows during active intervals is consistent with the previously found close association between tailwardmoving plasmoids and auroral brightenings [Ieda et al., 2001]. On the other hand, earthward flow can be due to either the distant or the near-Earth neutral line. Our result indicates that earthward flows are relatively more dependent of IMF than of geomagnetic conditions, when compared to tailward flows. Thus significant number of earthward flows are observed during southward IMF but not in the substorm expansion phase. One possible interpretation of this may be that such earthward flows come from the distant neutral line that is activated by southward IMF without substorms. Another possible interpretation is that earthward flows are also observed in the recovery phase [e.g., Baumjohann et al., 1999] and/or growth phase with southward IMF. Earthward flows in the expansion phase were not evident when compared to tailward flows, probably because there are earthward flows during the intervals other than the expansion phase. Summary On the basis of Geotail plasma and magnetic field observations, we have studied occurrence rates of fast flows in the plasma sheet from |X| = 8 to 32Rg. We found that earthward and tailward flows respond to geomagnetic and IMF conditions in different ways. In summary: (1) Both earthward and tailward flows were more often observed during active or southward IMF intervals than during quiet or northward IMF intervals, as expected. (2) On the other hand, relative importance of geomagnetic conditions to IMF conditions are different between earthward and tailward flows. Tailward flows are much (three times around 30 RE) more sensitive to geomagnetic conditions than earthward flows are. ACKNOWLEDGEMENTS This work was performed while A.I. held a research fellowship of the Japan society for the promotion of science for young scientists. The Kp and ASY index were provided through WDC-D2, Kyoto. OMNI data was provided by NSSDC. A.I. thank D. H. Fairfield, K. Liou, and T. Iyemori for their valuable comments. REFERENCES Angelopoulos, V., C. F. Kennel, F. V. Coroniti, R. Pellat, M. G. Kivelson, et al., Statistical characteristics of bursty bulk flow events, J. Geophys. Res.,99, 21,257-21,280, 1994. Baumjohann, W., M. Hesse, S. Kokubun, T. Mukai, T. Nagai, and A. A. Petrukovich, Substorm dipolarization and recovery, J. Geophys. Res., 104, 24,995-25,000, 1999. Ieda, A., S. Machida, T. Mukai, Y. Saito, T. Yamamoto, et al., Statistical analysis of the plasmoid evolution with Geotail observations, J. Geophys. Res.,103, 4453-4465, 1998. Ieda, A., D. H. Fairfield, T. Mukai, Y. Saito, S. Kokubun, K. Liou, C.-I. Meng, G. K. Parks, and M. J. Brittnacher, Plasmoid ejection and auroral brightenings, J. Geophys. Res.,106, 3845-3857, 2001. Iyemori, T., D. R. K. Rao, Decay of the Dst field of geomagnetic disturbance after substorm onset and its implication to storm-substorm relation Ann. Geophysicae,14, 608-618, 1996. Kokubun, S., T. Yamamoto, M. H. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The Geotail magnetic field experiment, J. Geomagn. Geoelectr.,46, 7-21, 1994. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa, et al., The Low Energy Particle (LEP) experiment onboard the Geotail satellite, J. Geomagn. Geoelectr.,46, 669-692, 1994. Nagai, T., M. Fujimoto, Y. Saito, S. Machida, T. Terasawa, et al., Structure and dynamics of magnetic reconnection for substorm onsets with Geotail observations, J. Geophys. Res.,103, 4419-4440, 1998. Paterson, W. R., L. A. Frank, S. Kokubun, and T. Yamamoto, Geotail survey of ion flow in the plasma sheet: Observations between 10 and 50 RE, J. Geophys. Res.,103, 11,811-11,825, 1998. Slavin, J. A., E. J. Smith, D. G. Sibeck, D. N. Baker, R. D. Zwickl, and S. -I. Akasofu, An ISEE 3 study of average and substorm conditions in the distant magnetotail, J. Geophys. Res.,90, 10,875-10,895, 1985. -189-
THE LOADING-UNLOADING PROCESS IN THE MAGNETOTAIL DURING A PROLONGED STEADY SOUTHWARD IMF Bz PERIOD T. Nagai1, R. Nakamura2, T. Hori3, and S. Kokubun4 'Tokyo Institute of Technology, Tokyo 152-8551, Japan Space Research Institute, Austrian Academy of Sciences, Schmiedlstr. 6, A-8042, Graz, Austria The Johns Hopkins University Applied Physical Laboratory, MD, U.S.A. 4 The University of Tokyo, Tokyo 113-0033, Japan
2
ABSTRACT The interplanetary magnetic field (IMF) Bz was continuously southward for 16 hours on April 18, 2002. During this period, the spacecraft Geotail was in the northern tail lobe at radial distances of 24-28 RE, and it observed sequences of increasing and decreasing in the magnetic field intensity in association with substorms. The increases in the magnetic field took approximately 40 minutes, and the decreases took approximately 50 minutes. The time scale of the increase is consistent with the typical time scale of the substorm growth phase, which usually begins with a clear southward turn of the IMF Bz. The time scale of the decrease agrees well with the typical time scale of the substorm expansion phase, although the substorm expansion phase frequently starts with a northward turn of the IMF Bz. It is suggested that the magnetotail has an intrinsic time scale for the loading-unloading process for substorms irrespective of the IMF Bz conditions.
Introduction It is known that the tail lobe magnetic field increases in the substorm growth phase and decreases in the substorm expansion phase (e.g., Caan et al., 1975, Fairfield et al., 1981, Kistler et al, 1993, Nagai et al., 1997). This process is interpreted as the loading and unloading of energy in the magnetotail for substorms (e.g., Baker et al., 1996). The loading process usually starts when the interplanetary magnetic field (IMF) Bz turns southward, and the major onset of a substorm occurs after the IMF Bz continues to be southward for 40-60 minutes. The unloading process usually corresponds to the expansion phase of the substorm, and it typically continues for 40 minutes on the ground. Since the major onset is frequently associated with a northward turn of the IMF Bz, the duration of 40 minutes for the unloading process is obtained for the northward IMF Bz period (e.g., Caan et al., 1975). On April 18, 2002, the IMF Bz was almost continuously southward, and substorms took place quasiperiodically. The spacecraft Geotail stayed almost continuously in the northern tail lobe at radial distances of 2428 RE. The observations provided a unique opportunity for examining the dynamics of the magnetotail for the steady southward IMF Bz conditions. We focus on the time scales of the loading-unloading process in this paper.
Observations The upper panel of Figure 1 presents the IMF and the solar wind velocity observed with the spacecraft ACE for the period from 1800 UT on April 17, 2002, through 1200 UT on April 19, 2002. ACE was located near (+224, +22, -23 RE) in GSM coordinates. The data are shifted 50 minutes for the propagation time. The IMF Bz was continuously southward for the period of 0100-1800 UT on April 18, 2002. The spacecraft Wind observed almost identical behaviors in the solar wind. The lower panel of Figure 1 presents the magnetic field observation obtained with Geotail (Kokubun et al., 1994). Geotail entered the magnetotail near 2100 UT on April 17, 2002. According
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to plasma measurements made with the low-energy plasma experiment (LEP) on Geotail (Mukai et al., 1994), Geotail stayed in the northern tail lobe after 0200 UT on April 18, 2002, except for a brief period in the tail lobe/plasma sheet boundary close to 1150 UT. For the period of 0400-1800 UT, there were five increase-decrease sequences in the total magnetic field, and each was associated with a southward turn of Bz in the tail lobe field. These magnetic field behaviors are known as traveling compression regions (TCRs) (Slavin et al., 1984; 1992).
Fig. 1. The upper panel presents the IMF and the solar wind velocity observed with the spacecraft ACE for the period from 1800 UT on April 17, 2002, to 1200 UT on April 19, 2002. ACE is located near (+224, +22, -23 RE) in GSM coordinates. The lower panel presents the magnetic field observation from Geotail in GSM.
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In order to find substorm activities, we examined ground magnetic field data from the mid-latitude stations Fredericksburg, Tucson, Honolulu, Kakioka, Urumqi, and Hermanus and from the auroral-zone stations Leirvogur, Narssarssuaq, Poste-de-la-Balaine, Fort Churchill, Yellowknife, College, Tixie Bay, and Kiruna. We also examined particle observations from the Los Alamos National Laboratory spacecraft and magnetic field observations from GOES 8 and GOES 10 at geosynchronous orbit. There were at least five major substorm onsets, at 0530, 0800, 1130, 1408, and 1632 UT for the period of 0300-1800 UT on April 18, 2002. These five onsets correspond well to the TCRs observed with Geotail in the magnetotail.
Fig. 2. The total pressure (magnetic pressure + plasma pressure) in the magnetotail observed with Geotail (thick curve) and the solar wind dynamic pressure observed with ACE (thin curve) on April 18, 2002. Figure 2 shows the total pressure (magnetic pressure + plasma pressure) in the magnetotail observed with Geotail and the solar wind dynamic pressure observed with ACE on April 18, 2002. The solar wind data are shifted 50 minutes for the propagation time. The dynamic pressure changed significantly for the period of 00000300 UT, so we do not discuss substorm-associated pressure changes in the magnetotail. For the 0530, 0800, 1130, and 1408 UT substorms, since the total pressure in the magnetotail did not follow the solar wind dynamic pressure well for most intervals, variations in the total pressure are attributable to the loading-unloading processes in the magnetotail. For the 1632 UT substorm, the time scale of the decrease in total pressure is used. Just after each substorm onset, total pressure showed a significant peak, which was produced by a spike in Bx. These peaks are probably caused by large pressure pulses, which are produced by high-speed tailward flows inside the plasma sheet (e.g., Nakamura et al., 1999). Since the peaks are considered to be transient phenomena in TCRs, we neglect them in our analysis of the loading-unloading process.
Fig. 3. The lower panel presents the changes in total pressure relative to the substorm onset (the zero epoch corresponds to each substorm onset), and the upper panel presents solar wind electric field variations.
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Figure 3 shows the changes in total pressure relative to substorm onset (the zero epoch corresponds to each substorm onset) as well as solar wind electric field variations. The solar wind electric field is calculated as Esw = Bs x V (Bs = 0 for the IMF Bz > 0, and Bs = -Bz for the IMF Bz < 0). The figure shows that the total pressure increased approximately 40 minutes before the substorm onset and reached to the lowest level approximately 50 minutes after the substorm onset. It is important to note that the solar wind electric field stayed at different levels for different events and did not show a correlation with any change in total pressure.
Discussion The Geotail observations on April 18, 2002, provided an opportunity for examining substorm processes in the magnetotail for the prolonged southward IMF Bz period. In this paper, we study the time scales of the loadingunloading process in the magnetotail for substorms. The loading process has a time scale of 40 minutes, even for different solar wind electric field levels, and it seems to start without any significant changes in solar wind conditions. Irrespective of the IMF Bz sign, the unloading process has a time scale of 50 minutes. These results are somewhat unexpected. A definite start time for the unloading process is not expected for a prolonged southward IMF Bz. If there is any critical limit to the total pressure in the tail for substorm onsets, we would expect the time scale of the loading process to be short for the high solar wind electric field. Furthermore, we would expect the unloading process to continue for a long time for a prolonged southward IMF Bz. The present results suggest that the magnetotail has an intrinsic time scale for the loading-unloading process for substorms. These results are still highly preliminary; however, no other periods in the Geotail observations are good for testing them. ACKNOWLEDGMENTS CDAWeb was used for the ACE, Wind, Los Alamos National Laboratory spacecraft, and GOES data. We thank N. Ness and D. J. McComas for the ACE data. We thank T. Mukai for providing the LEP data from Geotail. The ground magnetic field data are provided by WDC-C2, Kyoto University. REFERENCES Baker, D. N., T. I. Pulkkinen, V. Angelopoulos, W. Baumjohann, and R. L. McPherron, Neutral line model of substorms: Past result and present view, J. Geophys. Res., 101, 12975-13010, 1996. Caan, M. N., R. L. McPherron, and C. T. Russell, Substorm and interplanetary magnetic field effects on the geomagnetic tail lobes, J. Geophys. Res., 80, 191-194, 1975. Fairfield, D. H., R. P. Lepping, E. W. Hones, Jr., S. J. Bame, and J. R. Asbridge, Simultaneous measurements of magnetotail dynamics by IMP spacecraft, J. Geophys. Res., 86, 1396-1414, 1981. Kistler, L. M., W. Baumjohann, T. Nagai, and E. Mobius, Superposed epoch analysis of pressure and magnetic field configuration changes in the plasma sheet, J. Geophys. Res., 98, 9249-9258, 1993. Kokubun, S., T. Yamamoto, M. H. Acuna, et al., The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46,7-21,1994. Mukai, T., S. Machida, Y. Saito, et al., The low energy particle (LEP) experiment onboard the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 669-692, 1994. Nagai, T., T. Mukai, T. Yamamoto, A. Nishida, S. Kokubun, and R. P. Lepping, Plasma sheet pressure changes during the substorm growth phase, Geophys. Res. Lett., 24, 963, 1997. Nakamura, R., L. F. Bargatze, T. Mukai, et al., Response of the midtail electric field to enhanced solar wind energy input,/ Geophys. Res., 104, 17299-17310, 1999. Slavin, J. A., E. J. Smith, B. T. Tsurutani, et al., Substorm associated traveling compression regions in the distant tail: ISEE-3 geotail observations, Geophys. Res. Lett., 11, 657-660, 1984. Slavin, J. A., M. F. Smith, E. L. Mazur, et al., ISEE 3 plasmoid and TCR observations during an extended interval of substorm activity, Geophys. Res. Lett., 19, 825-828, 1992.
E-mail address [email protected]
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MAGNETOTAIL DEFLATION: GEOTAIL OBSERVATIONS H. Nakai'andY. Kamide2 'ibaraki High School, 12-1 Shinjo-cho, Ibaraki, Osaka, Japan Solar-Terrestrial Environment Laboratory, Nagoya University, 3-13, Honohara, Toyokawa, Aichi, Japan
2
ABSTRACT A sudden decrease in the total pressure in the mid-magnetotail associated with a substorm was designated as magnetotail deflation by Nakai and Kamide (2003). Utilizing UVI auroral image data from Polar and magnetic field and particle data from Geotail, a typical example of the magnetotail deflation is examined in detail. A series of auroral images during this event show that a precursory breakup first occurred, which was then followed by a major breakup. Although the foot point of Geotail was mapped within the auroral bulge formed during the precursory breakup, only minor changes in the magnetic field were observed at Geotail. In contrast, a clear magnetotail deflation occurred in conjunction with the major breakup. It is suggested that the major breakup is preferably initiated in the mid-tail premidnight region. INTRODUCTION When the interplanetary magnetic field (IMF) is directed southward, the Earth's magnetic flux is transported from the dayside magnetosphere to the magnetotail, resulting in a gradual increase in both the diameter and total pressure of the magnetotail. Storing magnetic energy in the magnetotail, the magnetic configuration evolves toward an unstable state. Stored energy is eventually unloaded, causing a sudden decrease in the tail radius and magnetic field magnitude (Fairfield and Ness, 1970; Maezawa, 1975). These processes can be observed typically as a sudden decrease in the magnetotail total pressure and the magnetic field dipolarization. Nakai and Kamide (2003) termed these phenomena as the "magnetotail deflation." Using data from ISEE 1, they showed that a deflation event tends to occur when the lobe magnetic field reaches a critical value that depends on the solar wind pressure and the Dst index. We have recently been able to monitor deflation events more closely than ever using data from Geotail and other spacecraft. Further, auroral images obtained by Polar have made it possible to investigate the relationship between auroral breakups and magnetotail deflations. Nakai et al. (2002) (Paper I, hereinafter) examined three deflation events, showing that auroral breakups can be classified into two types, depending on whether they are or are not associated with the deflation event. They argued that these two categories should be attributed to the pseudo and major breakups, respectively. In this paper we present another clear-cut example of the deflation event for a close examination and discuss the substorm-associated dynamics in the near-Earth and mid-distant magnetotail. OBSERVATIONS Figure 1 shows the three components of the magnetic field (left) and plasma parameters (right) obtained by Geotail from 0600 to 1000 UT on January 12, 1997. The elevation angle of the magnetic field is represented by the ratio of the Bz component to the magnetic field magnitude in the left, bottom panel. The total pressure, the ion density, and the X and Y components of the plasma bulk velocity are shown by solid lines in the right panels. The j3 value and the temperature are plotted by dashed lines in the top and middle panels, respectively. The location of the spacecraft is shown beneath the right panels and in Figure 2. The occurrence of a clear magnetotail deflation is seen at 0807 UT, where the total pressure abruptly began to decrease in association with the magnetic field dipolarization. The ion density (and temperature) continued increasing (and decreasing) during the first half of the time span concerned, indicating that the plasma sheet was thinning. Geotail entered the lobe region just before the onset of the deflation, and returned to the plasma sheet at 0855 UT. Thereafter the temperature tended to be higher than before the deflation, and high-speed earthward ion flows were frequently observed till about 1200
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Fig. 1. Variations in the magnetic field and plasma parameters obtained by Geotail.
UT (data are not shown between 1000 and 1200 UT). These features were commonly observed in the three examples reported in Paper I. Figure 3 shows variations in solar wind parameters, where PD denotes the solar-wind dynamic pressure. The data were plotted by shifting the traveling time (-26 minutes) calculated on the basis of the distance along the X coordinate from Wind to Geotail. The solar wind speed and the dynamic pressure were rather stable for the period concerned. Selected UVI auroral images from Polar are shown in Plate 1. An auroral breakup started at -0727 UT in the premidnight sector, and subsequently expanded both in poleward and longitudinal directions, creating a medium-scale auroral bulge, of which the maximum width was -15 degrees in geomagnetic latitude at 0746 UT. Geotail was located at (-28.82, 6.02, -3.05) RE. The foot point of Geotail on the northern hemisphere is estimated to be at (68 CGM latitude, 2300 MLT), using Tsyganenko's 89c magnetic field model (Tsyganenko, 1989). Thus, it is understood that the foot point of Geotail was within the auroral bulge. The auroral bulge rotated anticlockwise between 0749 and 0804 UT. Another breakup began at 0807 UT about 2100 MLT near the western edge of the existing auroral bulge. A new bulge grew, reaching a 20 degree width at 0823 UT. The second breakup appears to have occurred in conjunction with the first one. The first and second breakups are termed as precursory and major breakups, respectively. The onset times of these breakups are marked by thin and thick vertical dashed lines in Figure 1. It is noticed that the deflation event occurred almost simultaneously with the major breakup, while it did not occur in conjunction with the precursory breakup. Figure 4 shows magnetic field variations at GOES 8 and 9 in local s/c coordinates. The locations of the spacecraft at 0807 UT are shown in Figure 2. The Hp component abruptly increased at 0731 UT and 0740 UT at GOES 9 and 8, respectively. Comparing the 0727 UT image in Plate 1 with Figure 2, it is noticed that the Fig. 2. Locations of the spacecraft. precursory breakup was initiated in the vicinity of GOES 9.
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Fig. 3. Variations in the IMF and solar wind.
Fig. 4. Magnetic field variations observed by GOES 8 and 9.
Therefore, the time lag between the two spacecraft for the dipolarization is attributed to the local times at the spacecraft. Although GOES 9 was close to Geotail in local time, the magnetic field at geosynchronous orbits did not change evidently in association with the major auroral breakup. Therefore it is inferred that the onset region of the precursory breakup was located near geosynchronous orbit, while the source region of the major breakup was more distant from Earth, probably in the mid-magnetotail. After the precursory breakup the magnetic field fluctuated, and high-speed tailward ion flows appeared between 0724 UT and 0753 UT at Geotail. The flow directions were tailward and/or dawnward, implying a vortical motion of plasma. It should be noted that the plasma sheet was thinned further during this period. DISCUSSIONS Auroral breakups on January 12, 1997, have been studied in detail. Precursory and major breakups occurred sequentially. A deflation event was observed by Geotail in the premidnight, mid-tail region just after the major breakup, while it did not occurred in conjunction with the precursory breakup. The auroral bulge rotated eastward during the precursory breakup. The major breakup was initiated near the western edge of the existing bulge. These features are commonly observed in the events examined in Paper I. Plasma sheet expansions occurred within 10 minutes after the onsets of the deflation in Paper I. In contrast to this, the reentry of Geotail into the plasma sheet occurred about 47 minutes after the onset of the deflation. This difference is probably attributed to the difference in the radial distances of Geotail: Geotail was located further from Earth in the present event than it was at the events studied in Paper I. The definition of the pseudobreakup has relied primarily on the size of auroral bulges, while the absence of the unloading process in the magnetotail has been assumed to be one of the important signatures for the pseudobreakup (see Paper 1). Paper 1 and the present study demonstrate, however, that well-developed bulges do not necessarily assure the occurrence of the unloading process. To distinguish definitely the pseudobreakup from the major breakup it is useful to define them as the auroral breakups associated or not associated with a deflation, respectively. Tailward high-speed ion flows were observed in the mid-tail plasma sheet for about 29 minutes after the first onset, implying that particles were accelerated in the near-Earth region. A small reduction in the total pressure is probably attributed to the rarefaction behind the fast ion flows. Haland et al. (1999) reported the appearance of tailward flows near the lobe/plasma sheet boundary about 10 minutes after the 1656 UT precursory breakup on December 10, 1996. Since the Bz component was negative at the front of the ion flows as shown in the left panel
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Plate 1. UVI images obtained by Polar at 0727 UT, 0746 UT, 0804 UT, 0807 UT, and 0823 UT on January 12, 1997. in Figure 1, it is quite possible to interpret the flows as a small plasmoid. As mentioned above, however, the temperature was decreasing during the period the flows were observed. Since magnetic field energy is converted to thermal energy in association with the plasmoid formation, the decreasing temperature is inconsistent with the view that a plasmoid was formed in the near-Earth region. The mechanism which effectively accelerate particles in the magnetosphere is not known at present, except for the magnetic field reconnection. Thus, the cause of these flows is an open question. The present event comprises a precursory breakup and a major breakup, which are inferred to be initiated in the near-Earth magnetotail and the mid-magnetotail, respectively. Our results and Paper I strongly argue that the major breakup preferably begin near the western edge of the auroral bulge created during the precursory breakup. How this aspect is related to the dynamics of the magnetotail is one of the important questions which should be addressed. Acknowledgments. We thank the Geotail staff at the Institute of Space and Astronautical Science, Japan, NASA/Goddard Space Flight Center, USA, and the National Oceanic and Atmospheric Administration, USA, for providing the magnetic field and plasma data obtained by Geotail, the solar wind data from Wind, and the magnetic field data from GOES 8 and 9, respectively. We are grateful to M. Brittnacher and Yeh-Kai Tung for generously providing Polar UVI images. This work was carried out by a joint research program of the Solar-Terrestrial Environment Laboratory, Nagoya University. REFERENCES Fairfield, D. H., and N. F. Ness, Configuration of the geomagnetic tail during substorms, J. Geophys. Res., 75, 7032-7047, 1970. Haland, S., N. 0stgaad, J. Bjordal, et al., Magnetospheric and ionospheric response to a substorm: Geotail HEP-LD and Polar PIXIE observations,/ Geophys. Res., 104, 28,459-28,474, 1999. Maezawa, K., Magnetotail boundary motion associated with geomagnetic substorms, J. Geophys. Res., 80, 3543-3548, 1975. Nakai, H., Y. Kamide, and M. Brittnacher, Magnetic field and plasma variations in the mid- magnetotail associated with pseudo and major breakups, in press, Proc. ICS-6, 2002. Nakai, H. and Y. Kamide, Substorm-associated large-scale magnetic field changes in the magnetotail: A prerequisite for "magnetotail deflation" events, in press, Annales Geophys., 2003. Tsyganenko, N. A., A magnetospheric magnetic field model with a warped tail current sheet, Planet. Space Sci., 37, 5, 1989.
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COMPARISON OF ENERGETIC ION COMPOSITION BETWEEN RING CURRENT AND PLASMA SHEET M. Nose', R. W. McEntire2, and S. P. Christen3 1
Data Analysis Center for Geomagnetism and Space Magnetism, Graduate School of Science, Kyoto University, Oiwake-cho, Kitashirakawa, Sakyo-ku, Kyoto 606-8502, Japan 2 Applied Physics Laboratory, Johns Hopkins University, 11100 Johns Hopkins Road, Laurel, MD 20723-6099, USA 3 Focused Analysis and Research, Columbia, MD 21044, USA
ABSTRACT We calculated the energy density of H+, He+, and O+ ions in the plasma sheet during the development of a magnetic storm, using energetic (9-210 keV) particle flux data obtained by the suprathermal ion composition spectrometer (STICS) sensor of the energetic particle and ion composition (EPIC) instrument on the Geotail spacecraft. We found that the energy density ratio of O+/H+ stayed at -0.1 before storms, but increased as storms developed, reaching 0.5-1.0 at a peak of storms. The energy density ratio of He+/H+ was rather constant at 0.01-0.02. These results are comparable with those in the outer ring current reported by the previous studies. This implies that the ions of ionospheric origin (O+ and He+) are transported to the ring current through the plasma sheet.
INTRODUCTION Energetic ion composition of the ring current changes drastically during magnetic storms. It has been reported that H+ ions are dominant during quiet times, while ions of ionospheric origin, O+ ions in particular, become an important constituent during storms. Table 1 shows results by previous studies about the ion composition changes in the outer ring current (L=5-7) at the afternoon to evening sectors (1400-2200 hour of local time). In these studies the AMPTE/CCE satellite has been used. This satellite was operative from September 1984 to January 1989, during which the solar activity was minimum and in the rising phase. Thus the satellite encountered moderate magnetic storms with the minimum Dst of-50--124 nT, except for a big storm analyzed by Hamilton et al. (1988). From Table 1 we found that the O+/H+ energy density ratio during quiet times is 0.01-0.03, while that during storms is 0.3-0.6. The He/H + energy density ratio changed from 0.01-0.02 during quiet times to 0.02-0.03 during storms. Table 2 displays the ion composition change in the outer ring current reported by previous studies using the CRRES satellite. The CRRES satellite has made observations near the solar maximum period, that is, from July 1990 to October 1991. Therefore the satellite made an observation during intense magnetic storms (Dstmin—300 nT). The results showed that the O+/H+ energy density ratio is 0.1-0.3 during quiet intervals and 0.5-3.0 during storms. In respect of the He+/H+ energy density ratio, it increases from ~0.05 during quiet times to 0.2-0.25 during storms. The ion composition change in the ring current might be a consequence of the ion composition change occurring in the plasma sheet, because plasma in the plasma sheet is thought to be convected to the ring current region by the ExB drift. The drift velocities of H+, He+, and O+ ions are the same, thus this ion transport process will conserve the ion composition. However, it might be also possible to think that ions of ionospheric origin (i.e., O+ and He+) are directly supplied from the ionosphere to the ring current without traveling through the plasma sheet, resulting in the enhancement of the O+/H+ and He+/H+ energy density ratios. To answer which of the above two
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scenarios is more plausible, we investigated the ion composition change in the plasma sheet during a magnetic storm, using the Geotail satellite. We made a comparison of the ion composition between the ring current and the plasma sheet on the basis of the following idea. If both are similar, we can suppose that the first scenario is plausible. On the other hand, if the ion composition in the plasma sheet is different from that in the ring current, the second scenario may be plausible. Table 1. Ion composition change in the outer ring current found by theAMPTE/CCE measurements O + /H +
Energy References
Gloeckleretal. (1985) Krimigisetal. (1985) Hamilton et al. (1988) Feldstein et al. (2000) Greenspan and Hamilton (2002)
&}
.ange (keV) 5-315 0.005-5000 30-310 1-310 1.5-300
Quiet 0.029 <0.01
Storm 0.34 0.38 0.61 <0.45 <0.3
He + /H +
Quiet 0.020 <0.01
Storm 0.023 0.028 0.022
Minimum
Dst (nT) -124 -124 -312 -105 -50-100
Table 2. Ion composition change in the outer ring current found by the CRRES measurements Energy O+/H+ He7H + Minimum References Range (keV) Quiet Storm Quiet Storm Dst(nT) Roederetal. (1996a) 40-430 0.25 2.0 0.04 0.2 -314 Roederetal. (1996b) 40-430 0.20 0.50 0.06 0.25 -223 Daglis(1997) 50-430 0.1-0.3 3.0 -80--300 Daglis et al. (2000) 50-430 1.0-2.0 -280
GEOTAIL/EPIC/STICS INSTRUMENT The Geotail satellite has been in the near-Earth orbit (~9 RE perigee x 30 RE apogee) and frequently surveyed the plasma sheet on the night side since 1995. Geotail carried the suprathermal ion composition spectrometer (ST1CS) sensor of the energetic particle and ion composition (EPIC) instrument. The EPIC/STICS sensor can measure the H+, O+, and He+ ion fluxes in the energy range of 9-210 keV/e that is scanned by eight logarithmically spaced energy steps. This sensor has a nearly full directional coverage (i.e., ~47i sr). Information on the EPIC/STICS instrument is compiled well by Williams et al. (1994). OBSERVATION Figure 1 a presents the SYM-H index for the period 1600 UT on April 15, 2000 through 1200 UT on April 16, 2000. The SYM-H index is essentially the same the Dst index, except that it is derived from six mid-latitude stations with a high-time resolution of 1 min (Iyemori et al., 1992). A magnetic storm started around 2200 UT on April 15 and reached at a peak of-94 nT at 1100 UT on April 16. During the time interval shown in Figure la, the Geotail satellite moved from (XcSM, Y GSM H-22.2, -1.3) RE to (-2.1, -10.5) RE, as indicated at the bottom of the figure. Figure lb shows time profiles of energy density of H+, O+, and He+ in the plasma sheet, which were calculated from the Geotail/EPIC/STICS data. The H+ energy density is displayed with a heavy line to avoid being confused with that of O+. There is a small data gap at 1910-1930 UT on April 15 because of Geotail's retreat from the plasma sheet. Before the magnetic storm started, the energy density of all ion species was rather constant, that is, ~1.0 keV/cm3 for H+, ~0.1 keV/cm3 for O+, and ~0.02 keV/cm3 for He+. However, the energy density increased as the magnetic storm developed after 2200 UT on April 15. The energy density around the storm maximum was ~10keV/cm 3 for H+, ~5 keV/cm3 for O+, and -0.2 keV/cm3 for He+. We noticed that the O+ energy density increased more strongly than the H+ energy density. This strong enhancement of the O+ energy density can be seen clearly in Figure lc, in which the O+/H+ and He+/H+ energy density ratios are traced. The O+/H+ energy density ratio increased from -0.1 before the storm to 0.5-1.0 around the end of the storm main phase. However, the He/H energy density was nearly constant at 0.01-0.02 in this event.
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2000/04/15-04/16: SYM-H [1-min Dst]
universal Time 2000/04/15-04/16: Geotail EPIC/STICS {H+, O + , He+)
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Universal Time 2 0 0 0 / 0 4 / 1 5 - 0 4 / 1 6 : G e o t a i l EPIC/STICS
(O+/H+,
He+/H+)
16 17 18 19 20 21 22 23 24 01 02 03 04 05 06 07 08 09 10 11 12 Universal Time -16.9 -8.1 -2.1 -19.9 -6.7 -4.3 -10.4 -10.5 1.4 -0.4 -0.1 -0.6
Fig. 1. (a) The SYM-H index for the period 1600 UT on April 15, 2000 to 1200 UT on April 16, 2000. (b) The energy density of H + , O + , and He+ ions calculated from the Geotail/EPIC/STICS instrument. (c) The O + /H + and He7H + energy density ratios.
DISCUSSION Plasma Sheet Ion Composition In the April 15-16, 2000, event, we found that the O+/H+ energy density ratio changed from ~0.1 to 0.5-1.0 during the storm development, while the He+/H+ stayed at 0.01-0.02 (Figure lc). There are a limited number of previous studies that examined ion composition change in the plasma sheet during magnetic storms. Using the 1SEE-1 data at a geocentric distance o f - 2 2 RE, Ipavich et al. (1984) reported that the O+/H+ flux ratio at -130 keV increased from -0.03 to -0.35. Nose et al. (2001, 2003) examined ion composition change in the 1: Ring Current (AMPTE) 10.00 near-Earth and mid-tail plasma sheet (X=-9~-18 RE) Dst=-50~-124nT : 2: Ring Current (CRRES) Dst—3 00nT during six magnetic storms with Dstmjn= 3: Plasma Sheet (Geotail) Dst=-100—150nT T 4: Plasma Sheet (This Study) ' -100~-150nT by using the Geotail/EPIC/STICS Dst=-94nT 1.00 + + data. They found that the O /H energy density ratio 2 2 was 0.05-0.1 during quiet times but reached 0.3-1.0 I + + 3 during storms. The He /H energy density ratio 0.10 during quiet times was 0.01-0.02, while that during 1 134 f L storms was 0.02-0.1.
H1
0.01
in
0+ 'H+
Comparison of Ion Composition Between Ring Current and Plasma Sheet In Figure 2 we summarized the change of the
\ If ; i Quiet
Storm
i
;
He+/H+ Quiet
Storm
Fig. 2. Summary of the ion composition change in the ring current and the plasma sheet.
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ion composition in the ring current and plasma sheet. Vertical lines labeled " 1 " and "2" indicate the energy density ratios in the outer ring current that were reported by previous studies using the AMPTE/CCE and CRRES satellites, respectively (see Introduction, Tables 1 and 2). Lines 1 are for storms with Dstmill=-50~-]24 nT and lines 2 are for storms with Ds^^-SOO nT. Vertical lines labeled " 3 " correspond to results by Nose et al. (2001, 2003), who examined ion composition in the plasma sheet during magnetic storms with Dstmin=-100~-150nT. The result for the April 15-16, 2000, event (Dstmjn=-94 nT) is also displayed by vertical lines labeled "4". From Figure 2 we found that the results in the plasma sheet (lines 3 and 4) are comparable with those in the outer ring current reported by the previous studies using AMPTE/CCE (lines 1). The results by CRRES (lines 2) are larger than those in the plasma sheet. This would be because the CRRES satellite experienced more intense magnetic storms. Therefore we conclude that ions of ionospheric origin are possibly convected from the plasma sheet to the ring current. In future study we need to investigate how ionospheric ions are accelerated to the ring current energy (20-200 keV). It might be feasible that ions undergo mass-dependent acceleration by the dawn-to-dusk electric field in the current sheet, as proposed by Nose et al. (2001). Particle tracing in a realistic magnetotail model will be conducted to test this acceleration process. ACKNOWLEDGMENTS We thank D. J. Williams and S. R. Nylund for their help in analysis of the Geotail/EPIC/STICS data. The SYM-H index was provided by T. Iyemori at WDC for Geomagnetism, Kyoto. Thanks are due to S. Ohtani, K. Takahashi, and A. T. Y. Lui for their helpful comments. This work was partly supported by the Atmospheric Science Division of the National Science Foundation (grant ATM-0000255) to The Johns Hopkins University Applied Physics Laboratory and the Sasagawa Scientific Research Grant from The Japan Science Society.
REFERENCES Daglis, I. A., The role of magnetosphere-ionosphere coupling in magnetic storm dynamics, in Magnetic Storms, Geophys. Monogr. Ser., vol. 98, edited by B. T. Tsurutani et al., pp. 107-116, AGU, Washington, D. C , 1997. Daglis, I. A., Y. Kamide, C. Mouikis, et al., "Fine structure" of the storm-substorm relationship: Ion injections during Dst decrease, Adv. Space Res., 25(12), 2369-2372, 2000. Feldstein, Y. I., L. A. Dremukhina, U. Mall, et al., On the two-phase decay of the Dst variation, Geophys. Res. Lett., 27,2813-2816,2000. Gloecker, G., B. Wilken, W. Stiidemann, et al., First composition measurement of the bulk of the storm-time ring current (1 to 300 keV/e) with AMPTE-CCE, Geophys. Res. Lett., 12, 325-328, 1985. Greenspan, M. E., and D. C. Hamilton, Relative contributions of H+ and O+ to the ring current energy near magnetic storm maximum, J. Geophys. Res., 107(A4), doi: 10.1029/2001JA000155, 2002. Hamilton, D. C , G Gloeckler, F. M. Ipavich, et al., Ring current development during the great geomagnetic storm of February 1986,7. Geophys. Res., 93, 14343-14355, 1988. Ipavich, F. M., A. B. Galvin, G. Gloeckler, et al., Energetic (>100 keV) O ions in the plasma sheet, Geophys. Res. Lett.,11, 504-507, 1984. Iyemori, T., T. Araki, T. Kamei, et al., Mid-latitude geomagnetic indices ASY and SYM (Provisional) No.l 1989, Data Anal. Center for Geomagn. and Space Mang., Kyoto Univ., Kyoto, Japan, 1992. Klecker, B., E. Mobius, D. Hovestadt, et al., Discovery of energtic molecular ions ( N O and O2+) in the storm time ring current, Geophys. Res. Lett., 13, 632-635,1986. Krimigis, S. M., G Gloeckler, R. W. McEntire, et al., Magnetic storm of September 4, 1984: A syntesis of ring current spectra and energy densities measured with AMPTE/CCE, Geophys. Res. Lett., 12, 329-332, 1985. Nose, M., S. Ohtani, K. Takahashi, et al., Ion composition of the near-Earth plasma sheet in storm and quiet intervals: Geotail/EPIC measurements, J. Geophys. Res., 106, 8391-8403, 2001. Nose, M., R. W. McEntire, and S. P. Christon, Change of the plasma sheet ion composition during magnetic storm development observed by the Geotail spacecraft, J. Geophys. Res., 108, doi:10.1029/2002JA009660, in press, 2003. Roeder, J. L., J. F. Fennell, M. W. Chen, et al., CRRES observations of the composition of the ring-current ion populations, Adv. Space Res., 17(10), 17-24, 1996a. Roeder, J. L., J. F. Fennell, M. W. Chen, et al., CRRES observations of stormtime ring current ion composition, in -201-
Workshop on the Earth's trapped particle environment, AIP Conference Proceedings, vol. 383, edited by G D. Reeves, pp. 131-135, American Institute of Physics, New York, 1996b. Williams, D. J., R. W. McBntire, C. Schlemm 11, et al., Geotail energetic particles and ion composition instrument, J. Geomagn. Geoelectr., 46, 39-57, 1994. E-mail address of M. Nose
[email protected]
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SECTION 4: Microscopic Processes in Space Plasmas
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GEOTAIL, POLAR, AND WIND OBSERVATIONS OF AURORAL KILOMETRIC RADIATION Roger R. Anderson2'1, Hiroshi Matsumoto2, Kozo Hashimoto2, Hirotsugu Kojima2, Yasumasa Kasaba3, Michael L. Kaiser , Jean-Louis Bougeret , Jean-Louis Steinberg , and Gordon Rostoker 1
Department of Physics and Astronomy, The University of Iowa, Iowa City, IA 52242-1479, USA Radio Science Center for Space and Atmosphere, Kyoto University, Gokanosho, Uji, Kyoto 611-0011, Japan 3 The Institute of Space and Astronautical Science, 3-1-1, Yoshinodai, Sagamihara, Kanagawa 229, Japan 4 NASA/Goddard Space Flight Center, Laboratory for Extraterrestrial Physics, Code 695, Greenbelt, MD, 20771, USA 5 Observatoire de Paris, LESIA/CNRS, 5, Place Jules Janssen, 92195 Meudon, France ^Department of Physics, University of Alberta, Edmonton, Alberta, Canada T6G 2JI 2
ABSTRACT Aurora] kilometric radiation (AKR) is the plasma wave/radio phenomenon most clearly associated with substorms and increased geomagnetic activity. The GEOTAIL and POLAR Plasma Wave Instruments (PWI) both included sweep frequency receivers that had an upper frequency limit of 800 kHz and the WIND WAVES Thermal Noise Receiver (TNR) and Radio Receiver Band 1 (RAD1) went to 256 kHz and 1024 kHz, respectively. We have thus been able to observe the majority of the AKR spectrum in better detail than with earlier instrumentation and many important new discoveries have been made. Terrestrial low frequency (LF) bursts are a part of AKR observed during strong substorms. Although a limited portion of the LF burst spectrum is often detected on the dayside of the Earth and in the upstream solar wind, the complete spectrum is most frequently detected by spacecraft in the nightside magnetosphere or geomagnetic tail. Frequently these observations show that the LF bursts have a tapered tail centered on the present or recent past solar wind plasma frequency. We have found that on the dayside and in the upstream solar wind the high frequency AKR is detected during LF burst events only if the path from the AKR source is not blocked by the earth or dense plasmasphere. POLAR observations from high over the AKR source region show that the AKR increases in intensity and its lower frequency limits decrease when LF bursts are observed indicating that the AKR source region is expanding to higher altitudes. Frequently the upper frequency limit also increases indicating that the source region is then also expanding to lower altitudes. Data from both satellite and ground-based experiments show that the LF bursts are well correlated with expansive phase onsets and occur during very geomagnetically-disturbed periods. High resolution (in both time and frequency) data from the POLAR Wide Band Receiver have yielded exciting data on the fine structure of AKR as well as details on the structure of LF bursts. INTRODUCTION Plasma wave and radio measurements from the GEOTAIL, POLAR, and WIND spacecraft which were a part of the International Solar Terrestrial Physics/Global Geospace Science (ISTP/GGS) program (Acuna et al., 1995) have provided both in situ and remote observations of numerous plasma wave phenomena related to geomagnetic storms and substorms. Observations of auroral kilometric radiation (AKR) (Gurnett, 1974; Voots et al., 1977; Kaiser and Alexander, 1977b), the phenomenon most clearly associated with substorms and increased geomagnetic activity, provide remote or in situ (depending on the orbit) indicators of the timing, dynamics, and strengths of geomagnetic storms and substorms as well as characteristics of the source region and the generation mechanisms. The ISTP/GGS spacecraft were well instrumented to advance our understanding of AKR and related phenomena (Anderson et al., 1997, 1998, 2001). The GEOTAIL Plasma -205-
Wave Instrument (PWI) (Matsumoto et al., 1994) included a sweep frequency analyzer (SFA) that went up to 800 kHz with a sweep period of 8 seconds and a multichannel analyzer (MCA) that went up to 311 kHz and could produce a spectrum every 1/4 second. The POLAR PWI (Gurnett et al., 1995) included sweep frequency receivers (SFR) that also went up to 800 kHz with a sweep period of 2.4 seconds and an MCA that went up to 311 kHz and produced a spectrum every 1.3 seconds. The POLAR PWI also included a Wide Band Receiver (WBR) that provided both high-time resolution and high-frequency resolution data (up to 249,000 samples per second) with up to a 90 kHz bandwidth and a translatable frequency offset covering much of the AKR spectrum. The radio science experiment WAVES on WIND (Bougeret et al., 1995) included a Thermal Noise Receiver (TNR) that covered 4 kHz to 256 kHz and a Radio Receiver Band 1 (RAD1) that went to 1024 kHz but both had somewhat poorer time resolution than GEOTAIL or POLAR. The initiation or intensification of AKR occurs coincident with substorm onset. Fairfield et al., (1998,1999) showed that AKR onsets and intensifications observed by WIND WAVES, GEOTAIL PWI, and POLAR PWI were coincident and closely associated with high-velocity earthward flow bursts in the inner magnetotail identified with substorm onset. In order to calculate the sizes of plasmoids and the locations of the nearEarth X-line, Murata et al. (1995) used GEOTAIL PWI AKR observations to identify substorm onset (and the release or initiation of plasmoid flow) and in situ measurements of the magnetic fields from the GEOTAIL Magnetic Field Experiment (MGF) (Kokubun et al., 1994) and magnetic noise bursts from the MCA data to determine the timing of the passage of a plasmoid/flux rope over the spacecraft. Liou et al. (1999) in an investigation of relative timing in substorm signatures found that the start of enhanced AKR observed by POLAR when the spacecraft was in the midnight sector was coincident with a sudden auroral brightening indicative of an auroral breakup. These and other studies have shown that for spacecraft with an adequate view of the auroral zone, the initiation or intensification of AKR can be a very accurate identifier of substorm onset. A 38-month data set of GEOTAIL plasma wave observations of AKR from 100 kHz to 600 kHz was used by Kasaba et al. (1997) to examine the dependence of the angular distribution of AKR. While Green et al. (1977) and Green and Gallagher (1985) had found using the IMP-6 and Hawkeye MCA data (with a 178 kHz upper frequency limit) that the AKR illumination pattern broadens with increasing frequency from 56.2 kHz to 178 kHz, Kasaba et al. (1997) found using the GEOTAIL SFA data (with a 800 kHz upper frequency limit) that the AKR illumination pattern then becomes narrower above 300 kHz. Such differences are basically explained by propagation. The GEOTAIL study also found that the illumination region of AKR extends duskward as geomagnetic conditions become more disturbed especially for the low frequency range which suggests a duskward extension of the AKR source. Kasaba et al. (1997) speculated that the lack of such a feature in the high frequency range could be caused by insufficient density depression in the duskside auroral plasma cavity especially at lower altitudes. Hashimoto et al. (1998) compared AKR simultaneously observed by GEOTAIL PWI and WIND WAVES and found using ray tracing that most observations were explained by an AKR source around 21-22 MLT (Magnetic Local Time) but during substorms the source extended around to 19 MLT. Several studies using the POLAR and/or GEOTAIL PWI measurements in comparison with data from the POLAR imaging experiments have shown clear associations between electron precipitation and AKR. Plasma wave and bremsstrahlung x-ray data from the POLAR PWI and PIXIE (Imhof et al., 1995) experiments were used by Imhof et al. (1998) to find a 0.51 correlation coefficient over a six hour local time range in the pre-midnight sector for a satellite pass that had several short term enhancements in the intensities of both AKR waves from 60 kHz to 800 kHz and 2 to 12 keV x-ray emissions. Other POLAR PWI and PIXIE (Imhof et al., 1999; 2000) and GEOTAIL PWI and PIXIE (Imhof et al., 2001) correlative studies found that the cross-correlation coefficient of auroral x-rays and AKR emissions was enhanced over an MLT interval of six hours or less and the maximum occurred for x-rays emitted slightly before local midnight. In another GEOTAIL PWI and PIXIE correlative study of the dependence of AKR production on the intensity and energy spectra of auroral bremsstrahlung x-rays, Imhof et al. (2003) found that the x-rays with higher e-folding energy were correlated with higher AKR frequency (implying lower altitude of generation). This was interpreted as being due to the increased energy of the primary electrons resulting from acceleration through an increasing potential difference at lower altitude. Imhof et al. (2003) also found that higher x-ray fluxes were associated with lower AKR cutoff frequencies. The primary generation mechanism presently -206-
considered for AKR, the electron cyclotron maser instability (Wu and Lee, 1979), requires a small plasma frequency to cyclotron frequency ratio. AKR is generated near the local electron cyclotron frequency Fee such that high frequency AKR is generated at a lower altitude than the low frequency AKR. High time and frequency resolution measurements by the FAST satellite have confirmed that the AKR source is in a density depleted cavity and the AKR emissions are very close to and sometimes slightly below the cold plasma Fee down to the relativistic Fee (Ergun et al., 1998). The frequency of AKR thus identifies the location along a magnetic field line where it is generated. For example, on a 71 degree invariant latitude (L=9.4) field line typical for an AKR source region, Fee equals 500 kHz at a geocentric radial distance of 1.47 Re, 400 kHz at 1.58 Re, 300 kHz at 1.74 Re, 200 kHz at 1.98 Re, 100 kHz at 2.47 Re, 50 kHz at 3.09 Re, 30 kHz at 3.60 Re, and 15 kHz at 4.52 Re. Subtracting 1 Re yields the altitude. In the duskside plasmasphere, the electron density is enhanced so that the density in the auroral plasma cavity should be hard to decrease enough to satisfy the condition for the electron cyclotron maser instability. Therefore, generation of high-frequency AKR at lower altitudes is expected to be blocked on the duskside hemisphere. Kasaba et al. (1997) found that at both 200 kHz and 500 kHz the frequency of occurrence of AKR was positively correlated with the Kp index and that this was more evident for the 200 kHz data. The latter agreed with an earlier finding of Kaiser and Alexander (1977a) that the frequency of the peak flux of AKR was inversely correlated with the AE index. These results indicate that the AKR source region moves up in altitude under disturbed geomagnetic conditions. Kasaba et al. (1997) also found that the illumination region increased in size during active times with an equatorward extension common at both frequencies which is believed to be due to the equatorward shift of the auroral plasma cavity in the disturbed phase expected from the inward motion of the plasmapause just after the onset of substorms. For the duskside region and at higher magnetic latitudes, the illumination pattern became larger for 200 kHz than for 500 kHz as the geomagnetic activity increased. This was believed to be due to the influence of the evening plasmaspheric bulge on the AKR propagation especially at the lower altitudes. Another important new result from Kasaba et al. (1997) was that AKR is more active on the winter hemisphere especially for the high frequency range. Possible reasons include asymmetry of the population of precipitating electrons on the auroral field lines and insufficient density depression in the auroral plasma cavity on the summer hemisphere especially at lower altitudes which are most sensitive to ionospheric outflow. Terrestrial low frequency (LF) bursts are a part of AKR often observed during geomagnetic storms and strong substorms (Steinberg et al., 1988, 1990, 1998; Kaiser et al., 1996; Anderson et al., 1997, 1998, 2001). Frequently the LF bursts occur during a period of the lower cutoff frequency of AKR continually decreasing. Alexander and Kaiser (1976) first reported that the AKR lower cutoff frequency tended to move to lower frequencies during geomagnetically disturbed periods. Frequently the upper frequency limit increases indicating that the source region is then also expanding to lower altitudes. The changing upper and lower frequency limits observed during substorms can be used to study the dynamics of the plasma in the AKR source region. The POLAR WBR has yielded exciting data on the fine structure of AKR as well as details on the structure of LF bursts. Data from the plasma wave instruments on GEOTAIL and POLAR and the radio science experiment on WIND will be used here to study the characteristics and relationship of AKR and LF bursts. These three spacecraft and the CANOPUS (Rostoker et al., 1995) ground magnetometer network which provided data for this study were all parts of the ISTP/GGS program. We will concentrate first on quasi-periodic AKR low frequency enhancements well correlated with ground magnetometer observations and then on multi- spacecraft observations of the enhanced AKR resulting in LF burst observations in the magnetosphere and in the solar wind. We will examine events that illustrate the nature of AKR and LF bursts and their relationships to other measures of geomagnetic activity and we will compare simultaneous observations from different regions of space. OBSERVATIONS Quasi-periodic Intensifications An interesting characteristic of many of the LF bursts that we have observed is that they occur as a part of a series of quasi-periodic AKR bursts whose lower cutoff frequencies progressively decrease. After the LF burst occurs, a series of quasi-periodic AKR bursts whose lower cutoff frequencies increase progressively to higher frequencies often follows. An example of this is shown in Figure 1 which displays the GEOTAIL PWI -207-
Sweep Frequency Analyzer (SFA) data on February 23, 1994, from 03:00 UT to 05:00 UT. GEOTAIL was about 85 Re behind and towards the dusk side of the Earth in the solar wind just outside the early evening magnetosheath at GSE (X, Y, Z) = (-47.7 Re, 69.5 Re, -4.6 Re) and GSM (X, Y, Z) = (-47.7 Re, 64.1 Re, 27.4 Re). WIND and POLAR had not yet been launched. Five distinct events occur between 03:30 UT and 04:45 UT in which the AKR intensifies and its lower frequency cutoff progressively decreases (until 04:00 UT) and then increases (after 04:05 UT). The decreasing lower cutoff frequency means that the conditions required for the generation of AKR are moving to higher altitudes. These conditions include the local Fp in the source region being significantly less than the local Fee and that there be a source of free energy from a positive slope in the perpendicular electron velocity distribution.
Fig. 1. The GEOTAIL PWI electric field SFA data for 03:00 UT to 05:00 UT on February 23, 1994. The data shown are on linear frequency scales from 12.5 kHz to 100 kHz and from 100 kHz to 800 kHz. Five distinct quasi-perioidic enhancements in the AKR are evident between 03:30 UT and 04:45 UT that include the lower frequency cutoff significantly decreasing until 04:00 UT and then increasing after 04:05 UT. These enhancements are coincident with negative bay onsets and Pi2 oscillations shown in Figure 2. The diffuse tail of a terrestrial LF burst is evident beginning near 30 kHz at 04:00 UT which reaches 15 kHz before 04:05 UT. The nearly-constant-frequency emission line near 15 kHz is at the local electron plasma frequency Fp. The frequent enhancements in intensity are Langmuir waves excited when the region is magnetically connected to the Earth's bow shock (Kasaba et al., 2000). The weak narrow emission line just above 30 kHz is the 2Fp line which is generated near the Earth's bow shock on or just downstream of magnetic field lines tangent to the Earth's bow shock (Reiner et al., 1996) The plasma frequency in the near-Earth magnetosheath is 2Fp at the bow shock and gradually decreases downstream until it equals the local solar wind Fp. A LF burst is evident beginning near 30 kHz at 04:00 UT which reaches 15 kHz before 04:05 UT. The LF burst after falling below the 2Fp line approaches the Fp line with increasingly longer delay times. This long diffuse tail is an identifiable characteristic of the LF bursts. A very weak LF burst also is detected between 30 kHz and 15 kHz beginning at 03:48 UT and lasting for about five minutes.
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Table 1. CANOPUS Magnetometer Sites
Location Rankin Inlet Eskimo Point Fort Churchill Gillam Island Lake Pinawa Rabbit Lake Contwoyto Lake Fort Smith Fort McMurray Fort Simpson Dawson
Acronym RANK ESKI FCHU GILL ISLL PINA RABB CONT FSMI MCMU FSIM DAWS
Geodedic Lat. 62.8 61.1 58.8 56.4 53.9 50.2 58.2 65.8 60.0 56.7 61.7 64.1
Long. 267.9 266.0 265.9 265.4 265.3 264.0 256.3 248.8 248.1 248.8 238.8 220.9
Pace Lat. 73.7 71.9 69.7 67.4 64.9 61.2 67.8 73.4 67.9 64.8 67.6 65.9
Long. -29.0 -31.8 -30.8 -30.9 -30.3 -31.6 -45.0 -61.2 -57.3 -54.4 -69.9 -90.1
L
12.4 10.2 8.2 6.7 5.5 4.3 6.9
12.4 7.1 5.5 6.8 5.9
Inv. Lat. 73.5 71.7 69.6 67.3 64.8 61.2 67.6 73.5 68.0 64.8 67.4 65.7
The geomagnetic conditions were moderately active with Kp = 3+ and DST = -57 nT. Slightly more than one day earlier a strong storm produced Kp = 8- and DST = -144 nT. The north-south (X) components of the CANOPUS array ground magnetometer data for 0300 UT to 0500 UT on February 23, 1994, are shown in Figure 2. The locations of the twelve stations for which there are data are listed in Table 1.
Fig. 2. The CANOPUS array magnetometer north-south (X) component data for 03:00 UT to 05:00 UT on February 23, 1994. Panel A, the left panel, is unfiltered and Panel B, the right panel, is high-pass filtered with fc = 7 mHz which highlights Pi2 oscillations associated with expansive phase onsets. Full scale in Panel A is 650 nT and full scale in Panel B is 180 nT. The quasi-periodic negative bay onsets and Pi2 oscillation enhancements correspond well with the five AKR LF enhancements in Figure 1.
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Strong negative bays of the order of several hundred nT and strong Pi2 oscillations are evident especially in the Fort Churchill and Eskimo Point data for the first three distinct events and also in the higher latitude Rankin Inlet data for the final two. Both are indicative of expansive phase onsets. The LF burst beginning at 04:00 UT and the AKR lower cutoff frequency reaching its lowest value are coincident with the strongest Pi2 oscillation observed for the two-hour period. The solar wind speed measured by the GEOTAIL Comprehensive Plasma Instrument (CPI) (Frank et al., 1994) was high at about 700 km/sec (K. L. Ackerson, Private Communication). Desch et al. (1996) and Desch (1997) found that the occurrence of LF bursts were well correlated with the solar wind speed and somewhat less correlated with the azimuthal direction of the IMF. The LF bursts tended to occur in the sectors when the IMF was pointed towards the sun. They concluded that the LF bursts were a signature of a large-scale process of magnetospheric energy dissipation. They suggested a viscous-like interaction between the solar wind and the magnetosphere. For this event, the solar wind speed was higher than normal but the measurements were made in an away sector. AKR generated in the auroral region above twice the solar wind plasma frequency can propagate directly through the magnetosheath into the solar wind. The diffuse tails of the LF bursts are believed to be due to the lowest frequency portion of the AKR waves scattering far down the tail where the plasma frequency of the magnetosheath reaches the solar wind plasma frequency. Nearer the Earth these waves are reflected by the high density encountered at the magnetopause. As an electromagnetic wave approaches a region where the plasma frequency is near the wave frequency, the group velocity decreases and the delay time increases. The closer the frequency of the wave is to the plasma frequency, the slower it travels and the more likely it is to be scattered. As the magnetosheath density decreases further down the tail, it approaches the solar wind density. The scattering can thus carry the waves out into the solar wind. This results in a very large apparent source size for the radiation which would have no spin modulation. Multi-Spacecraft Observations On January 28, 1997, the WIND spacecraft was 174 Re upstream of the Earth near the Earth-Sun line at GSE X = 172 Re, GSE Y = -20 Re, and GSE Z = -16 Re.
Fig. 3. The WIND/WAVES TNR data for 13:00 UT to 15:00 UT on January 28, 1997 when WIND was 174 Re upstream of the Earth. The data shown are on a logarithmic frequency scale from 4 kHz to 256 kHz. The color-coded intensity scale covers 10 dB relative to the receiver background. The LF burst beginning at 13:42 UT has a distinctive diffuse tail. The solar wind speed (from the WIND comprehensive plasma experiment SWE (Ogilvie et al., 1995)) was high at about 650 km/s such that the local solar wind parameters would reach Earth in about 28 minutes.
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Fig. 4. The GEOTAIL PWI electric field SFA data for 13:00 UT to 15:00 UT on January 28, 1997 when GEOTAIL was about 30 Re down the tail. The data are displayed in five linear frequency bands: 25 Hz to 200 Hz, 200 Hz to 1600 Hz, 1.6 kHz to 12.5 kHz, 12.5 kHz to 100 kHz, and 100 kHz to 800 kHz. The intensity is color coded over a 70 dB dynamic range. The black and white line near the bottom of the spectrogram indicates the local electron cyclotron frequency determined from the GEOTAIL MGF experiment. Just before 13:40 UT the lower cutoff frequency of the AKR began to fall and reached below 20 kHz at 13:44 UT. A tapered diffuse tail centered very near 20 kHz is evident for about ten minutes after 13:44 UT.
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Two LF bursts that were detected at WIND in the second half of the day will be examined here to better understand their relationship to the AKR observed. One began at 13:42 UT and the second began around 19:05 UT. Geomagnetic activity was moderately active with with Kp from 12 UT to 21 UT being 4, 4, and 4+. Figure 3 shows the WIND/WAVES TNR spectrogram for the 4 kHz to 256 kHz logarithmic frequency range from 13:00 UT to 15:00 UT. The intensity is color coded over a 10 dB range relative to the receiver background. The emission line near 20 kHz at the beginning of the plot is at the local electron plasma frequency Fp. The narrow emission line near 40 kHz at the beginning of the plot and which decreases to about 35 kHz by the end of the plot is the 2Fp line which is generated near the Earth's bow shock. The LF burst began between 40 and 70 kHz in frequency at 13:42 UT and fell below the 2Fp line at 13:44 UT and then approached the Fp line with increasingly longer delay times. The other emissions above 2Fp are AKR. The GEOTAIL and POLAR PWI data for the same two-hour period are shown in Figures 4 and 5, respectively. GEOTAIL was about 30 Re down the tail at GSE X = -28.3 Re, GSE Y = 10.2 Re, and GSE Z = -3.2 Re and was moving toward the center of the tail. The 2 kHz lower cutoff of the continuum radiation (which here extends from 2kHz to 12 kHz) is at the local plasma frequency and indicates low number densities around 0.04 e/cc. Several times throughout this period the lower cutoff frequency of the AKR fell below 100 kHz.
Fig. 5. The POLAR PWI SFR Ez data for 13:00 UT to 15:00 UT on January 28, 1997 when POLAR was in the pre-dawn sector ( 4.5 MLT) approaching the plasmasphere. The data shown are on a logarithmic frequency scale from 10 kHz to 800 kHz. The intensity is color coded over a 50 dB dynamic range. The white line which crosses 10 kHz just before 14:00 UT is the electron cyclotron frequency determined from the ambient magnetic field measurements of the POLAR MFI experiment (Russell et al., 1995). The LF burst with the tapered tail centered on 20 Khz is clearly evident beginning at 13:44 UT.
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In the period from 13:30 UT to 15:00 UT we see that when the AKR emissions appeared below 100 kHz that the AKR above 100 kHz was enhanced and that strong extremely low frequency (ELF) emissions (from 25 Hz to 1-3 kHz) occurred simultaneously. These elf emissions accompany the bursty bulk flows that have been identified with substorms and AKR intensifications (Fairfield et al., 1998,1999). However, only one of these episodes produced a clear LF burst that was also detected by WIND. Just before 13:40 UT the lower cutoff frequency of the AKR began to fall and reached below 20 kHz at 13:44 UT. A tapered diffuse tail centered very near 20 kHz is evident for about ten minutes after 13:44 UT. Only when the AKR reaches a frequency significantly below the maximum magnetosheath frequency (twice the solar wind plasma frequency at the bow shock) is a LF burst observed. Figure 5 displays the POLAR PWI SFR Ez (the electric field measured by the antenna along the spacecraft spin-axis) data logarithmically from 10 kHz to 800 kHz. The LF burst with the tapered tail centered on 20 Khz is clearly evident beginning at 13:44 UT. Note that the strong AKR observed on GEOTAIL after 14:40 UT is not present in the POLAR data. It has been refracted away by the dense plasmasphere. The WIND/WAVES TNR data from 18:00 UT to 20:00 UT over a 12 dB range relative to the receiver background are shown in Figure 6. The LF burst begins above 256 kHz and falls below the 2Fp line at 19:05 UT and approaches the Fp line with increasingly longer delay times. Fp and 2Fp have fallen slightly from five hours earlier to 18 kHz and 36 kHz, respectively. The GEOTAIL PWI SFA data over a 70 dB dynamic range for the same two-hour time period are shown in Figure 7. A LF burst with a tapered tail centered on 20 Khz is clearly evident from about 19:05 UT to 19:15 UT. Note that the strong ELF waves believed to be associated with bursty bulk flows occur from 18:50 UT to 19:13 UT simultaneously with the lower cutoff frequency of AKR being very low. The POLAR PWI SFR Ez data are shown in Figure 8. Polar near dusk ( 18 MLT) has exited the plasmasphere and is headed for its north pole apogee. Diffuse emissions down to about 15 kHz are present from 19:05 UT to 19:15 UT. Very intense AKR reaches as low as 32 kHz at about 19:13 UT.
Fig. 6. The WIND/WAVES TNR data for 18:00 UT to 20:00 UT on January 28, 1997. The color-coded intensity scale covers 12 dB relative to the receiver background. The LF burst beginning at 19:05 UT has a distinctive diffuse tail.
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Fig. 7. The GEOTAIL PWI SFA data for 18:00 UT to 20:00 UT on January 28, 1997. A LF burst with a tapered tail centered on 20 Khz is clearly evident from about 19:05 UT to 19:15 UT. Strong ELF waves associated with bursty bulk flows occur from 18:50 UT to 19:13 UT simultaneously with the lower cutoff frequency of AKR being very low. POLAR PWI WBR data in the 0-90 kHz mode were available for the 19:05 UT LF burst and showed that there was a distinct difference in the higher frequency AKR in this case as compared to the lower frequency AKR. The higher frequency AKR was more intense and the discrete structure was spin modulated and at least initially was predominantly rising in frequency. In the lower frequency portion the discrete emissions
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were dominated by falling frequency stripe-like features that were not always well spin-modulated. Near the end of the WBR data after about 19:11 UT the lower cutoff frequency of the more intense higher frequency AKR dropped to about 32 kHz and the weaker lower frequency portion extended below 20 kHz. What we have seen here is that POLAR does observe AKR that is detected as a LF burst by GEOTAIL and WIND but that the lowest frequency portion is less intense and has different discrete structure. One possibility is that a more intense lower frequency portion is generated near or earlier than local midnight that is more visible to GEOTAIL situated close to local midnight.
Fig. 8. The POLAR PWI SFR Ez data for 18:00 UT to 20:00 UT on January 28, 1997. POLAR was near dusk ( 18 MLT), had exited the plasmasphere, and was headed for its north pole apogee. Diffuse emissions down to about 15 kHz are present from 19:05 UT to 19:15 UT. Very intense AKR reaches as low as 32 kHz at about 19:13 UT. DISCUSSION AND SUMMARY GEOTAIL and POLAR PWI measurements along with those from the WIND WAVES experiment, the CANOPUS ground optical data and magnetograms, the IMAGE Network magnetograms, and POLAR auroral imaging have provided new information on AKR and LF bursts and their relationships to the magnetospheric plasma dynamics. AKR measurements when a satellite can view the nightside hemisphere provide excellent indications of substorm onsets. The absence of higher frequency AKR in some upstream observations of LF bursts can be attributed to propagation blockage by the Earth and dense plasmasphere of the portion of the AKR generated at the lowest altitudes on the night side. The ideal situation for substorm onset identification would include the simultaneous availability of multiple spacecraft observations of AKR covering all local times, worldwide ground all sky camera, magnetometer, and meridian scanning
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photometer measurements, and space-based auroral imaging. The combination of the high time-resolution and remote sensing capabilities provided by the plasma wave measurements make them very important for studying the triggering, near triggering, and burstiness of substorms and related geomagnetic disturbances. We have found that LF bursts are the lower frequency portion of AKR and are produced simultaneously with intense isolated substorms. Spacecraft observations often show that the AKR increases in intensity and its lower frequency limits decrease when LF bursts are observed indicating that the AKR source region is expanding to higher altitudes. Since AKR is generated near the local electron cyclotron frequency, the frequency of AKR identifies the location along a magnetic field line where it is generated and the lower the frequency, the higher the altitude. This relationship allows us to investigate some source characteristics such as the average speed for the upward movement of the AKR source region. Frequently the upper frequency limit also increases indicating that the source region is then also expanding to lower altitudes. In the future we will examine source speeds for all the LF events with measurable upper and lower cutoff frequencies. We will look for geophysical parameters that might correlate with the speeds. In a limited number of cases that we have been able to examine in detail, CANOPUS ground magnetometer and meridian scanning photometer data show that during the LF burst events the expansive phase onsets start at unusually low latitudes and move poleward and westward. The data also show that the LF bursts occur when the expansive phase onset signatures are most intense. Many of the LF bursts occurred related to CME events observed by SOHO which were identified by the NOAA SEC as being highly geoeffective. Magnetometer data from geosynchronous satellites usually show increased magnetic field dipolarization and the presence of field-aligned currents during LF burst events. Large injections of protons and electrons have also been detected by the GOES and LANL geosynchronous satellites during LF burst events. AKRproduced low frequency bursts which occur during the more intense geomagnetic disturbances thus provide a space weather marker of the geoeffectiveness of the events. We will continue to investigate the global activity that is associated with and leads to LF bursts. The already productive correlative studies with the POLAR x-ray and optical imaging are continuing and will now concentrate on LF burst events. We will also continue the investigation of the dimensions and shape of the geomagnetic tail. Our analysis using our tapered-tail LF burst observations and solar wind speed and density measurements from available spacecraft have yielded tail lengths fron 200 Re to over 2000 Re. Desch (1997) calculated travel times for LF burst emissions and found that observations required the emissions having had to travel as much as 2000 Re down the tail. Steinberg et al. (1998) used WAVES direction finding measurements to find that the spin-modulated high frequency portion of LF bursts exited the bow shock from 100 Re to 460 Re down the tail. Desch and Farrell (2000) studied a LF burst that was occulted by a strong increase of solar wind plasma density 97 Re down the tail. Thus the LF burst had to have entered the solar wind beyond 97 Re down the tail. Steinberg et al. (2003) analyzed 119 LF bursts observed by WIND WAVES when the spacecraft was near the Lagrange Point LI and concluded that the bow shock still exists beyond 1000 Re. A topic for future study is determining what solar wind parameters control the length of the Earth's geomagnetic tail. Another worthwhile topic to be pursued is trying to determine the reasons for the quasiperiodic nature of many geomagnetic disturbances and whether or not they contradict the importance of self organized criticality in magnetospheric dynamics. We have well demonstrated in this paper that GEOTAIL, POLAR, and WIND are still valuable resources for the study of the Earth's magnetosphere and the heliosphere. ACKNOWLEDGMENTS We thank all members of the ISTP team for the high quality data and for the successful spacecraft operation. R. R. Anderson appreciates support for his participation in this research under his Visiting Professorship at RASC, Kyoto University, and at The University of Iowa by NASA Contract NAS5-30371 and NASA Grants NAG5-2346, NAG5-7110, NAG5-7943, and NAG5-11707, all with Goddard Space Flight Center. Portions of this research were supported by Grant-in-Aid for Scientific Research (A) 12304026 from Japan Society for the Promotion of Science (JSPS).
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REFERENCES Acuna, M. H., K. W. Ogilvie, D. N. Baker, S. A. Curtis, D. H. Fairfield, and W. H. Mish, The Global Geospace Science Program and its investigations, Space Sci. Rev., 71, 5-21, 1995. Alexander, J. K., and M. L. Kaiser, Terrestrial Kilometric Radiation 1. Spatial Structure Studies, J. Geophys. Res., 81, 5948, 1976] Anderson, R. R., D. A. Gurnett, H. Matsumoto, K. Hashimoto, H. Kojima, Y. Kasaba, M. L. Kaiser, G. Rostoker, J.-L. Bougeret, J.-L. Steinberg, I. Nagano, and H. Singer, Observations of low frequency terrestrial type III bursts by GEOTAIL and WIND and their association with isolated geomagnetic disturbances detected by ground and space-borne instruments, Planetary Radio Emissions IV, Proc. Graz Conf., ed. by H. O. Rucker, S. J. Bauer, and A. Lecacheux, Austrian Academy of Sciences Press, Vienna, 241-250, 1997. Anderson, R. R., D. A. Gurnett, L. A. Frank, J. B. Sigwarth, H. Matsumoto, K. Hashimoto, H. Kojima, Y. Kasaba, M. L. Kaiser, G. Rostoker, J.-L. Bougeret, J.-L. Steinberg, I. Nagano, T. Murata, H. J. Singer, T. G. Onsager, and M. F. Thomsen, Geotail, Polar, Wind, CANOPUS, and ISTP associated geosynchronous satellite observations of plasma wave emissions and related magnetospheric phenomena during substorms, in Proceedings of International Conf. on Substorms-4, ISC-4, edited by S. Kokubun and Y. Kamide, pp. 567-572, Kluwer Academic Publishers, Dordrecht, London, Boston, and Terra Scientific Publishing Company, Tokyo, 1998. Anderson, R. R., H. Matsumoto, K. Hashimoto, H. Kojima, I. Nagano, Y. Kasaba, M. L. Kaiser, J.-L. Bougeret, and J. L. Steinberg, Using Geotail, Wind, and Polar observations of solar, interplanetary, and terrestrial plasma wave and radio emissions to identify source characteristics, in Planetary Radio Emissions V, ed. H. O. Rucker, M. L. Kaiser, and Y. Leblanc, pp. 297-310, 2001. Bougeret, J.-L., M. L. Kaiser, P. J. Kellogg, R. Manning, K. Goetz, S. J. Monson, N. Monge, L. Friel, C. A. Meetre, C. Perche, L. Sitruk, and S. Hoang, WAVES: The radio and plasma wave investigation on the WIND spacecraft, Space Sci. Rev., 71, 231-263, 1995. Desch, M. D., Terrestrial LF Bursts: Source and solar wind connection, in Planetary Radio Emissions IV, Proc. Graz Conf., ed. by H. O. Rucker, S. J. Bauer, and A. Lecacheux, Austrian Academy of Sciences Press, Vienna, 251-258, 1997. Desch, M. D., and W. M. Farrell, Terrestrial LF Bursts: Escape paths and Wave Intensification, in Radio Astronomy at Long Wavelengths, AGU Geophysical Monograph 119, R. G. Stone, K. W. Weiler, M. L. Goldstein, and J.-L. Bougeret, editors, pp. 205-211, 2000. Desch, M. D., M. L. Kaiser, and W. M. Farrell, Control of terrestrial low frequency bursts by solar wind speed, Geophys. Res. Lett., 23, 1251-1254, 1996. Ergun. R. E., C. W. Carlson, J.P. McFadden, F. S. Mozer, G. T. Delory, W. Peria, C. Chaston, M. Temerin, R. Elphic, R. Strangeway, R. Pfaff, C. A. Cattell, D. Klumpar, E. Shelly, W. Peterson, E. Moebius, and L. Kistler, FAST satellite wave observations in the AKR source region, Geophys. Res. Lett., 25, 2061, 1998. Fairfield, D. H., T. Mukai, A. T. Y. Lui, C. A. Cattell, G. D. Reeves, T. Nagai, G. Rostoker, H. J. Singer, M L. Kaiser, S. Kokubun, A. J. Lazarus, R. P. Lepping, M. Nakamura, J. T. Steinberg, K. Tsurda, D. J. Williams, and T. Yamamoto, Geotail observations of substorm onset in the inner magnetotail, J. Geophys. Res., 103, 103-117, 1998. Fairfield, D. H., T. Mukai, M. Brittnacher, G. D. Reeves, S. Kokubun, G. K. Parks, T. Nagai, H. Matsumoto, K. Hashimoto, D. A. Gurnett, and T. Yamamoto, Earthward flow bursts in the inner magnetotail and their relation to auroral brightenings, AKR intensifications, geosynchronous particle injections and magnetic activity, J. Geophys. Res., 104, 355-370, 1999. Frank, L. A., K. L. Ackerson, W. R. Paterson, J. A. Lee, M. R. English, and G. L. Pickett, Comprehensive plasma instrumentation (CPI) for the GEOTAIL spacecraft, J. Geomag. Geoelectr., 46, 23-37, 1994. Green, J. L., D. A. Gurnett, and S. D. Shawhan, The angular distribution of auroral kilometric radiation, J. Geophys. Res., 82, 1825, 1977. Green, J. L., and D. L. Gallagher, The detailed intensity distribution of the AKR emission cone, J. Geophys. Res., 90, 9641, 1985.
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Gurnett, D. A., The earth as a radio source: Terrestrial kilometric radiation, J. Geophys. Res., 79, 4227, 1974. Gurnett, D. A., A. M. Persoon, R. F. Randall, D. L. Odem, S. L. Remington, T. F. Averkamp, M. M. DeBower, G. B. Hospodarsky, R. L. Huff, D. L. Kirchner, M. A. Mitchell, B. T. Pham, J. R. Phillips, W. J. Schintler, P. Sheyko, and D. R. Tomash, The POLAR plasma wave instrument, Space Sci. Rev., 71, 597-622, 1995. Hashimoto, K., H. Matsumoto, T. Murata, M. L. Kaiser, and J.-L. Bougeret, Comparison of AKR simultaneously observed by the GEOTAIL and WIND spacecraft, Geophys. Res. Lett., 853-856, 1998. Imhof, W.L., K.A. Spear, J.W. Hamilton, B.R. Higgins, M.J. Murphy, J.G. Pronko, R.R. Vondrak, D.L. McKenzie, C.J. Rice, D. J. Gorney, D.A. Roux, R.L. Williams, J A. Stein, J. Bjordal, J. Stadsnes, K. Njoten, T.J. Rosenberg, L. Lutz, D. Detrick, The Polar Ionospheric X-ray Imaging Experiment, (Reprinted in The Global Geospace Mission, ed. by C.T. Russell, Kluwer Academic Publishers,1995), Space Science Reviews, 71, Nos. 1-4, 1995. Imhof, W. L., D. L. Chenette, D. W. Datlowe, J. Mobilia, M. Walt, and R. R. Anderson, The correlation of AKR waves with precipitating electrons as determined by plasma wave and X-ray image data from the POLAR spacecraft, Geophys. Res. Lett., 289-292, 1998. Imhof, W. L., R. R. Anderson, D. L. Chenette, J. Mobilia, S. M. Petrinec, and M. Walt, The correlation of rapid AKR variations with changes in the fluxes of precipitating electrons, Adv. in Sp. Res., 23, 1747-1752, 1999. Imhof, W. L., M. Walt, R. R. Anderson, D. L. Chenette, J. D. Hawley, J. Mobilia, and S. M. Petrinec, Association of electron precipitation with auroral kilometric radiation, J. Geophys. Res., 105, 277-289, 2000. Imhof, W. L., M. Walt, R. R. Anderson, J. D. Hawley, M. J. Brittnacher, S. M. Petrinec, and H. Matsumoto, Relationship between X-ray, ultraviolet, and kilometric radiation in the auroral region, J. Geophys. Res., 106, 10,479- 10,492, 2001. Imhof, W. L., R. R. Anderson, M. Walt, J. D. Hawley, S. M. Petrinec, J. Mobilia, and H. Matsumoto, The dependence of AKR production on the intensity and energy spectra of auroral bremsstrahlung, J. Geophys. Res., accepted for publication, Manuscript Number: 2002JA009274, 2003. Kaiser, M. L., and J. K. Alexander, Terrestrial kilometric radiation 3. Average spectral properties, J. Geophys. Res., 82, 3273, 1977a. Kaiser, M. L., and J. K. Alexander, Relationship between auroral substorms and the occurrence of terrestrial kilometric radiation, J. Geophys. Res., 82, 5283, 1977b. Kaiser, M. L., M. D. Desch, W. M. Farrell, J.-L. Steinberg, and M. J. Reiner, LF Band Terrestrial Radio Bursts Observed by Wind/WAVES", Geophys. Res. Lett., 23, 1283-1286, 1996. Kasaba, Y., H. Matsumoto, K. Hashimoto, and R. R. Anderson, The angular distribution of auroral kilometric radiation observed by the Geotail spacecraft, Geophys. Res. Lett., 24, 2483-2486, 1997. Kasaba, Y., H. Matsumoto, Y. Omura, R. R. Anderson, T. Mukai, Y. Saito, T. Yamamoto, and S. Kokubun, Statistical studies of plasma waves and backstreaming electrons in the terrestrial electron foreshock observed by Geotail, J. Geophys. Res., 105, 79-103, 2000. Kokubun, S., T. Yamamoto, M. Acuna, K. Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46,7,1994. Liou, K., C.-I Meng, T. Y. Lui, P. T. Newell, M. Brittnacher, G. Parks, G. D. Reeves, R. R. Anderson, and K. Yumoto, On relative timing in substorm onset signatures, J. Geophys. Res., 104, 22807-22817, 1999. Matsumoto, H., I. Nagano, R. R. Anderson, H. Kojima, K. Hashimoto, M. Tsutsui, T. Okada, I. Kimura, Y. Omura, and M. Okada, Plasma wave observations with GEOTAIL spacecraft, J. Geomag. Geoelectr., 46, 59-95, 1994. Murata, T., H. Matsumoto, H. Kojima, A. Fujita, T. Nagai, T. Yamamoto, and R. R. Anderson, Estimation of tail reconnection lines by AKR onsets and plasmoid entries observed with GEOTAIL spacecraft, Geophys. Res. Lett., 22, 1849-1852, 1995.
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Ogilvie, K. W., D. J. Chornay, R. J. Fitzenreiter, F. Hunsaker, J. Keller, J. Lobell, G. Miller, J. D. Scudder, E. C. Sittler, Jr., R. B. Torbert, D. Bodet, G. Needell, A. J. Lazarus, J. T. Steinberg, J. H. Tappan, A. Mavretic, and E. Gergin, SWE, A comprehensive plasma instrument for the WIND spacecraft, Space Sci. Rev., 71, 55-77, 1995. Reiner, M. J., M. L. Kaiser, J. Fainberg, M. D. Desch, and R. G. Stone, 2Fp radio emissions from the vicinity of the Earth's foreshock: WIND observations, Geophys. Res. Lett., 23, 1247-1250, 1996. Rostoker, G., J. C. Samson, F. Creutzberg, T. J. Hughes, D. R. McDiarmid, A. G. McNamara, A. Vallance Jones, D. D. Wallis, and L. L. Cogger, CANOPUS - A ground-based instrument array for remote sensing the high latitude ionosphere during the ISTP/GGS program, Space Sci. Rev., 71, 743-760, 1995. Russell, C. T., R. C. Snare, J. D. Means, D. Pierce, D. Dearborn, M. Larson, G. Barr, G. Le, The GGS/Polar magnetic-fields investigation, Space Science Reviews, 71 (1-4), 563-582, 1995. Steinberg, J.-L., C. Lacombe, and S. Hoang, A new component of terrestrial radio emission observed from ISEE-3 and ISEE-1 in the solar wind, Geophys. Res. Lett., 15, pp 176-179, 1988. Steinberg, J.-L., S. Hoang, and J. M. Bosqued, Isotropic kilometric radiation: A new component of the Earth's radio emission, et al., Ann. Geophysicae, 8, pp 671-686, 1990. Steinberg, J.-L., C. Lacombe, and S. Hoang, sounding the flanks of the Earth's bow shock to -230 Re: ISEE-3 observations of terrestrial radio sources down to 1.3 times the solar wind plasma frequency, J. Geophys. Res., 103, 23,565-23,579, 1998. Steinberg, J.-L., C. Lacombe, P. Zarka, S. Hoang, and C. Perche, Terrestrial Low frequency bursts: escape paths of radio waves through the bow shock, Planetary and Space Science, submitted, 2003. Voots, G. R., D. A. Gurnett, and S. -I. Akasofu, Auroral kilometric radiation as an indicator of auroral magnetic disturbances, J. Geophys. Res., 82, 2259, 1977. Wu, C. S., and L. C. Lee, A theory of the terrestrial kilometric radiation, Astrophys. J., 230, 621, 1979. e-mail: [email protected] and [email protected], [email protected], [email protected], [email protected], [email protected], [email protected], [email protected], [email protected], and [email protected]
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OCCULTATIONS OF AURORAL KILOMETRIC RADIATION IN THE VICINITY OF THE EARTH K. T. Murata1, W. Kurth2, K. Hashimoto3, and H. Matsumoto3 'Faculty of Engineering, Ehime University, 3 Bunkyo-cho, Matsumayama, Ehime 790-8577, JAPAN 2 Dept. of Physics and Astronomy, University of Iowa, Iowa City, I A, USA 3 Radio Science Center for Space and Atmosphere, Gokasho, Uji, Kyoto611-0011, JAPAN
ABSTRACT Auroral kilometric radiation (AKR) occupations in the vicinity of the Earth are studied using two observations by GEOTAIL and POLAR. We compared the dynamic spectra of both satellites for eight months paying attention to times and frequencies at which AKR is observed simultaneously. Then, we carefully examined the AKR illumination regions using the POLAR two-month orbit data. Two distinct regions where the AKR is occulted are found during the period. One is the region on the night side of the Earth, where the AKR does not propagate at frequencies >400 kHz. The other region is in the vicinity of the plasmapause, on both the day and night side of the Earth.
INTRODUCTION Auroral kilometric radiation (AKR) was discovered about 30 years ago by Benediktov et al. [1965]. Dunckel et al. [1970] related this radiation with auroral phenomena and Kurth et al. [1975] named the emission based on this relation. The first extensive study of the properties of AKR was provided by Gurnett [1974]. An important aspect of the study of AKR is its propagation characteristics. Using simultaneous observations by two or more satellites, Green and Gallagher [1985] and Green and Gurnett [1979] ascertained the extent of illumination cones of the emission. Additional studies of AKR propagation properties have been carried out either by satellite observations or by ray tracing computer simulations. The former studies have mainly employed dynamic spectrum data from multichannel analyzers, the latter were ray-tracing calculations. These studies have revealed that AKR propagation on the day and the night sides of the Earth is strongly dependent on the location of the plasmapause boundary and the AKR source location. OVERVIEW OF DYNAMIC SPECTRA OBTAINED BY GEOTAIL ANS POLAR In the present study, we investigate the regions illuminated by AKR. We compare the dynamic spectrum from two satellites, GEOTAIL and POLAR, which are usually in positions favorable for the observation of AKR [Matsumoto et al., 1994]. In this study, data from one satellite are used as a reference for the observations by the other. First, we carefully compare dynamic spectra for eight months (from April to November 1996). Figure 1 shows the dynamic spectra provided by GEOTAIL (lower panel) and POLAR (upper panel). Panel A is a case in which both satellites observe AKR simultaneously. It is notable that AKR is observed by both spacecraft at the same times and over the same frequency range. We found this correspondence in time and frequency on many days. Note that the signals in the POLAR wideband spectra around 15:30 is not AKR but emissions associated with the plasmapause.
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In panel B, AKR observations do not match either in time or in frequency. Between 12:00 and 16:00 UT both of the satellites observe AKR, but the upper frequency extent is different; POLAR does not observe AKR above 400 kHz. We found several cases of upper frequency cutoff differences during the examined period. The cutoff frequencies in these cases are always about 400 kHz and last longer than 3 hours. In other cases AKR occultation occurs in the vicinity of the plasmapause as found at 5:00 UT and 22:00 UT. We often find occultation periods of a few tens of minutes before or after the entry of POLAR into the plasmapause during which AKR is not observed. We examine additional POLAR examples similar to those in panel B in the present paper.
Fig.1 One day plots of simultaneous observations of AKR by POLAR (upper panel) and GEOTAIL (lower panel): (A) 13 May 1996 and (B) 4 May 1996, respectively. Color contour is in unit of dBV/m/Hz1'2.
AKR OCCULTATION ON THE POLAR ORBIT We first compared AKR dynamic spectra provided by the satellites during April and May 1996 when POLAR situated in the vicinity of the meridian plane. The differences in the upper cutoff frequency were found only when GEOTAIL observed AKR extending above 400 kHz. In Fig.2, the symbols represent POLAR positions (in SM coordinates) when the upper cutoff of the AKR observed is lower than 400 kHz. Small dots are given to show the entire POLAR orbital position coverage for these two months. We find that the region where high frequency AKR (> 400 kHz) does not propagate is concentrated between 18 and 24 MLT. The magnetic latitude range of this region extends from 20 to 80 degrees. It should be noted that the geocentric radial distance of the POLAR orbit in this region is 7 to 9 RE. Also of interest is the AKR occultation in the vicinity of the plasmapause as seen in panel B of Fig. 1. The orbital inclination of POLAR is almost 90 degrees, which places its apogee over the north polar region at a geocentric radial distance of 9 RE and its perigee over the south polar region at 2 RE. Thus the satellite encounters the day side and night side plasmapause once per orbit. Almost always before and after its plasmapause entry, POLAR passes through the AKR occultation region in the northern hemisphere. In the present paper, we discuss only the northern hemisphere since the satellite moves quickly in the southern hemisphere and it is difficult, therefore, to estimate the satellite position with high accuracy. Figure 3 shows the high latitude boundaries of the AKR occultation region measured at 250 kHz. At the lower magnetic latitude of each point, POLAR did not observe the AKR. These illumination boundaries are between 30 and 50 MLAT on the day side, and between 10 and 30 MLAT on the night side, respectively. The geocentric radial distance of the spacecraft at these positions is about 4 RE. CONCLUSIONS AND DISCUSSION In the present paper, we compare dynamic spectra of AKR provided by GEOTAIL and POLAR for a twomonth interval. Attention is paid to the AKR propagation in the near-Earth magnetosphere. This kind of multi-221-
satellite observation is important since one AKR observation can be used as a reference for the other excitation. Our observational results are summarized as follows: 1. In the high latitude region, we found that the high frequency AKR (> 400 kHz) does not often propagate for the examined period (April to November 1996), and shows no seasonal dependency. This high frequency AKR occultation region is found mainly in the night side at lower region than 60 MLAT. 2. In the vicinity of the plasmapause, there are cavities of the AKR illumination both in the night and day side of the Earth. The cavity is between 30 and 50 MLAT on the day side and between 10 and 30 MLAT on the night side. Green et al. [1977] and Hashimoto [1984] have performed 2-D and 3-D ray-tracing simulations for the AKR in the vicinity of the Earth, respectively. Hashimoto suggested that the lowest latitude of L-0 mode waves at 250 kHz, which tends to be lower than R-X mode, varies from 50 to Fig.2 The tiny dots represent the POLAR orbit in April and 10 degrees on the night side. This variation May 1996. The large symbols are the POLAR locations depends on the wave normal angles and source when it did not observe high frequency AKR (> 400 kHz). location (MLT) of the AKR. This ray-tracing result is consistent with the present results. However, further detailed comparisons of event to event correspondence between the AKR source locations and illumination boundaries are required. Green et al. [1977] have shown in their 2-D ray-tracing simulations that the highest latitudes for AKR propagation on the day side are independent of the size of the plasmapause. This does not correspond to the present results. The boundary latitude varies by as large as 20 degrees as shown in Fig. 3. This variation should be explained by the fact that the AKR source locations move, mainly from the midnight sector to the day side, as the aurora progresses. We have found that in the vicinity of the midnight meridian there is a region where high frequency AKR (> 400 kHz) does not propagate. The ray-tracing simulations so far don't explain why only the low frequency AKR (< 400kHz) propagates to this region. More collaborative studies between simulations and observations are required to understand this result. The present conclusions are derived from twomonth data analyses. We need longer duration Fig.3 Each dot represents MLAT and MLT of the AKR research to obtain general trend on the AKR illumination boundary on the POLAR orbit. At the lower occultation region in the vicinity of the Earth. latitude of each point, POLAR is not illuminated by AKR.
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REFERENCES Benediktov, E. A., G. G. Getmantsev , Yu. A. Sazonov, and A. F. Tarasov, Preliminary Results of Measurements of the Intensity of Distributed Extraterrestrial Radio-Frequency Emission at 725 and 1525-kHz Frequencies by the Satellite Electoron-2 , Kosm. Issled., 3, 614, 1965. Dunckel, N., B. Ficklin, L.Rorden, and R.A. Helliwell, Low frequency noise observed in the distant magnetosphere with OGO 1, J. Geophys. Res., 75, 1854-1862, 1970. Green, J. L. and D. A. Gumett, A Correlation Between Auroral Kilometric Radiation and Inverted V Electron Precipitation, J. Geophys. Res., 84, 5216, 1979. Green, J. L., D. A. Gurnett and S. D. Shawhan, The Angular Distribution of Auroral Kilometric Radiation, J. Geophys. Res., 82, 1825, 1977. Green, J. L. and D. L. Gallagher, The Detailed Intensity Distribution oftheAKR Emission Cone, J. Geophys. Res.,
90,9641, 1985. Gurnett, D. A., The Earth as a Radio Source: Terrestrial Kilometric Radiation, J. Geophys. Res., 79, 4227, 1974. Hashimoto, K., A reconciliation of propagation modes of auroral kilometric radiation, J. Geophys. Res., 89, 7459, 1984. Kurth, W. S., M.M. Baumback, and D.A. Gurnett, Direction-finding measurements of auroral kilometric radiation, J. Geophys. Res., 80, 2764-2770, 1975. Matsumoto, H., I. Nagano. R. R. Anderson, H. Kojima, K. Hasimoto, M. Tsutsui, T. Okada, I. Kimura, Y. Omura and M. Okada, Plasma Wave Observations with GEOTAIL spacecraft, J. Geomag. Geoelectr., 46, 59, 1994.
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LOBE TRAPPED CONTINUUM RADIATION GENERATED IN THE DISTANT MAGNETOTAIL H. Takano1, I. Nagano1, S. Yagitani1, and H. Matsumoto2 1
Graduate School of Natural Science and Technology, Kanazawa University, 2-40-20 Kodatsuno, Kanazawa 920-8667, Japan 2 Radio Science Center for Space & Atmosphere, Kyoto University, Gokasho, Uji, Kyoto 611-0011, Japan ABSTRACT Lobe trapped continuum radiation (LTCR) has been observed by the plasma wave instrument (PWI) onboard the GEOTAIL spacecraft at frequencies as low as 1 kHz in the distant geomagnetic tail region. Prom the direction finding analysis with the wave form capture (WFC) data, the arrival directions of LTCR are almost parallel to the dawn-dusk direction. A 3-D ray tracing analysis shows that the initially radiated Earth-tail ray directions of LTCR are transformed into the dawn-dusk directions by the reflection at the cylindrical tail magnetopause. The propagation characteristics of LTCR can give us very important information on the macroscopic structure such as magnetotail flapping. By comparing the results of the direction finding with the 3-D ray tracing analysis, we find that the possible source regions for LTCR are located at the plasma sheet boundary layer away from the nominal tail axis and the low latitude boundary layer.
INTRODUCTION Continuum radiation (CR) has been widely observed by many spacecraft in the magnetosphere. CR is detected in the wide frequency range from 5 to 100 kHz, and is believed to be generated at the geomagnetic equator of the plasmapause from 4 to 14 hours local time zone (Gurnett and Shaw, 1973; Gurnett, 1975). Gurnett (1975) classified CR into lower-frequency trapped components (5~20 kHz) and higher-frequency escaping components (20~100 kHz). CR is believed to be generated through a linear mode conversion from an electrostatic wave near the upper hybrid frequency (fuHR) hi a large density gradient perpendicular to the geomagnetic field at the plasmapause (Jones, 1976). The frequency range of CR indicates the electron density at the plasmapause, where fuHR is close to the local plasma frequency(/p). On the other hand, the low frequency component of CR has been observed by ISEE-3 and GEOTAIL spacecraft in the distant magnetotail (Coroniti et al., 1984; Nagano et al., 1994). Nagano et al. (1994) have shown that the low frequency component of CR at frequencies extending down to 1 kHz was observed along with an electron cyclotron harmonic (ECH) wave in the lobe region, and suggested that such a low frequency component of CR, called "Lobe Trapped CR" (LTCR), is generated locally from the ECH wave near the plasma sheet boundary layer (PSBL). However, it has not been clear how the source locations of LTCR distribute in the distant tail region. In this paper, we estimate the distribution of source regions for LTCR in the distant magnetotail by comparing the direction finding results with the 3-D ray tracing calculations. GEOTAIL OBSERVATION OF LTCR The sweep frequency analyzer (SFA) of the plasma wave instrument (PWI) onboard GEOTAIL can measure the LTCR spectrum in the distant magnetotail. A typical observation of LTCR is shown in Figure 1 which is the dynamic spectrum of the electric field component measured by SFA from 13:00 UT to 15:00 UT on January 26, 1993. In this plot GEOTAIL mainly surveys the tail lobe region at 90 RE away from -224-
the Earth after GEOTAIL traverses the transition region between the PSBL and tail lobe at 13:10 UT. The normal CR which consists of the trapped and escaping components of CR generated at the equatorial plasmapause is observed continuously in the frequency range from 6 kHz to 40 kHz. The LTCR is detected from 13:22 to 13:45 UT, 14:00 to 14:15 UT, and 14:20 to 15:00 UT in the frequency range from 1 kHz to 4kHz. The LTCR spectrum is similar to the CR spectrum with a continuous structure. LTCR is trapped in the low-density tail lobe region by being reflected at high density regions such as the plasma sheet and the magnetopause.
Fig. 1. The dynamic spectrum of LTCR measured by SFA from 13:00 UT to 15:00 UT on January 26, 1993.
DIRECTION FINDING AND 3-D RAY TRACING ANALYSIS A direction finding analysis has been used to estimate the source positions of electromagnetic waves. The wave form capture (WFC) which is another subsystem in the PWI can simultaneously measure the wave forms of two electric field and three magnetic field components for 8.7 seconds. WFC observation provides the wave normal direction, Poynting flux and polarization of electromagnetic waves as well as the detailed spectral structures (Matsumoto et al., 1994). Unfortunately, the 3-D wave normal direction and Poynting flux of LTCR cannot be calculated from WFC data because of insufficient sensitivity of the magnetic search coil for measurement of LTCR magnetic field components. The arrival direction of LTCR on the spin plane is calculated only from the spin-modulated electric filed intensities (Manning and Fainberg, 1980). The arrival directions calculated with WFC data on the X-Y plane of the GSE coordinate system are shown in Figure 2 (a). The circles are the observation positions and the black bars show the arrival directions averaged over the LTCR frequency band. The angular range of arrival direction is obtained only from 0 to 180° because of the symmetry of the dipole antenna pattern. The arrival directions of LTCR are almost parallel to the dawn-dusk direction. Figure 2 (b) shows the occurrence distribution of the arrival directions of LTCR obtained by the direction finding analysis. "Event" is the frequency-averaged arrival direction and "All" shows the arrival direction obtained every 10 Hz. The dawn-dusk direction is (p = 90° and the Earth-tail direction is <j> = 0 or 180°. The occurrence is highest around
Fig. 2. The direction finding results of the LTCR with WFC data, (a) The arrival directions of the LTCR on the X-Y plane of the GSE coordinate system, (b) The occurrence distribution of the arrival directions of the LTCR obtained by the direction finding analysis. -225-
A 3-D ray tracing calculation provides us with an additional tool to probe the LTCR propagation in the distant magnetotail. Here the computer code of ray-tracing calculations developed by Nagano et al., (1998) is modified for the analysis of the LTCR propagation in the distant magnetotail. The ray equations of cold plasma waves are integrated with fourth-order Runge-Kutta and predictor/corrector scheme with step-size control. The electron density and magnetic field models are based on the results of the 3-D MHD simulation of the solar wind/magnetosphere interaction (Ogino et al., 1994). LTCR would be generated through a linear mode conversion from a Z-mode electromagnetic wave in the PSBL. With a linear mode conversion (Jones et al., 1987), LTCR is generated as L-0 mode waves emitted near the local plasma frequency from the radio window centered at #RVV = tan~ 1 (/ p e // c e ) with respect to the local magnetic field line, where fpe and / c e are the local electron plasma and cyclotron frequencies. Figure 3 (a) and (b) show the results of the 3-D ray tracing calculations for LTCR in two cases. In this calculation, the initial ray locations are chosen as (X,Y,Z) = (-65 R E , 1 R E , 2 R E ) in Figure 3 (a) and (-65 RE, 10 RE, 2 R E ) in Figure 3 (b). In the case of Y = 1 RE, the LTCR emitted tailward mainly propagates along the Earth-tail direction and is trapped in the lobe region with multiple reflections between the magnetopause and the plasma sheet. On the other hand, the ray radiated at Y = 10 RE in Figure 3 (b) propagates to the Earth-tail and dawn-dusk directions. The initially radiated Earth-tail ray directions are transformed into the dawn-dusk directions by the reflections at the cylindrical tail magnetopause. The arrival directions obtained from WFC measurement are consistent with the ray directions emitted from the plasma sheet boundary layer away from the nominal tail axis along the Earth-tail direction.
Fig. 3. The results of 3-D ray tracing calculation. The initial rays are radiated at (a) Y = 1 RE and (b) Y = 10 RE- The thick lines show the 3-D ray paths and thin lines are their projection onto the X-Y, Y-Z and X-Z planes. The circle shows the source location and the arrows drawn on the planes are the initial ray directions.
DISTRIBUTION OF SOURCE REGION FOR LTCR The estimation of the possible source region of LTCR is carried out by comparing the results of the direction finding with the 3-D ray tracing calculation. The direction finding analysis with WFC data indicates that the arrival directions of LTCR are almost parallel to the dawn-dusk direction (65° < <j> < 115°). By using the 3-D ray tracing analysis, the regions radiating the ray consistent with the direction finding result are estimated as the possible source locations. The estimation method of the source regions for LTCR is shown as follows. The probability of the time interval with the ray azimuth angle
locations are 1 RE in XQSE and 0.5 RE in YCSE- Figure 4 shows the spatial distribution of the source regions of LTCR at (a) Z = 1.6 RE and (b) Z = 2.0 RE. The elongated distribution with the high probability at 13 RE < I^GSE! ^ 23 RE and —130 RE < XQSE <• ~30 RE is observed as the possible source regions. Such regions are the transition regions between PSBL and LLBL. The PSBL away from the nominal tail axis is another possible source region. In the case of Z = 1.6 RE, the PSBL as the source region of LTCR is widely spread in the distant tail. On the other hand, the PSBL source region of LTCR radiated away from the central plasma sheet {Z = 2.0 RE) is concentrated around the middle tail from —65 RE to —85 RE in XQSEThe distribution of PSBL source region varies with the distance from the central plasma sheet.
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Fig. 4. The spatial distribution of source region for LTCR on the X-Y plane of the GSE coordinate system at (a) Z = 1.6 RE and (b) Z = 2.0 RE- The probability within the range of 0 to 1 is shown by the gray scale on the right of each figure.
CONCLUSION The direction finding analysis with WFC data and 3-D ray tracing calculation have provided the propagation characteristics of LTCR and the distribution of possible source regions in the distant magnetotail. The direction finding analysis with WFC data has shown that LTCR mainly propagates along the dawn-dusk direction. By comparing the results of the direction finding with the 3-D ray tracing calculation, the possible source regions of LTCR are found to be located at the PSBL away from the nominal tail axis and the LLBL. The estimation of the LTCR source regions makes it possible to study the generation mechanism of LTCR in the distant magnetotail. ACKNOWLEDGEMENTS We would like to thank Prof. T. Ogino for providing the electron density and magnetic field data obtained by his simulation. We thank GEOTAIL/PWI members for their advice and suggestions. Part of the computation in the present study was provided by and performed with the KDK system of Radio Science Center for Space and Atmosphere (RASC) at Kyoto University as a collaborative research project. REFERENCES Coroniti, F. V., F. L. Scarf, C. F. Kennel, and D. A. Gurnett, Continuum radiation and electron plasma oscillations in the distant geomagnetic tail, Geophys. Res. Lett,, 11, 661-664, 1984. Gurnett, D. A., and R. R. Shaw, Electromagnetic radiation trapped in the magnetosphere above the plasma frequency, J. Geophys. Res., 78, 8136-8149, 1973. Gurnett, D. A., The Earth as a radio source : The nonthermal continuum, J. Geophys. Res., 80, 2751-2763, 1975. Jones, D., Source of terrestrial nonthermal radiation, Nature, 260, 686-689, 1976. Jones, D., W. Calvert, D. A. Gurnett, and R. L. Huff, Observed beaming of terrestrial myriametric radiation, Nature, 328, 391-395, 1987. Manning, R., and J. Fainberg, A new method of measuring radio source parameters of a partially polarized distributed source from spacecraft observations, Space Science lustrum., 5, 161-181, 1980. Matsumoto, H., I. Nagano, R. R. Anderson, H. Kojima, K. Hashimoto, M. Tsutsui, T. Okada, I. Kimura. Y. Omura, and M. Okada, Plasma wave observations with GEOTAIL spacecraft, J. Geomag. Geoelectr.. 46, 59-95, 1994. Nagano, I., S. Yagitani, H. Kojima, Y. Kakehi, T. Shiozaki, H. Matsumoto, K. Hashimoto, T. Okada, S. Kokubun, and T. Yamamoto, Wave form analysis of the continuum radiation observed by GEOTAIL, Geophys. Res. Lett, 21, 2911-2914, 1994. Nagano, I., X.-Y. Wu, S. Yagitani, K. Miyamura, and H. Matsumoto, Unusual whistler with very large dispersion near the magnetopause: Geotail observation and ray-tracing modeling, J. Geophys. Res., 103, 11827-11840, 1998. Ogino, T., R. J. Walker and M. Ashour-Abdalla, A global magnetohydrodynamic simulation of the response of the magnetosphere to a northward turning of the interplanetary magnetic field, J. Geophys. Res., 99, 11027-11042, 1994. -227-
WHISTLER MODE CHORUS OBSERVED AROUND THE PLASMAPAUSE DURING MAGNETIC STORMS Y. Kasahara1, H. Uchiyama2, and Y. Goto1 1
Kanazawa University, 2-40-20 Kodatsuno, Kanazawa, 920-8667, Japan 2 Kyoto University, Yoshida, Sakyo-ku, Kyoto, 606-8501, Japan ABSTRACT
Drastic changes of relativistic electron population in the outer radiation belt during magnetic storms have been observed by many satellites. The relativistic electron flux decreases during a main phase of magnetic storm, and in some cases increases to above the prestorm level during the recovery phase. One of the plausible mechanisms of the relativistic electron enhancement during the recovery phase of magnetic storm is the wave-particle resonant diffusion by whistler mode chorus. In the present paper, we statistically studied chorus emissions observed by Akebono in the vicinity of the outer radiation belt associated with magnetic storms. In order to evaluate the contribution of whistler mode chorus to the relativistic electron enhancement, we performed direction finding analysis of the chorus using wave distribution function (WDF) method and quantitatively estimated the energies of whistler mode wave and relativistic electrons during recovery phase of magnetic storm.
INTRODUCTION More than 13 years have passed since the Akebono satellite was launched. The VLF instruments on board the Akebono satellite were designed to investigate the plasma waves in ELF/VLF range below 20 kHz (Kimura et al., 1990; Hashimoto et al., 1997). The instruments also have a function to determine the wave normal vector of these waves. Using the long period observation datasets obtained by Akebono, the time variation and spatial distribution of VLF waves in the inner magnetosphere was statistically investigated (Kasahara et al., 2001). According to their statistical study, there are two active regions of VLF waves above 1kHz in the inner magnetosphere; one is along the auroral region and the other is along the plasmapause. They showed that the wave activity along the auroral field lines is dominantly auroral hiss and the wave in the vicinity of the plasmapause is dominantly chorus emission. The strongest chorus is observed in the dawn to noon sector, and the chorus in the higher frequency range is distributed along the smaller L-value in the earlier local time region. On the other hand, drastic changes of relativistic electron population in the outer radiation belt during magnetic storms have been observed by many satellites. The relativistic electron flux decreases during main phase of magnetic storm, and in some cases increases to above the prestorm level during the recovery phase (e.g., Li et al., 1997; Obara et al., 2000, and references therein). The acceleration mechanism of relativistic electron was discussed by many authors for a long time. Fujimoto and Nishida (1990) proposed a recurrent type of acceleration, but there are so many types of magnetic storm that their acceleration mechanism is not always consistent with all types of relativistic electron variability during magnetic storms. According to the review papers by Li and Temerin (2001) and Friedel et al. (2002), several mechanisms for the enhancement of relativistic electrons were introduced but the primary causes have not been clarified yet. One of the plausible mechanisms is the wave-particle resonant diffusion by whistler mode chorus (Summers et al., 1998). Summers et al. (1998) suggested that whistler mode chorus is generated by the ring current -228-
electrons in the energy range from a few tens keV to 100 keV during recovery phase of magnetic storms and the population of relativistic electrons in the outer radiation belt gradually increases related to the generation of chorus. Summers and Ma (2000) theoretically showed that the acceleration process through the gyroresonant electron-whistler mode chorus interaction could account for the gradual flux increase during the long-lasting recovery phase. In their model, however, they assumed the wave spectra to be Kolmogorov and the propagation direction of the wave to be parallel to the geomagnetic field line. From the observational point of view, Meredith et al. (2002) investigated plasma wave and electron data observed by the CRRES satellite. In particular, they examined the temporal evolution of the spectral response of the electrons and the wave during October 1990 geomagnetic storm and suggested that wave-particle interaction contributed to the reformation of the relativistic electrons and this mechanism can be most effective when the recovery phase is characterized by prolonged substorm activity. On the other hand, Miyoshi et al. (2003) introduced one observational example of such magnetic storms using the data from the NOAA and Akebono satellites. They showed that the relativistic electron flux recovered gradually around the outside of plasmapause with their energy spectrum changing from soft to hard from the inner portion of the outer belt, and intense whistler mode chorus was simultaneously observed in the whole outer belt. Miyoshi et al. (2003) concluded that this result is an evidence of the acceleration of relativistic electrons by whistler mode chorus generated in the outer radiation belt, although the contribution of other acceleration mechanisms such as ULF waves and adiabatic process should be also taken into account to understand the dynamics of outer radiation belt in individual magnetic storms. Summers et al. (2002) developed the model proposed by Summers and Ma (2000). They utilize both particle and wave data observed by CRRES satellite. They concluded that their model agreed with the observational data well and the mechanism of stochastic acceleration by whistler mode chorus is a viable candidate for generating relativistic electrons. It should be also noted, however, that Miyoshi et al. (2003) and Summers et al. (2002) also assumed that the propagation direction of the waves were parallel to the geomagnetic field lines. In order to discuss the contribution of whistler mode chorus to the acceleration process of relativistic electrons quantitatively, we need to estimate the total wave energy of the chorus generated during the magnetic storms. Furthermore, wave normal direction of the whistler mode chorus is quite important information for the estimations of total wave energy and energy transfer rate from chorus to relativistic electrons by wave-particle interaction. In the present paper, we firstly investigate time variation of chorus emissions in the vicinity of the outer radiation belt associated with magnetic storms. As a next step, we perform direction finding analysis of the chorus using wave distribution function (WDF) method. Finally we quantitatively estimate the energies of whistler mode wave and relativistic electrons during recovery phase of magnetic storm and evaluate the contribution of whistler mode chorus to the relativistic electron enhancement. GLOBAL DISTRIBUTION OF CHORUS DURING MAGNETIC STORMS The VLF instrument onboard Akebono consists of five kinds of subsystems (Kimura et al., 1990; Hashimoto et al., 1997). The MCA (Multi Channel Analyzer) is one of the subsystems of the VLF instrument. MCA measures one component of each electric and magnetic field with 16 channels of band-pass filters in a frequency range from 3.18 Hz to 17.8 kHz with 0.5 sec time resolution. Using the data from MCA observed from 1989 to 1998, the time variation and spatial distribution of the VLF wave activity during magnetic storms is statistically studied. In the analysis, we superposed 41 events of magnetic storms occurred from 1989 to 1998 adjusting the zero point at the beginning time of each storm recovery and investigated averaged wave intensity as a function of time from the zero point. Based on the statistical results by Kasahara et al. (2001), we used the data which was obtained in the magnetic local time region between 0 and 13, where the chorus was mainly observed. Fig.l shows the averaged electric wave intensity in the frequency range from 1 to 10 kHz between two days before and five days after the beginning of the storm recovery. In the figure, the vertical axis indicates the L-value where the waves were observed. It is found that the wave intensity is enhanced from the late main phase and the intensity becomes largest just after the beginning of the storm recovery. These intense VLF waves were dominantly chorus emissions by examining the detailed spectra of the waves. An important point to be noted is that the active region of the chorus is located at L ~ 2.5 in the beginning of the recovery -229-
Fig. 1. Time variation of electric wave intensity during magnetic storm.
Fig. 2. Chorus observed on May 2, 1991. The thick line at 10.2 kHz shows the tuned frequency of the PFX system.
phase and gradually shifts toward the larger L-value region (L ~4) in the late recovery phase. The projected map to the equatorial plane of the active region of the chorus possibly corresponds with the region of ring current electrons, which are thought to be the energy source of the chorus. DIRECTION FINDING The VLF instrument onboard Akebono has another subsystem named PFX (Poynting FluX) analyzer. PFX measures two components of electric field and three components of magnetic field with an output bandwidth of 50 Hz in a frequency range from 100 Hz to 12.75 kHz. The center frequency of PFX is either automatically stepped by the minimum step of 50 Hz or kept fixed at a constant frequency. Using these five components of wave data, the wave normal direction and Poynting vector can be determined. In the present study, we determined the wave normal and Poynting direction of the chorus using wave distribution function (WDF) method (Storey and Lefeuvre, 1979; 1980), which is the most suitable technique for direction finding for natural waves. The WDF is derived from the concept that observed signals consist of a number of elementary plane waves and can be defined as the distribution of the wave energy density relative to the angular frequency ui and to the variables 9 and <j>, where 6 and
f* aijiuAftFiuA^ainOdOdt,
(i,j = 1,2,3,4,5),
(1)
JO
where aij(b},6,) are integration kernels introduced by Storey and Lefeuvre (1980), and F(u),0,(p) is the distribution function for a wave with its frequency u) and the wave normal direction (0,<j>). The integration kernels a^ are calculated theoretically from the local plasma frequency and cyclotron frequency. Knowing the integration kernels a^ and the spectrum matrix 5y, we can estimate the function F by solving the set of integral equations (1). However, the WDF method is so called ill-posed problem and we need some hypotheses to solve (1) such as one-or-two direction model, or maximum entropy model. In the present paper, we utilized a Gaussian WDF model (Goto et al., 2000). Because the time duration of each chorus component is so short (typically, a few hundreds milli-seconds) and furthermore the frequency variation of each chorus element is so fast (several kHz/sec), we need to know instantaneous intensity and phase of the wave in order to determine the wave normal direction. Therefore we used analytic signal method (ASM) instead of normal FFT analysis in the calculation of spectral matrix. In the ASM method, firstly the positive part of FFT spectra is doubled and the negative part of FFT spectra -230-
Fig. 4. Whistler mode wave energy with its electric amplitude of lmV/m versus wave frequency as a function of wave normal angle.
Fig. 3. Wave normal and Poynting directions of the chorus determined by wave distribution function method. Fig. 5. Time variation of wave energy during magnetic storm.
is deleted, and secondly the reverse-FFT is performed so as to determine the instantaneous amplitude and phase of the signal separately. In the present study, we analyzed the chorus observed at altitude of 6700km at latitude of 32°. The wave spectrum of the chorus used for the analysis is shown in Fig.2. In the event, PFX is tuned at 10.2 kHz (solid line in Fig.2). The upper panel in Fig.3 shows the spectrum of the chorus observed by PFX around the fixed frequency at 10.2 kHz. The next two panels in Fig.3 show the wave normal directions (0, (j>) determined by the WDF method. The wave normal directions of the chorus have quite large angle from the magnetic field line in the meridian plane and some of them directed nearly resonance angle. The last two panels in Fig.3 show the Poynting directions (8p, p). Poynting direction is almost along magnetic field lines in the meridian plane. We also confirmed that these chorus elements were all propagating away from the equatorial plane. Thus it is found that the chorus emissions in this event generally propagate along the same magnetic field line with quite large wave normal directions from lower to higher latitude region. This condition is quite favorable for the interaction between wave and relativistic electrons because waves can interact the larger energetic electrons in the same field lines. ESTIMATION OF WAVE ENERGY In order to evaluate the contribution of wave-particle resonant diffusion mechanism, it is necessary to estimate the total energy of VLF chorus and compare the increase of energy of relativistic electrons during recovery phase of magnetic storm. Fig.4 shows the relation between whistler mode wave energy versus frequency of the wave as a function of wave normal direction from 60° to 80°. The wave energy in the vertical axis is normalized by the electric wave amplitude of 1 mV/m. As shown in the figure, the wave energy depends on both wave frequency and wave normal directions. Using the same datasets shown in Fig.l, we statistically derived the averaged amplitude of chorus at each frequency point of the MCA subsystem as a function of time and L-value. Then we can estimate the averaged -231-
Fig. 6. Integrated wave energy at each L-shell region.
Fig. 7. Generation region of VLF chorus in the equatorial plane.
energy of the VLF chorus generated during magnetic storm multiplying the square of wave amplitude to the normalized wave energy at each frequency point and integrating them. Fig.5 shows the integrated wave energy from 1 to 10kHz at each time and L-shell region. The vertical and horizontal axes are the same used in Figl. In the calculation, we assumed the wave normal angle of the chorus to be 70°based on the direction finding results. Secondly, by integrating the wave energy with respect to time from zero to 5 days after the beginning of storm recovery at each L-shell region, we statistically estimated total Poynting flux passing through at each L-shell region as shown in Fig.6. For example, 48 hours after the beginning of the storm recovery, the integrated energy at the L-shell region between 4 and 5 is about 6.44 x 10~3[J/m2]. It is noted that Poynting flux is represented by the energy passing through a unit area, we therefore need to assume the area of the distribution region of the chorus in order to estimate the total wave energy generated in the equatorial region. Based on the statistical results (Kasahara et al., 2001), we assumed that chorus is generated in the magnetic local time region from 0 to 13 (see Fig.7). Then we derived that the total wave energy at the L-shell region between 4 and 5 for 48 hours during storm recovery phase is about 8.0 x 1012 [J] On the other hand, we roughly estimated the energy increase of the relativistic electrons during recovery phase of the magnetic storm occurred in early November, 1993. Referring the data from the radiation monitor (RDM) onboard Akebono introduced by Obara et al. (2000) (see Fig.2 in their paper), we took into account the conversion factors of the RDM (private communication, Y. Miyoshi and T. Nagai, 2002) and derived the energy density of relativistic electrons, which is represented by the energy per unit volume. By multiplying the volume of the region where the relativistic electrons are filled in, we finally derived that the total energy increase of relativistic electrons in 48 hours after the beginning of the storm recovery in this event was about 1.3 x 1011 [J]. SUMMARY In the present paper, we statistically investigated the time variation and spatial distribution of VLF chorus emissions in the vicinity of the outer radiation belt associated with magnetic storms. Our statistical study clarified that the generation region of the chorus during the recovery phase of the magnetic storms is firstly located around the inner edge of the outer radiation belt, and gradually shifts toward the outer region. Our direction finding study showed that the chorus emissions generally propagate downward (toward higher latitude region) along the geomagnetic field line with large wave normal angles. It is generally accepted that chorus emissions are generated with their wave normal almost parallel to the geomagnetic field lines at the source region, but it should be noted that our direction finding analysis was performed to the data observed off-equatorial region. Therefore these results support the conventional theory that the chorus emissions are generated by the ring current electrons injected around the equatorial region and they propagate toward higher latitude region. -232-
Using the information on the time variation of the wave intensities associated with magnetic storms (Fig.l) and wave normal directions (Fig.3), we evaluated the total energies of whistler mode chorus and relativistic electrons. Although the estimation in the present study is quite rough, we found that the energy of VLF chorus is several tens times larger than the energy increase of relativistic electrons. It is also noted that we used the statistically averaged wave amplitude in the calculation, the wave energy would be possibly under estimation in some cases. As a future study, we need to evaluate more quantitatively for each storm events and discuss the detailed relationship between wave and particle distribution. For example, we only show one example of the direction finding analysis in the present paper, but we need to examine the spatial dependence of the wave normal direction of the wave statistically by using the enormous datasets of the Akebono observation. Furthermore we should also take into account the energy transfer rate from chorus to relativistic electrons by waveparticle interaction. The Akebono satellite is still successfully operated and the datasets of Akebono will be an important clue to understand a generation mechanism of relativistic electrons in the recovery phase of magnetic storm. ACKNOWLEDGMENTS The authors wish to thank Y. Miyoshi and T. Obara for their helpful comments and suggestions throughout the present work. We would like to thank I. Kimura, I. Nagano and the other Akebono/VLF members for their contribution to the VLF instrument and its data analysis. We also acknowledge K. Tsuruda, T. Mukai and the other Akebono mission project members for their support. REFERENCES Freidel, R. H. W., G. D. Reeves, and T. Obara, Relativistic electron dynamics in the inner magnetosphere - a review -, J. Atmos. Solar Terr. Phys., 64, 265-282, 2002. Hashimoto, K., I. Nagano, M. Yamamoto, T. Okada, I. Kimura, H. Matsumoto, and H. Oki, EXOS-D (Akebono) very low frequency plasma wave instruments (VLF), IEEE Trans. Geoelectr. Remote Sens., 35, 278-286, 1997. Kimura, I., K. Hashimoto, I. Nagano, T. Okada, M. Yamamoto, T. Yoshino, H. Matsumoto, M. Ejiri, and K. Hayashi, VLF observations by the Akebono (EXOS-D) satellite, J. Geomagn. Geoelectr., 42, 459-478, 1990. Kasahara, Y., I. Nagano, and I. Kimura, Global dynamics of the inner magnetosphere derived from ELF/VLF waves observed by Akebono, Conference Digest of the 2001 Asia-Pacific Radio Science Conferece, , H5-04, 2001. Li, X, D. N. Baker, M. Temerin, T. E. Cayton, G. D. Reeves, R. A. Christensen, J. B. Blake, M. D. Looper, R. Nakamura, and S. G. Kanekal, Multisatellite observations of the outer zone electron variation during the November 3-4, 1993, magnetic storm, J. Geophys. Res., 102, 14,123-14,140, 1997. Li, X, and M. Temerin, The electron radiation belt, Space Sci. Rev., 95, 569-580, 2001. Meredith, N. P., R. B. Home, D. Summers, R. M. Thorne, R. H. A. lies, D. Heynderickx, and R. R. Anderson, Evidence for acceleration of outer zone electrons to relativistic energies by whistler mode chorus, Annales Geophysicae, 20, 967-979, 2002. Miyoshi, Y., A. Morioka, T. Obara, H. Misawa, T. Nagai, and Y. Kasahara, Rebuilding process of the outer radiation belt during the November 3, 1993, magnetic storm - NOAA and EXOS-D observations, J. Geophys. Res., 108, doi:10.1029/2002ja007542, 2003. Obara, T., T. Nagatsuma, M. Den, Y. Miyoshi, and A. Morioka, Main-phase creation of "seed" electrons in the outer radiation belt, Earth Planets Space, 52, 41-47, 2000. Storey, L. R. 0., and F. Lefeuvre, The analysis of 6-component measurements of a random electromagnetic wave field in a magnetoplasma, 1, The direct problem, Geophys. J. R. Astron. Soc, 56, 255-269, 1979. Storey, L. R. O., and F. Lefeuvre, The analysis of 6-component measurements of a random electromagnetic wave field in a magnetoplasma, 2, The integration kernels, Geophys. J. R. Astron. Soc, 62, 173-194, 1980. Summers, D., R. M. Thorne, and F. Xiao, Relativistic theory of wave-particle resonant diffusion with application to electron acceleration in the magnetosphere, J. Geophys. Res., 103, 20,487-20,500, 1998. -233-
Summers, D., and C. Ma, A model for generating relativistic electrons in the Earth's inner magnetosphere based on gyroresonant wave-particle interactions, J. Geophys. Res., 105, 2625-2639, 2000. Summers, D., C. Ma, N. P. Meredith, R. B. Home, R. M. Thorne, D. Heynderickx, and R. R. Anderson, Model of the energization of outer-zone electrons by whistler-mode chorus during the October 9, 1990 geomagnetic storm, Geophys. Res. Lett, 29, doi:10.1029/2002GL016039, 2002.
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MAGNETOSPHERIC ACTIVE WAVE MEASUREMENTS P. Song, B. W. Reinisch, and X. Huang Center for Atmospheric Research, Department of Environmental, Earth and Atmospheric Sciences, University of Massachusetts Lowell, 600 Suffolk Street, Lowell, Massachusetts, USA J. L. Green Goddard Space Flight Center, NASA, Greenbelt, Maryland, USA ABSTRACT The magnetospheric community has conducted active wave measurements, which are made by radiating known signals to space and measuring the dispersion of the signals when propagating through magnetospheric plasma, for more than three decades. Because of the advances in space electronics and signal processing technology, active wave experiments can now cover a large dynamic range in frequency and be made in a large spatial range for versatile applications. They are becoming a promising new technique to probe the space plasma conditions. In this review we demonstrate the capability of magnetospheric sounding technique employed by the radio plasma imager (RPI) on board the IMAGE satellite. This new technique, combined with the mathematical density inversion algorithm, measures the plasma density in situ at the satellite location and remotely and instantaneously along the magnetic field line from one hemisphere to the other down to as low as half an Earth radius (Re) in altitude. The technique has been formally validated. The database from the RPI active measurements covers all local times from 1.5 Re to 5 Re under different geomagnetic activities. Empirical models that are being developed specify the density as functions of radial distance, latitude, local time, distance along the field line from the earth's surface, solar wind conditions, geomagnetic indices, and other possible variables that affect the density distribution. The models can describe the statistical behaviors of the plasma distribution and can also provide snapshots of the plasma conditions on occasions. Dynamical processes that cause variations from the average models, such as depletion/refilling processes and plasma convection tail formation, can be studied. When applied to multiple satellites, or a constellation of satellites, the active wave measurements also make it possible for magnetospheric tomography, in which transmissions and reception of various waves within the constellation are used to derive the plasma density and magnetic field component in the constellation plane. 1. INTRODUCTION Conventional magnetospheric measurements of the electron number density Ne are made in situ on spaceborne platforms. These measurements provide Ne at each location at different times along the satellite trajectory. To interpret a measured Ne time series as the Ne distribution along the trajectory, one assumes that the plasma is in steady state. Often this assumption may not be valid in many regions in the magnetosphere because plasma structures may move rapidly. Furthermore, many of the plasma detectors cover only a limited energy range. In many cases, a significant portion of the plasma at lowest energies cannot be measured. At the lowest energies, the spacecraft potential becomes important to the total measured flux. Therefore, measuring Ne accurately by in situ particle detectors is not trivial [e.g., Song and Russell, 1999]. The Ne can also be derived from the changes in the characteristics of the waves near the plasma frequency or upper hybrid frequency. This method has been confirmed in some cases by comparison with other plasma density measurement methods [e.g., Carpenter et al., 1981; Song et al., 1993]. Furthermore, Ne can be derived through active wave measurements, in which known signals are transmitted into the space plasma and the dispersion or phase/time delays of the signals at the receivers are measured [Harvey et al., 1978; Etcheto et al., 1981; Oya et al., 1987, 1990; Perraut et al., 1990; Song et al., 1993; Kurth et al., 2001; Decreau et al., 2001]. This paper focuses on recent results from the sounding method used by the radio plasma imager (RPI) [Reinisch et al., 2000] on the IMAGE satellite.
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At the end of the paper, we further explore the possibility of magnetospheric tomography, a new exciting possibility of global imaging technique using a constellation of satellites. Physically, the plasma distribution is affected by magnetospheric dynamic processes and global convection in addition to the variations in the ionospheric sources. Simultaneous measurements of the plasma distribution is a large spatial region can help us to understand the mechanisms that determine the distribution. In the magnetosphere, the electromagnetic force dominates the plasma distribution processes. The Ne profile may be determined by the loss/refilling and other dynamic processes, and hence may vary substantially according to its stage relative to a storm/substorm onset. In the inner magnetosphere, transient structures, such as shoulders, fingers, "bite-outs", troughs, "donkey ears", and plasma convection tails [Grebowsky, 1970; Chen et al., 1975; Oya, 1991; Sandel et al, 2001; Burch et al., 2001b; Goldstein et al., 2002] can also be formed. Therefore, instantaneous or simultaneously measurements of Ne from regions other than the single location at a satellite are critical to the understanding of magnetospheric dynamic processes. In situ measurements, in order to make such measurements, are limited by the number of satellites and to the locations of the satellites. We introduce and demonstrate a new technique that instantaneously and remotely measures Ne along a magnetic field line from one hemisphere to the other. As the satellite flies through different field lines, a 2-D snapshot of the global Ne distribution can be taken during each orbit. The variations of the distribution can then be studied. Before discussing the active wave measurements, it is useful to briefly review the cold plasma wave theory because it is the foundation of the method to be discussed below. Figure 1 shows the cold plasma dispersion relations when the ion effects are neglected [Budden, 1988]. The vertical-axis is the square of the index of refraction, n = c/v where v and c are the phase speed of the wave and the speed of light in free space, respectively. The wave cannot propagate when n is negative and it reflects at n = 0, at which the corresponding frequency is referred to as the cutoff frequency where the phase velocity goes to infinity and the group velocity goes to zero. From left to right the three cutoff frequencies, where the shaded regions cross the x-axis, are referred to as the X, O, and Z cutoffs, respectively. The three shaded regions above the x-axis from left to right are known as X, O, and Z modes, respectively. Each shaded region indicates the range of the propagation angle 9 (angle between B and the wave normal) from parallel (0 = 0) to perpendicular propagation (9 = 7t/2) as labeled. The quantity plotted on the horizontal axis is the squared ratio of the plasma frequency to the wave frequency. At a fixed point in space, say at the satellite location, the scan with an increasing frequency is from the right to the left as labeled at the bottom of the figure. The presentation of Figure 1 is particularly convenient when discussing wave propagation in a spatially varying plasma. For waves with a fixed frequency, Ne increases (decreases) along the curve to the right (left), also labeled at the bottom of the figure. Since the reflection point (n = 0) is at the right-hand side of the propagation region of each curve which is the upper portion, echoes can be produced only after being reflected from the region of a higher Ne than at the satellite location. The echoes can be received at the source location only if the gradient in the Fig. 1. Cold plasma dispersion relation without collisions index of refraction is parallel to the incident wave at [Budden, 1988]. the reflection point. One important feature in this
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figure is that for parallel propagations (9 = 0) the O and Z modes intercept at point P. For nearly parallel propagation, it is possible for a mode conversion to occur near this point. Namely, the wave "jumps" from one curve to the other [Oya, 1971; Jones, 1981; Budden, 1988]. When an O mode wave propagates into a higher Ne region, it can convert to the Z mode near point P and then propagate into the region of even higher Ne. After reflection at the Z mode cutoff point, the wave is converted back to the O mode again at point P. The echo trace formed by this particular mode coupling process has been referred to as the "Ztrace" [Budden, 1988]. To avoid confusion, we will refer to it as the O-Z-0 mode. In summary, when a wave propagates in the direction of increasing Ne, the wave is reflected at places where the wave frequency equals the cutoff frequency of the wave mode. The cutoff frequency of a given mode is, in general, a function of Ne and field strength B. Therefore, from the frequency of an echo the plasma conditions at the reflection point can be determined, and from the echo time delays of the signals with various frequencies the distance to the reflection points can be calculated. 2. IMAGE ACTIVE WAVE MEASUREMENTS 2.1. RPI Remote Measurement Technique The radio plasma imager (RPI) [Reinisch et al., 2000] on IMAGE, which is on a polar orbit with an apogee altitude of ~7 Re [Burch et al., 2001a], applies the radio sounding technique to measure Ne remotely. The sounding method has been used extensively in ionospheric Ne measurements. The RPI uses three orthogonal thin-wire dipole antennas. The length of the two dipoles in the spin plane is nominally 500 meters tip-to-tip, and the dipole along the spin axis is 20 meters long. The spin period of the satellite is about 2 minutes. In the active measurement, the RPI transmits coded signals from 3 kHz to 3 MHz in a broad dipole pattern and measures the delay time of returned signals called echoes as a function of frequency. Echoes that experience the same dispersion in the propagation form a distinct trace in the "plasmagram", a display of signal amplitude as functions of frequency and echo delay time as shown in Figure 2a. However, the exact mode of a trace is unknown yet. Since there are only a limited number of modes, we can test them each individually and compare with the observation. Combining the measured dispersion and the assumed theoretical dispersion relations, the plasma conditions in the regions where the waves propagate, can be derived. If using the derived plasma conditions and a different mode another observed trace is reproduced, the plasma conditions and the two modes used in the calculations are confirmed.
Fig. 2a. Plasmagram at 1750 UT, Oct. 24, 2000, showing color-coded signal amplitude in dB as function of frequency and virtual range [Reinisch et al., 2001b]. The virtual range is the time delay of an echo multiplied by one half the speed of light.
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In the example shown in Figure 2a, the RPI transmitted a pulse and listened for echoes, then repeat it one more time for improved signal-to-noise ratios. This process was taken at each of 114 frequencies, which were logarithmically spaced, producing a plasmagram. A complete measurement is made in near one spin period. The time resolution for the echo delay time is 3.2 ms (given by the pulse width) translating to -500 km virtual range. In addition to the active plasma wave measurement, passive measurements are also made in the same way as those of conventional plasma wave instruments, when RPI alternates between two active measurements, each of which produce a plasmagram like that in Figure 2a [Reinisch et al., 2001b].
Although the RPI radio sounding, in principle, can receive echoes from all directions [Green et al., 2000], the most distinct long-range echoes arrive along the magnetic field lines [Reinisch et al., 2001a], Although field-aligned HF propagation had previously been reported for the topside sounder observations [Muldrew, 1963; Dyson and Benson, 1978], the possibility of long-range propagations over 6 Re was beyond expectation. This unexpected phenomenon becomes a unique opportunity. When IMAGE is in the polar region where the field line is nearly radial, the Ne profile in altitude can be derived [Reinisch et al., 2001a]. When IMAGE is in the equatorial region, field-aligned Ne profiles for both hemispheres can be derived [Reinisch et al., 2001b]. Below, we discuss the example shown in Figure 2a when IMAGE was in the equatorial region. We discuss in detail how the Ne profile along the magnetic field from one hemisphere to the other can be derived with the observed echo traces. In Figure 2a, the color-coded signal amplitudes are shown as functions of frequency and time delay, which is multiplied by cfl and presented as the so-called virtual range. IMAGE was at 22.8 magnetic local time, -11.2° magnetic latitude, and L=3. The most striking echo traces are the two at middle and long ranges between 300 and 600 kHz. The wave normal direction of the arriving signals can be determined from the amplitude and phase measurements on the three antennas [Reinisch et al., 1999]. The phase measurements of these two traces, which will be referred to as the primary traces, show that the echoes are arriving along the model magnetic field line through the spacecraft location. Notice that for waves of frequencies just above the plasma frequency the group velocity can be significantly smaller than c, but with increasing frequency it gets close to c. The shorter virtual echo ranges in a trace at higher frequencies actually correspond to longer real distances than those at lower frequencies in the same trace. The vertical lines near 35, 40, 70, and 140 kHz in Figure 2a result from electrostatic echoes and resonances of the local plasma. IMAGE does not carry a magnetometer to measure the electron gyrofrequency. Based on a magnetic field model [Tsyganenko and Stern, 1996], the local electron gyrofrequency is 34.93 kHz. Therefore, the observed resonances at 35, 70, and 140 kHz are the first, second and fourth harmonics of the electron gyrofrequency, respectively. The missing third harmonic may be caused by the gap in the frequency stepping. The measured electron gyrofrequency is used in the Ne profile inversion procedure. 2.2. Inversion of Plasmagrams to Plasma Density Profiles Plasma density inversion techniques from echo traces have been developed in radio science and applied to ionospheric physics [e.g., Jackson et al., 1980; Huang and Reinisch, 1982]. They are based on the radio wave propagation theory in a cold plasma as shown in Figure 1. If one starts with the plasma density and field measurements immediately next to the satellite, assuming a mode and a propagation direction relative to the magnetic field, in principle, the group velocity of the wave can be determined from the dispersion relations. Using the calculated group velocity and the measured time delay of the echo, the distance of the reflection point is determined. The plasma density can be determined by the reflection condition, namely, the cutoff frequency for the chosen mode, assuming a global magnetic field model. This
Fig. 2b. Reconstruction of the echo traces, black lines, using the inverted density [Reinisch et al., 2001b]. The color code shows the enhanced observed traces in Figure 2a. Labels S, N, X, and Z denote southern and northern hemispheres and X and O-Z-0 traces, respectively. The NX trace has been extrapolated to higher freauencies.
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procedure can be continuously applied to points further away from the satellite, i.e. including echoes at increasingly higher frequency, and hence the Ne -distance profile can be derived. In practice, our inversion is performed as a best fit to an observed echo trace, solving an integral equation. Two assumptions are needed for an inversion: the propagation mode and the propagation angle. For a given echo trace, one may try the dispersion for each of the three modes shown in Figure 1, and also look at possible mode coupling mechanisms as discussed at the end of the introduction. There are two possible outcomes for each try. First, using the selected mode no Ne profile can be calculated that reproduces the observed trace. This means that the chosen mode cannot propagate in the given direction and is a wrong choice. Second, the calculation provides a Ne -location profile that reproduces the trace. This Ne profile is then used to calculate the expected echo traces for other modes. If these calculated traces match the observations, then the correct profile has been determined. For the inversion, two initial values — Ne and B at the satellite location ~ are needed. The Ne near the satellite is measured in situ by the resonance at the plasma frequency or by the extrapolation of the echo trace. The latter is used for our case. The global magnetic field is determined according to an empirical model [Tsyganenko and Stern, 1996] and verified by the in situ electron gyrofrequency measurements as discussed above. The echo traces (echo delay as function of frequency) in Figure 2a, global magnetic field model, and initial values are input into the inversion program to test for various possible modes and possible propagation angles. Assuming the two strongest traces in Figure 2 are X-modes propagating parallel and anti-parallel, respectively, to the magnetic field, the Ne distribution along the field is shown in Figure 3. Using the Ne profile shown in Figure 3, the recalculated traces, solid black lines in Figure 2b, are labeled SX and NX, with the NX trace consisting of X-mode echoes from the northern hemisphere, and the SX trace of X-mode echoes from the southern hemisphere. This is consistent with a picture in which Ne is approximately symmetric about its lowest value at the equator and an IMAGE location in the southern hemisphere. Here we should point out that no echo is reflected from the equatorial region where Ne is lower than that at the satellite although the total electron content in this region is measured by the virtual range jump in the start frequency of the NX trace. The Ne in this region is derived by assuming a continuous N e gradient, namely a smooth Ne profile constrained by the total electron content. Because the total electron content in this region is very small compared with the total electron content at higher latitudes, the Ne profile at higher latitudes in the NX trace, nevertheless, is not affected by the assumption made to interpolate Ne near the equator. The successful reproduction of the X traces does not provide a proof of the derived Ne profile to be correct. The proof is given in the following procedure. Using the Ne profile in Figure 3 derived from the primary traces, SX and NX, we calculate the OZ-0 traces using the same propagation direction and plot them in Figure 2b labeled SZ and NZ, respectively. The OZ-0 traces are formed by waves that undergo two mode conversions, as discussed in the introduction. Mode Fig. 3. Density distribution along the field line for the case shown in Figure 2, (a) as coupling is most efficient function of latitude with height marks, and (b) as functions of the magnetic field strength. The dashed-line segment in (a) is based on extrapolation of the NX trace. The solid (dashed) line in (b) is for the southern (northern) hemisphere [Reinisch et al., 2001b].
,
Wtien wave
n •
1S
n T^ - lfle cannot
Sma11
power
completely be converted in the mode conversion and some wave power is lost. Therefore the amplitudes of the O-Z-O traces
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should be smaller than that of the X traces. The calculated SZ and NZ traces in Figure 2b match the observed weak-amplitude traces very well, supporting the derived Ne profile and correct identification of the mode. We should point out that the density profile shown in Figure 3 and the trace reproductions shown in Figure 2b were obtained after extensive tests of various modes, mode conversions, and propagation angles. Other possibilities which are not shown cannot reproduce the whole set of the observed traces. The differences are quite obvious by any standards, in the shape of the trace and in the level of the virtual range. Figure 3 also shows the calculated Ne distribution along the field as a function of strength of the magnetic field. Since the measurements of the echo traces were made in about 20 s, the Ne profile can, for most purposes, be considered as being obtained instantaneously. The satellite traveled only about 100 km in that time. Therefore the measured profile represents the flux tube with a radius of less than 50 km near the equator. The profile function in Figure 3a is generally consistent with the statistical observational results at low and mid latitudes by Gallagher et al. [2000]. The Ne increases relatively slowly near the equator and much faster at higher latitudes. 2.3. 2-D Plasmaspheric Density Distributions As we demonstrated in the last subsection, a field-aligned Ne distribution can be derived with measurements made in about 20 seconds every 2 minutes. As IMAGE flies through regions of different L-shell, multiple such Ne distributions can be measured and hence a 2-D snapshot of the magnetospheric Ne distribution can be derived. Figure 4 shows an example when 6 field-aligned Ne distributions are measured in 20 min in the morning sector, noting that not every plasmagram contains identifiable traces. After exhausted tests of different functional forms, we derive an empirical expression of the Ne distribution based on the requirements for parallel propagation: Fig. 4. In an IMAGE pass, six complete measurements were made in 20 min near 0800 LT. The solid lines show the inverted density and the dashed lines the best fit to an empirical expression
N =Nn\l
+ K— 1
derived based on the requirements for parallel propagation. The L-shell values, where measurements were made, are from top 2.2,2.4,2.5, 2 7 2 9 and 3 0
cos n \
9
3
where X and Xmw, are the magnetic latitude and magnetic invariant latitude of a point in space, respectively. In . . , , ., , .. fi e , TU principle, both Neo and n are functions of L. The equatorial densities shown in Figure 4 indicate a 1/L dependence. The factor n is only weakly dependent on L. The dashed lines in Figure 4 show the model prediction based on the above equation for each corresponding L-value of the measurements. The colorcoded Ne on the right-hand-side of the Earth in Figure 5 shows the empirical model for this snapshot. 2.4. Polar Region Density Distributions The Ne distribution in the polar region is different from that in the plasmasphere because the field lines in this region may be either open or have not been closed for a long time. In contrast, the plasmasphere consists of field lines that have convected or corotated around the Earth for some time. Examination of a large number of events indicates that Ne is highly dependent on the radial distance R and then the geomagnetic activity. In order to determine the functional form of each dependence, the data is first binned according to the radial distance as shown in panel 6a. The linear relationship in the log-log scale between Ne and radial distance indicates a power-law relationship with a power index near - 5 . The observed Ne is then normalized by this power law to remove the radial distance dependence and then
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binned according to the Kp index as shown in panel 6b. The relationship between the normalized electron density and Kp is relatively linear in a linear-log scale, indicating an exponential relationship. Then, the radial distance and Kp dependences are removed by a similar normalization. The binning of the normalized data according to invariant latitude X.;nv shows a weak decrease near the pole. The polar region empirical density model in this database [Nsumei et al., 2003] can be expressed as N.{R,Kp,Xlm) Fig. 5. Plasma density distribution derived from RPI measurements. Low latitude region density is derived from a single snapshot made in 20 min on June 8, 2001. The polar region density is based on a statistical study for Kp= 4 [Nsumeietal, 2003].
= N0Rre!'Kl**<- .
A tri-variant fit of radial distance, Kp, and invariant latitude gives Ne(R,Kp,Zim.) = 68510/r4-9V°-220*"-0-038^) The color-coded density distribution above the Earth in Figure 5 shows this empirical model
under the average Kp (=4) condition. When comparing this model with previous empirical models [Persoon et al., 1982, Gallagher et al., 2000], which were developed based on in situ passive wave and some particle measurements, the densities in the previous models are consistent with those for the lowest Kp cases in the RPI model. This difference may be due to the whistler-mode wave cutoff technique employed by Persoon et al. [1982], In their passive measurements, the plasma density is determined by the plasma frequency "cutoff. This plasma frequency cutoff occurs only when the plasma frequency is lower than the electron gyrofrequency. Otherwise, the cutoff occurs at the electron gyrofrequency. Observationally, the two cutoffs can be relatively easily distinguished by more fluctuations in the plasma frequency cutoff and a less fluctuating electron gyrofrequency cutoff that follows the electron gyrofrequency predicted by a magnetic field model. Therefore, Persoon et al.'s method may only produce reliable Ne values when the electron plasma
Fig. 6. Polar plasma densities, (a) Density binned as function of radial distance, (b) Normalized density, by removing the radial distance dependence, binned as function of Kp. (c) Normalized density, by removing the radial distance and Kp dependence, binned as function magnetic invariant latitude [Nsumei et al., 2003].
frequency is less than the electron gyrofrequency. As shown in Figure 6b, the density is higher for higher
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Kp. Thus the results of Persoon et al. [1982], which were also incorporated in the Gallagher et al. [2000] model, would likely have excluded high Kp cases when Ne is so high that the plasma frequency is greater than the electron gyrofrequency. Furthermore, when the effects of finite source regions and the possible oblique propagation are considered, the cutoff frequency used in the passive measurements is an even lower underestimate of Ne [Nsumei et al., 2003], 2.5. Local Resonance Frequencies. One interesting phenomenon in active measurements is enhanced power at local plasma characteristic frequencies, as shown in Figure 7. From the empirical magnetospheric field model [Tsyganenko and Stern, 1996], the electron gyrofrequency is 19.5 kHz. Clear enhancements occur at 19.5 and 39 kHz and correspond to the first and second harmonics of the electron gyrofrequency. In addition, an enhancement occurs at 22.4 kHz. It corresponds to the plasma frequency fpe, as can be verified by the enhancement at 29.7 kHz which corresponds to the upper hybrid frequency fuh = (fee2 + fpe2)1'2Furthermore, using the above values, the X-cutoff frequency fx = [fce+ (J-*2 + 4fpe2)1/2]/2, which corresponds to the positive branch of X=l-Y at X =0 in Figure 1, is 34.2 kHz at the satellite location. The beginning of the X-trace is at 0.3 Re from the satellite. The projected cutoff at zero virtual range for the curved trace in Figure 7 is consistent with the calculated X-cutoff. Fig. 7. Active wave measurement showing resonance at characteristic frequencies, fce, 2f«, £_ and fuh.
_.
„„.
T h e R P 1 Can
, . . ... P«>gram the transmissions With even
finer frequency resolution and broader frequency range, while reducing the virtual range or listening time, for local resonance analyses. This technique is being used for extensive studies. Benson et al. [2002] reported observations of many different modes and harmonics. 2.6. Direct Echoes From Plasma Boundaries In sections 2.1-2.4, we have focused on fieldaligned propagation modes. Transmitted signals propagate in all directions and each signal can be decomposed into orthogonal polarizations and propagate as different modes. It is possible for a mode to radially propagate and reflect back. If Ne varies relatively slowly in space, echoes at the same frequency may be reflected at different distances in space forming rather diffused echo "trace" as can be seen in Figure 2a below the discreet trace in virtual ranges around 3 Re in the frequency range between 300 and 400 kHz. However, if at a plasma boundary where Ne Fig. 8. An example of direct echo traces that are reflected at changes rapidly in space, sharp echo traces can shar P d e n s i t v gradients, such as the magnetopause and also be observed. These traces are flat on a P lasma P ause plasmagram because it is formed by a large change in Ne in a relatively small spatial distance range. Figure 8 shows an event when traces were formed by echoes from the magnetopause and the plasmapause when IMAGE was near 7 Re radial distance, descending from apogee. Green and Reinisch [2002] provided detailed accounts for this event.
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3. MAGNETOSPHERIC TOMOGRAPHY 3.1. Phase and Faraday Rotation Measurements For a constellation of satellites, each satellite transmission can be received by other satellites in the constellation. As shown in Figure 1, the phase velocity of a wave is function of Ne and frequency. If two frequencies are transmitted at the same time, the phase difference between the two signals at one of the receivers is [Davies, 1990]
° = — {^-%fiAS^.uU\-f^if^[fl] c
\J\
Jl )
J\
where /i,2 and S are the two transmission frequencies and the separation between the transmitting and receiving satellites, respectively. The brackets indicate ray averages and S = 2 Re was used to obtain the last expression (frequency in kHz). When <£ is measured and/i a n d ^ are known, the average plasma frequency and hence the density between the two satellites can be determined. However, because the phase difference has a 27i-ambiguity, an additional frequency, which provides two more phase difference measurements but adds one more equation, is required to make an absolute Ne measurement. The magnetic field in the plasma will produce the Faraday rotation in the polarization direction. The Faraday rotation is used to measure the magnetic field component in the constellation plane [Reinisch et al, 1999] eF = ^SUlML
cos^] = 0 . 5 4 ^ [ / M cos>F]
where *F is the angle between the ray and the magnetic field. The average plasma frequency can be taken from the phase difference measurement if the Faraday rotation and the phase difference measurements are made nearly at the same time. Similar to the phase difference, the Faraday rotation has a 7i-ambiguity (since *P is defined between 0 and n/2, the ambiguity for 9p is n and not 2n) and requires an additional frequency to resolve. The phase difference method has been employed in the ISEE [Harvay et al., 1979; Song et al., 1993]. The long range Faraday rotation has been measured between the IMAGE and Wind satellites [Cummer et al., 2000]. 3.2. Tomographic Inversion In a constellation, mutual transmission and reception among satellites provide Ne measurements along m(m-l)/2 independent paths, where m is the number of satellites. From these measurements, the underlying Ne distribution in the constellation plane can be inverted using the tomography method which is widely used in medical imaging. Therefore, largescale plasma distribution can be imaged. Several groups [Ganguly et al., 2000; Ergun et al., 2000; Green, 2001] have explored the possibilities to take tomographic images of the magnetosphere. The power of transmitters and the sensitivity of receivers may not be a big challenge for a tomography mission with today's digital technologies, as we have shown
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in this review that although RPI transmits only 10 W, its echoes can be easily received more than 5 Re away, translating to more than 10 Re in distance. Note that some signals loss may occur at reflection. Signal-to-noise ratio can be improved by transmitting longer pulses. The spatial resolution of the tomography is proportional to the number of independent paths and hence is proportional to the square of the number of the satellites. It decreases with an increase in the area covered in tomography. A coorbiting constellation of 7 satellites may provide about 1 Re resolution in regions separated by as much as 10 Re, as shown in Figure 9. The large dynamic range of Ne, from magnetospheric cavity of 10"1 cm"3 to magnetosheath of 100 cm"3, and the range of satellite separation, from few Re to more than 10 Re, may require more than three frequencies to resolve the n- or 27r-ambiguiry. Quantitative assessments on the above issues can be found in Green [2001]. Simultaneous transmissions of multiple frequencies may be challenging. Ergun et al. [2000] proposed to use the fundamental and third harmonic of a square-wave as two probing frequencies. The frequency separation in this scenario may be too large to be useful in magnetospheric tomography. 4. SUMMARY AND COMMENTS Active wave measurements have brought some revolutionary methods to space measurements. These measurements include remote sensing plasma boundaries, remote measurements of plasma density along the magnetic field, local resonance frequencies, and possible magnetospheric tomography. Among these, the field-aligned Ne measurements have provides tremendous scientific advances in understanding ionosphere-magnetosphere coupling. An empirical model magnetosphere density from 1.5 Re to 5 Re is under active construction [Nsumei et al., 2002; Reinisch et al., 2003; Huang et al., 2003]. The IMAGE RPI instrument has been based on technologies developed from ground-based ionospheric sounding and low altitude ionospheric topside sounding. The hardware, instrument operation software, and data process software have been tested extensively. However, transplant of these technologies to magnetospheric explorations reveals many new scientific challenges. For example, the sounding technique probes only locations with higher densities than that at the sounder's location. In ionospheric sounding, the high-density regions are relatively easy to determine and the spatial scales of the regions are relatively small. In magnetospheric sounding, on the other hand, the satellite(s) flies through a large density range. Very often, high-density regions can be in any direction to the satellite. Furthermore there exist many not well known large-scale moving plasma structures. Mysterious echo traces may be received some times. It may take years of experience before all observations can be fully understood. These pose both challenges and opportunities. The power efficiency of active wave measurements has been improved enormously through digital technology and coding technology. Signals from a 10-W transmitter can be easily received from 10 Re away with normal magnetospheric background wave activities. IMAGE-RPI has employed very long antennas, 500 m each, in order to improve power efficiency at lower frequencies and to be able to measure the phase differences of longer wavelengths. While IMAGE-RPI was successful to deploy such long antennas, which is very challenging in engineering, shorter antennas may be sufficient for many missions of different scientific objectives. For example, in a tomography mission, the antennas can be much shorter than that of IMAGE-RPI. Among major unknowns in RPI observations are the physical reasons why the echo-signals are likely to propagate along the magnetic field. Presently, there are two possible answers. First, there may be plasma irregularities along the magnetic field forming magnetic cavities. Signals could be trapped in the cavities and propagate along one of them in each transmission. Second, in a large spatial range, the large-scale plasma and field gradients are more along the field leading to the contours of refraction index to be perpendicular to the field. A normal incident wave will then propagate along the field according to SnelPs law of refraction. Each of the two possibilities can explain many observed features but with some others unexplained. Major undertakings are required to understand this problem.
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ACKNOWLEDGMENTS The presented paper summarizes the works done by, or in collaborations with, the following members of the talented and supporting RPI team: R. Benson, D. Carpenter, S. Fung, I. Galken, G. Khmyrov, and P. Nsumei. The work was supported by NASA under subcontract 83822 from Southwest Research Institute, and by NSF under Awards NSF-ATM9729775 and NSF-ATM0077655. REFERENCES Benson, R. F., V. A. Osherovich, J. Fainberg, and B. W. Reinisch, J. Geophys. Res., submitted, 2002 Budden, K. G., The propagation of Radio Waves, 669 pp., Cambridge Univ. Press, New York, 1988. Burch, J.L. et al., Views of Earth's magnetosphere with the IMAGE satellite, Science, 291, 541, 2001a. Burch, J. L., D. G. Mitchell, B. R. Sandel, P. C. Brandt, and M. Wuest, Global dynamics of the plasmasphere and ring current during magnetic storms, Geophys. Res. Lett., 28, 1159, 2001b. Carpenter, D. L., R. R. Anderson, T. F. Bell, and T. R. Miller, A comparison of equatorial electron density measured by whistler and by satellite radio techniques, Geophys. Res. Lett., 8, 1107, 1981. Chen, A. J., J. M. Grebowsky, and H. A. Taylor, Jr., Dynamics of mid-latitude ion trough and plasma tail, J. Geophys. Res., 80, 968, 1975. Cummer, S. A., M. J. Rener, B. W. Reinisch, M. L. Kaiser, J. L. Green, R. F. Benson, R. Manning, and K. Goetz, A test of magnetospheric radio tomographic imaging with IMAGE and WIND, Geophys. Res. Letts. 28, 1131, 2001. Davis, K., Ionospheric Radio, Peter Peregrinus Ltd., London, UK, 1990. Decreau, et al., Early results from the Whisper instrument on Cluster: an overview, Ann. Geophysicae, 19, 1241-1258, 2001. Dyson, P. L., and R. F. Benson, Topside sounder observations of equatorial bobbles, Geophys. Res. Lett., 5, 795, 1978. Ergun, R. E., et al., Feasibility of a multisatellite investigation of the Earth's magnetosphere with radio tomography, /. Geophys. Res., 105, 361, 2000. Etcheto, J., H.d. Feraudy, and J.G. Trotignon, Plasma resonance stimulation in space plasmas, Adv. Space Res., 1, 183-196, 1981. Gallagher, D.L., P. D. Craven, and R. H. Comfort, Global core plasma model, J. Geophys. Res., 105, 18,819, 2000. Ganguly, S., G. H. Van Bavel, and A. Brown, Imaging electron density and magnetic field distributions in the magnetosphere: A new technique, J. Geophys. Res., 105, 16063, 2000. Goldstein, J., M. Spasojevic, P. H. Reiff, B. R. Sandel, W. T. Forrester, D. L. Gallagher, and B. W. Reinisch, Identifying the plasmapause in IMAGE EUV data using IMAGE RPI in situ steep density gradients, J. Geophys. Res., submitted, 2002. Grebowsky, J. M., Model study of plasmapause motion, J. Geophys Res., 75, 4329, 1970. Green, J. L., Magnetospheric Imaging Constellation (MAGIC) Medium-class Explorer (MIDEX) Proposal, AO-01-OSS-03, 2001. Green, J. L., and B. W. Reinisch, An overview of results from RPI on IMAGE, Space Science Reviews, accepted, 2002. Green, J. L., et al., Radio plasma imager simulations and measurements, Space Sci. Rev., 91, 361, 2000. Harvey, C. C, J. Etcheto, Y. De Javel, R. Manning, and M. Petit, The ISEE electron density experiment, IEEE Trans. Geosci. Electr., GE-16, 231, 1978. Huang, X. and B. W. Reinisch, Automatic calculation of electron density profiles from digital ionograms. 2. True height inversion of topside ionograms with the profile-fitting method, Radio Sci., 17, 4, 837, 1982. Huang, X., B. W. Reinisch, P. Song, P. Nsumei, J. L. Green, D. L. Gallagher, Empirical models of the plasma density in the inner magnetosphere, Adv. Space Res., submitted, 2003. Jackson, J.E, E.R. Schmerling, and J.H. Whittacker, Mini-review on topside sounding, IEEE Trans. Propag., AP 28(2), 284,1980. Jones, D., Nature, 228, 225-229, 1981. Kurth, W.S., G.B. Hospodarsky, D.A. Gumett, M.L. Kaiser, J.-E. Wahlund, A. Roux, P. Canu, P. Zarka, and Y. Tokarev, An overview of observations by the Cassini radio and plasma wave investigation at Earth, J. Geophys. Res., 106 (A12), 3023930252,2001. Muldrew, D. B., Radio propagation along magnetic field-aligned sheets of ionization observed by the Alouette topside sounder, J. Geophys. Res., 68, 5355, 1963. Nsumei, P. A., X. Huang, B. W. Reinisch, P. Song, V. M. Vasyliunas, J. L. Green, S. F. Fung, R. F. Benson, and D. L. Gallagher, Electron density distribution over the northern polar region deduced from IMAGE/RPI sounding, /. Geophys. Res., 108No.A2, 10.1029/2002JA009616, 2003. Oya, H., Conversion of electrostatic plasma waves into electromagnetic waves: Numerical calculation of the dispersion relation for all wavelength. Radio Sci., 6, 1131, 1971. Oya, H., Studies on plasma and plasma waves in the plasmasphere and auroral particle acceleration region by PWS on board the EXOS-D (Akebono satellite,/ Geomag. Geoelectr., 43, Suppl, 369, 1991. Oya, H., and T. Ono, Stimulation of plasma waves in the magnetosphere using satellite JIKIKEN (EXOS-B). Part II: Plasma density across the plasmapause, J. Geomagn. Geoelectr., 39, 591-607, 1987.
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Oya, H., A. Morioka, K. Kobayashi, M. Iizima, T. Ono, H. Miyaoka, T. Okada, and T. Obara, Plasma wave observation and sounder experiments (PWS) using the Akebono (EXOS-D) satellite - Instrumentation and initial results including discovery of the high altitude equatorial plasma turbulence, J. Geomagn. Geoelectr., 42, 411-442, 1990. Perraut, S., H.d. Feraudy, A. Roux, P.M.E. Decreau, J. Paris, and L. Matson, Density measurements in key regions of the Earth's magnetosphere: Cusp and auroral region, J. Geophys. Res., 95, 5997-6014, 1990. Persoon, A. M., D. A. Gurnett, and S. D. Shawhan, Polar cap electron densities from DE 1 plasma wave observations, J. Geophys. Res., 88, 10,123, 1983. Reinisch, B. W., G. S. Sales, D. M. Haines, S. F. Fung, W. W. L. Taylor, Radio wave active Doppler imaging of space plasma structures: Arrival angle, wave polarization, and Faraday rotation measurements with the radio plasma imager, Radio Sci., 34, 1513, 1999. Reinisch, B.W. et al., The radio plasma imager investigation on the IMAGE spacecraft, Space Sci, Rev., 91, 319, 2000. Reinisch, B.W. et al., First results from the radio plasma imager in IMAGE, Geophys. Res. Lett, 28, 1167, 2001a. Reinisch, B. W., X. Huang, P. Song, G. Sales, S.F. Fung, J.L. Green, D. L. Gallagher, and V. M. Vasyliunas, Plasma Density Distribution Along the Magnetospheric Field: RPI Observations From IMAGE, Geophys. Res. Lett, 28, 4521, 2001b. Reinisch, B. W., X. Huang, P. Song, S. F. Fung, J. L. Green, V. M. Vasyliunas, D. L. Gallagher, and W. R. Sandel, Plasmaspheric mass loss and refilling as a result of a magnetostorm, Submitted to /. Geophys. Res., 2003JA009948, 2003. Sandel, B. R., R. A. King, W. T. Forrester, D. L. Gallagher, A. L. Broadfoot, and C. C. Curtis, Initial results from the IMAGE extreme ultraviolet imager, Geophys. Res. Lett., 28, 1439, 2001. Song, P. et al., Structure and Properties of the Subsolar Magnetopause for Northward IMF: Multiple Instrument Particle Observations,/ Geophys. Res., 98, 11319, 1993. Song, P. and C. T. Russell, Time series data analyses in space physics, Space Science Reviews, 87, 387-463, 1999. Tsyganenko, N.A. and D.P. Stem, Modeling the global magnetic field and the large-scale Birkeland current systems, J.Geophys.Res., 101, 27187-27198, 1996.
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2/ p RADIO SOURCE IN GEOTAIL OBSERVATIONS AND NUMERICAL SIMULATIONS -MICROSCOPIC VIEWY.Kasaba 1 , H. Matsumoto2, Y. Omura2, and T. Muka?
'The Institute of Space andAstronautical Science (ISAS), Sagamihara,Kanagawa229-851O.Japan RadioScienceCenterfor Space andAtmosphere (RASC), Kyoto University, Uji, Kyoto 611-0011,Japan
2
ABSTRACT We have studied several topics related to the % radiation generated in the terrestrial electron foreshock. Our investigation started from the macroscopic geometry d" the radio source, and is expanding to the microscopic processes. In this paper, we present a summary of latter studies, especially about the generation mechanism of electrostatic and electromagnetic 2/p waves and the electron acceleration at the quasi-perpendicular shock.
INTRODUCTION: MACROSCOPICVIEWOF 2/ p RADIOSOURCE "2/p radiation" is frequently observed in the terrestrial upstream region at twice the electron plasma frequency. Its source is thought at the electron foreshock. Energetic electrons are accelerated at quasi-perpendicular shocks and backstreaming along the interplanetary magnetic field (IMF) lines. These electrons generate Langmuir waves by electron beam instability. Electromagnetic 2fp radiation is finally generated from these Langmuir waves. The geometry of this radio source has been estimated by two indirect methods. One is "direction finding analysis", statistics based on single spacecraft observation (Kasaba et al., 2000) and the triangulation by two spacecraft (Reiner et al., 1997). Another method is "frequency variation" (Kasaba et al., 1997). Associated with the fluctuation of the plasma density at the radio source, new lfv line at twice the new plasma frequency appears and the old 2/p line vanishes. The extension of the radio source can be estimated by the timing of these variations. These results indirectly suggest that the radio source should be the electron foreshock. Direct evidence is provided by the first global mapping of 2/p radio flux around the foreshock (Kasaba et al., 2000). Figure 1 shows that the center of 2^ flux density is superposed on the region along the IMF lines connected to the quasi-perpendicular shocks in which strong Langmuir waves and energetic electrons are confined. The 2/p radio flux distribution also leads to two questions. One is the different diffusion rate at the region distant from the shock. The decrease rate of the Langmuir wave and electron flux is smaller than that in 2/p radio flux. The other is weak 2^ radio flux in the region near the perpendicular shock. Quasi-perpendicular shock region supplies energetic electrons to the foreshock, but the region near the shock seems not to have strongest radio emissivity. These questions are related to microscopic views. The former is concerning to the generation mechanism of the radiation. The radiation near the electron plasma frequency is common in solar and stellar radio burst, but its mechanism has not been established enough. The latter is concerning to the acceleration of electron beam. Electron acceleration at quasi-perpendicular shock is common in interplanetary and interstellar shocks. The electron foreshock provides us a natural laboratory for fundamental processes related to these phenomena. GENERATIONMECHANISM Several mechanisms have been proposed for 2/p radiation [Ref: in Kasaba et al., 2001]. Strong instability (ex. Langmuir soliton) was not supported by observations, so two weak instability models are remained, 'modeconversion' and 'direct-conversion'. The former is based on the conversion from electrostatic % wave, but its
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excitation and the conversion to electromagnetic 2/p in space plasma are not established The latter is based on the conversion from Langmuir wave. It consists of two processes: Backscattering of Langmuir wave by the interaction with ion acoustic waves, and the coupling of the backscattered and primary Langmuir waves to the electromagnetic 2fv. However, the expected growth rate is too small to account for the production of 2/p waves in the foreshoek [cf. Yoon et al., 1994]. There are two observational evidences for the radiation mechanism. First is the relation between 2/p radio flux and Langmuir wave activity. We find that 2^ radio emissivity shows nonlinear relation to the foreshoek Langmuir wave activity in the diffusion in the region distant from the shock (Figure 1) and the relation with the solar wind kinetic energy flux (Kasaba et al., 2000). It suggests that the radiation mechanism is non-linear process. However, both candidate processes are non-linear, so that this cannot distinguish two models. Another evidence is typical wave spectrum in the electron foreshoek (Figure 2). At the most leading region of the foreshoek, there is enhancement like electrostatic 2fp and low-frequency ion acoustic wave associated with intense Langmuir wave (Lacombe et al., 1988; Kasaba et al., 2000). We do not have a confidence for the existence of electrostatic 2fp wave because the harmonic noise can be generated with the saturation by intense Langmuir waves. On the other hand, low-frequency electrostatic wave seems real, because different instruments aboard GEOTAIL (PWI and EFD) get them independently. Although we are not convincing yet that this wave is ion acoustic mode, these results favor the direct conversion model associated with backscattering of Langmuir wave. We have tried to generate 2/p waves in self-consistent PIC code simulations (Kasaba et al., 2001). Figure 3 are ? -k diagrams of E* and Bz in one- and two-dimensional cases. Magnetic field and electron beam is along the Xaxis. In both cases, the electron beam generates intense Langmuir wave in positive wave number (L). Langmuir wave backscattered by ion is also found in negative wave number (L'). Electrostatic fp wave (ES-2fp) appears at twice the wave number of beam-excited Langmuir wave. We find that this wave is enhanced in the early stage and independent of electromagnetic 2/p wave with slower growth rate. This result suggests that electrostatic 2fp wave is directly generated from beam-excited Langmuir waves and independent of electromagnetic 2fp radiation. On the other hand, electromagnetic $p wave (EM-2fp) is only excited in twc-dimensionalcase. The growth of electromagnetic 2^ is slow and related to the backscattering process of Langmuir wave. In our results, backward Langmuir wave is correlated to the product of beam-excited Langmuir wave and ion acoustic wave (ES-LF). And electromagnetic 2p wave is correlated to the product of beam-excited and backward Langmuir waves. They are favorable to the direct conversion based on successive two processes, the back-scattering of Langmuir wave by electron-ion coupling and the wave-wave coupling between beam-excited and backscattered Langmuir waves. Generation through the strong instability by extreme Langmuir wave is not supported in our simulation. In this work, the growth rate problem is not fully resolved yet by two problems: 1) Large beam energy generates the wave s about 2-3 orders stronger than the values observed in the foreshoek region. 2) Insufficient number of particles in cell enhanced backscattering process. They are originated from the limitation of running time and memory size. The reality of the most delicate processes related to electrostatic 2/p wave and backscattered Langmuir wave will be solved by future low-noise simulations (ex. electromagnetic 2-D Vlasov code). ELECTRONACCELERATION Electron acceleration at perpendicular shock is another problem. We started the investigation of the origin of "weak 2/p radio activity" near the quasi-perpendicular shock (Figure 1). We had three ideas for this weakness: First is the evolution of electron beam in the foreshoek. "Time of flight" and "electric field drift" effects generate and enhance 'bump' in the electron velocity space, and induce strong wave instability (cf. Fitzenreiter, 1995). Therefore, electron beam instability might not be induced enough in the region very close to the shock surface, if electrons initially ejected from the quasi-perpendicular shock do not have a clear bump yet. The second idea is "quadrupole directivity of 2^ radiation". Quadrupole pattern is predicted in the excitation modek, supported in the simulation (Kasaba et al., 2001), and suggested in some solar burst observations (Thejappa, 2002). Quadrupole pattern is not confirmed in the terrestrial foreshoek observation yet, but the weakness in the direction perpendicular to the tangential IMF line might cause observed weak flux near the shock. The third idea is the lack of accelerated electrons near perfectly perpendicular shock. The classic al electron acceleration model, "loss-cone type" electrons are produced from nearly perpendicular shock (Leroy and Mangeney, 1984). Its distribution function depends on the angle between the shock normal and magnetic field, 0B,,. Near the perfectly perpendicular shock, energy per charge increases exponentially, but electron flux drastically decreases. Product of the normalized flux and energy per charge becomes maximum at 0 B n ~ 86-87°.
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Fig. 1. The distribution of the amplitude of plasma waves and the population of energetic electrons: (a) 2fp radiation; (b) Langmuir wave; (c) electrons from 1 to 3 keV. [Spatial resolution: 2x2 RE|
Fig. 2. Electric field spectrum (top) and electron population bottom) around the electron foreshock observed by GEOTAIL on Apr. 6, 1995.
Fig. 3. ? -k diagrams of Ex and Bz in (a) oneand (b) two-dimensional cases.
Fig. 4. (left) Foreshock electric field spectrum and electron energy distribution at 5 shock crossings (Yellow arrow) observed by GEOTAIL on 5 Apr. 1995. (right) A schematic view of the electron foreshock. Red line is the region connected to quasi-perpendicular shocks with 0B n = 80°~85°.
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Therefore, energetic electrons are supplied from nearly perpendicular, but not truly perpendicular shock region. Figure 4 shows some examples of the energy distribution of foreshock electrons just after the acceleration at the shock. Left panels show electric field spectrum and electron distributions around 5 shock crossings with different 0^ angles. GEOTAIL was in the upstream of the shock in 16:55-17:14, 17:27-17:36, and 17:54(-18:00). At 3 almost perfectly perpendicular shocks with 0^ > 85° (89.4° at 16:55, 88.2° at 17:27, and 89.6° at 17:54), no Langmuir and electron enhancements are found in the upstream region just in front of the bow shock. With G^n = 80°~85° (84.5° at 17:36), narrow-band intense Langmuir wave becomes evident, and high-energy field-aligned electrons are supplied from the bow shock When 0en is less than 80° (76.8° 17:14), Langmuir wave becomes weak and wide-band. It is not shown in Figure 4, but low frequency magnetic field turbulence is also found. Right panel of Figure 4 is a schematic view of the electron foreshock derived from 46 shock crossings. In this panel, "the electron foreshock" (red area), with keV-order electrons and large amplitude Langmuir waves, is assumed in the area connected to the quasi-perpendicular shock region with 0B,, = 80°~85°. Although the upper and lower limits of 6^n is not quantitatively evaluated because the accuracy of 0 Bn determination is not enough, "no-active region" in this case appears at about severalJ?E around the perfectly perpendicular shocks. Namely, the third idea seems qualitatively enough to account for "weak 2^ radio activity" near the shock in Figure 1. If all the features of the electron acceleration are matched with the Leroy and Mangeney's model, it is the end of the electron acceleration studies at the Earth's shock. But, in the recent two-dimensional PIC simulation with curved shock (Savoini and Lembege, 2000), field-aligned component was also found in the expected loss cone distributions. It might be generated by the interaction with strong wave turbulence at the shock surface. We will do the further quantitative studies in wave and electron observations just upstream of the shock: 1) precise determination of ©a, at the strongest Langmuir wave generation, and 2) the detailed electron distribution function and wave features along the IMF lines from quasi- and perfectly perpendicular shocks.
ACKNOWLEDGMENTS We gratefully acknowledge whole GEOTAIL team for successful operations. The computer simulation was done on the KDK system at Radio Science Center for Space and Atmosphere (RASC), Kyoto University. REFERENCES Fitzenreiter, R. J.,The electron foreshock, Adv. Space Res., 15, 9, 1995. Lacombe, C , C. C. Harvey, S. Hoang, A. Mangeney, J.-L. Steinberg, and D. Burgess, ISEE observations of emission at twice the solar wind plasma frequency, Ann.Geophys., 1, 113, 1988. Leroy, M. M., and A. Mangeney, A theory of energization of solar wind electrons by the Earth's bow shock, Ann. Geophys., 2, 4, 449, 1984. Kasaba, Y., H. Matsumoto, and R. R. Anderson, GEOTAIL observation of % emission around the terrestrial foreshock region, Adv. SpaceRes., 20, 5, 699-702, 1997. Kasaba, Y., H. Matsumoto, Y. Omura, R. R. Anderson, T. Mukai, Y. Saito, T. Yamamoto, and S. Kokubun, Statistical studies of plasma waves and backstreaming electrons in the terrestrial electron foreshock observed by Geotail, J.Geophys.Res., 105, Al, 79-103, 2000. Kasaba, Y., H. Matsumoto, and Y. Omura, One- and two-dimensional simulations of electron beam instability: Generation of electrostatic and electromagnetic 2^p waves, J. Geophys.Res., 106, 18693-18711, 2001. Reiner, M. J., Y. Kasaba, M. L. Kaiser, H. Matsumoto, I. Nagano, and J.-L. Bougeret, 2/j, radio source location determined from WIND/GEOTAIL triangulation, Geophys.Res.Lett, 24, 919-22, 1997. Savoini, P., and B. Lembege, B., Two-dimensional simulations of a curved shock: self-consistent formation of the electron foreshock, J. Geophys. Res., 106, 12975,2001. Thejappa, G., private communication, 2002. Yoon, P. H., C. S. Wu, A. F.-Vinus, M. J. Reiner, J. Fainberg, and R. G. Stone, Theory of 3Mpe radiation induced by the bow shock, J. Geophys.Res., 99, 23481-23488, 1994. E-mail address of Y.Kasaba kasaba(2jstp.isas.ac.jp
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WEAK LANGMUIR TURBULENCE Peter H. Yoon Institute for Physical Science and Technology, University of Maryland, College Park, MD 20723, USA ABSTRACT Langmuir wave turbulence generated by a beam-plasma interaction has been studied since the early days of plasma physics research. In particular, mechanisms which lead to the quasi power-law spectrum for Langmuir waves and the superthermal electron tail population have been sought. Meanwhile, the generation of harmonic Langmuir waves has been known for some time, in both laboratory and computer-simulated experiments. However, the phenomenon has not been adequately explained in terms of theory, nor has it been fully characterized by means of numerical simulations. In this paper, a theory of harmonic Langmuir wave generation is put forth and tested against the Vlasov simulation results. It is found that the harmonic Langmuir mode spectra can indeed exhibit quasi power-law feature, implying a multi-scale structure in both frequency and wave number space spanning several orders of magnitude. Moreover, the new hyper-diffusion process is shown to be responsible for the generation of superthermal electron population. These findings indicate that the harmonic excitation process and nonlinear hyper-diffusion may be important in understanding the high-frequency turbulent processes in plasmas. INTRODUCTION It is often said that the turbulence is one of the unsolved problems of classical physics. By "turbulence," it is meant here as the turbulence in neutral fluids governed by Navier-Stokes equation. Most of the studies on turbulence is in this context (Kolmogorov, 1941; McComb, 1990). The turbulence in plasmas is even less well-understood. For low-frequency turbulence in magnetized plasmas, at least the macroscopic magnetohydrodynamic (MHD) equation, which is similar to Navier-Stokes equation in neutral fluids is applicable, and as such, many conceptual and theoretical tools originally developed for the fluid turbulence can be employed (e.g., Iroshnikov, 1964; Kraichnan, 1965). However, for high-frequency plasma turbulence where microscopic wave-particle interaction becomes important, the situation is quite different from the low-frequency MHD or fluid turbulence. Many physical systems exhibit turbulent behavior, which can be characterized by quasi scale-free structure associated with the fluctuations as exemplified by the power-law spectral distribution. Modern research on turbulence, which began with the pioneering work of Kolmogorov in the 1940s (Kolmogorov, 1941), stems largely from neutral fluid turbulence, and terminologies such as inertial and/or dissipation range behaviors, etc., often arise out of the context of the Navier-Stokes equation. However, many physical systems including the plasma, exhibit behaviors associated with fluctuations which are superficially similar to the fluid turbulence. Hence, the notions developed in the context of the fluid turbulence are sometimes indiscriminately applied to plasmas on the basis of gross morphological features, even though one does not really understand the underlying dynamics. A BRIEF HISTORY OF LANGMUIR TURBULENCE Serious investigations of (high-frequency) plasma turbulence can be said to have begun with the works of the scientists largely from the former Soviet Union in the early 1960s (e.g., standard monographs on the subject are, Kadomtsev, 1965; Sitenko, 1967; Vedenov, 1968; Sagdeev and Galeev, 1969; Tsytovich, 1970; Davdison, 1972; Kaplan and Tsytovich, 1973; Hasegawa, 1975; Akhiezer et al., 1975; Tsytovich, 1977a, 1977b; Melrose, 1980; Sitenko, 1982). Efforts by these pioneers, which came to be known as the plasma weak turbulence theory, continued on through the 1970s and 1980s. It should be noted that although the weak turbulence formalism is quite general, and in principle it can be applied to a wide variety of problems, in practice however, it is almost exclusively applied to the bump-on-tail (or weak beam-plasma) instability problem, which is one of the simplest plasma instabilities. Hence, the Langmuir turbulence problem became the testbed for various plasma turbulence theories. During the trail-blazing days, scientists in the West were largely following the Soviet scientists' lead, but a few made important contributions of their own. For instance, Dupree (1966, 1972), Weinstock (1969), and others, suggested the renormalized turbulence theories, which is an effort to go beyond the weak turbulence perturbation scheme and take the higher-order terms into account. (In the early days, the renormalized kinetic theories were called "the strong turbulence" theories, but we use the term "renormalized" to distinguish them from the later theory of the same name by Zakharov). However, findings from these sophisticated theories and earlier weak turbulence theories were directly challenged by results from numerical simulations, which often showed coherent nonlinear effects (such as particle trapping by large-amplitude waves) playing an important, if not the dominant, role (Dawson and Shanny, 1968; Morse and Nielson, 1969).
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Compounding the inability of the weak or renormalized turbulence theories to account for the coherent nonlinear dynamics was the fact that these theories simply could not produce quasi scale-free power-law type of spectrum associated with the Langmuir turbulence. The major reason is the dominance of linear physics in plasmas, unlike the fluids. In neutral fluids, the Navier-Stokes equation can be Fourier transformed into a form given by
(Jt + uk^jWj(k) = Mijm(k) I dk' vj(k') vm(k - k'), where ^i(k) is the Fourier component of the perturbed fluid velocity vector, v is the fluid viscosity, and My m (k) = -(i/2) P i j m (k),
Pijm(k) = km Py(k) + kj Pim(k),
Py(k) = <Sy -
kikj/k2,
is the coupling factor. In fluids, v is a small parameter, and many problems can be discussed by entirely ignoring the linear dissipation term, i/fc2. In Vlasov plasmas, on the other hand, the Fourier component of the perturbed distribution function obeys an equation of the form
(-J - <(•* - k - v)+ li\ /k(v) = - ^ k - 2 § W + !£ [dk> k'. | - U , / k _ k , - <^k), where ^>k is the perturbed electrostatic field which must be determined self-consistently from Poisson's equation, and u> = o;k + i7 k is the complex wave dispersion relation, 7 k being the Landau damping factor. Unlike the fluids, the dominant term in the plasma is the linear term, especially, the Landau damping factor. Owing to the insignificance of dissipation in fluids, the dominant-scale eddies can freely break into smaller and smaller eddies, thus creating the well-known stationary Kolmogorov fc~5/3 scaling in the inertial range. In plasmas, however, the cascading of the bump-on-tail generated Langmuir waves to shorter wavelength modes is prevented by the strong Landau damping. On the other hand, the upside of the dominance of the linear physics is that it makes the perturbation expansion (i.e., weak turbulence theory) to become quite valid. For Navier-Stokes turbulence, in contrast, as seen above, there is no dominant linear term around which a suitable perturbation expansion can be employed. As a result, the fluid turbulence problem necessitates some sort of renormalization or another at the outset. Then, in 1972, Zakharov proposed a semi-phenomenological theory of plasma turbulence, which came to be known as the strong turbulence theory. In his theory, the collapse of intense Langmuir wave packet plays the prominant role. The strong turbulence theory ignores the wave-particle effect, and is a macroscopic theory. But the beauty of the theory is that it predicts a certain power-law scaling resulting from the collapse of the wave packet (Zakharov, 1972; Galeev et al., 1975). Because of this, the attention in the community gradually shifted to the Zakharov's strong turbulence theory (see the reviews by Goldman, 1984; Robinson, 1997). However, as mentioned already, Zakharov theory ignores the microscopic wave-particle effect (e.g., Landau damping), and to this date the theory remains controversial, as various numerical simulations and experiments to confirm the theory are inconclusive (Robinson and Newman, 1990; Vyacheslavov et al., 2002). In short, by the mid 1980s, the community began to realize that the notion of plasma turbulence as nonlinear interactions of Fourier eigenmodes (weak turbulence picture) is insufficient, and/or that semi-phenomenological approach of Zakharov (strong turbulence picture) remains inconclusive. Moreover, the issue of coherent nonlinear effects remained wide open. However, various attempts were made to improve upon the early renormalized kinetic theories to bring the particle trapping effects (i.e., coherent nonlinear physics) into the picture. As a result, Adam et al. (1979) and Laval and Pesme (1983) claimed that the so-called turbulent trapping effect invalidates the quasilinear/weak turbulence theory at the lowest order. However, efforts to confirm their theoretical prediction either by means of numerical simulations or by experiments did not produce positive results (Theilhaber et al., 1987; Tsunoda et al., 1987). Approxmately a decade passed when Liang and Diamond (1993) performed a correct analysis of renormalized kinetic theory to show that the effects predicted by Adam et al. (1979) and Laval and Pesme (1983) were insignificant, thus reestablishing the validity of the original quasilinear/weak turbulence approach. However, the fact still remained that the weak turbulence theory could not account for the dominant coherent nonlinear dynamics observed in early simulations. In 1990, however, Dum (1990a,1990b,1990c) carried out detailed particle-in-cell simulations to show that the dominant particle trapping behavior observed in early simulations were partly owing to the insufficient mode resolution and small system size. He proceeded to demonstrate with his refined simulations that quasilinear/weak turbulence theories are actually quite good for certain parameter regime. Specifically, for a weak and warm beam, the weak turbulence theory provided an acceptable first-order description of the nonlinear behavior of the system. For a more recent discussions on the preponderance of incoherent versus coherent nonlinear effects in beam-plasma interactions, see the discussion by Omura et al. (1996). Notwithstanding these developments, however, the issue of the lack of mechanism in the weak turbulence theory to generate power-law turbulent spectra still remained outstanding. Another unsolved problems associated with the Langmuir turbulence is the generation of nonthermal energetic tail electron population frequently seen in experiments as well as in simulations (Levitskii and Shashurin, 1967; Dawson -252-
and Shanny, 1968; Whelan and Stenzel, 1985). Despite a number of plausible explanations put forth over the years, the essential issue of the physical mechanism responsible for the acceleration/heating of the electrons in the tail population remained unsolved.
HARMONIC LANGMUIR MODES In the mean time, over the past four decades or so, evidence from laboratory and spaceborne experiments (Apel, 1967, 1969; Malmberg and Wharton, 1969; Mizuno and Tanaka, 1972; Gentle and Lohr, 1973; Mori, 1973; Seidl et al., 1976; Boswell and Kellogg, 1983; Llobet et al., 1985; Kellogg et al., 1986) as well as from numerical computer simulations (Joyce et al., 1971; Goldstein et al., 1978; Pritchett and Dawson, 1983; Klimas, 1983, 1990; Akimoto et al., 1988; Nishikawa and Cairns, 1991; Yin et al., 1998; Vinas et al., 2000; Schriver et al., 2000; Kasaba et al., 2001) accumulated which showed that the real Langmuir turbulence involves the so-called harmonic mode generation, and that the power-law spectrum associated with the Langmuir turbulence involves these harmonic modes. Such a phenomenon cannot be accounted for on the basis of available theories, despite some early efforts (Manheimer, 1971; O'Neil et al., 1971; O'Neil and Winfrey, 1972). It seemed that the true picture of plasma turbulence involves the harmonic Langmuir modes, which are nonlinear eigenmodes of turbulent plasmas. In contrast, the traditional plasma turbulence theories assume that the basic modes are linear eigenmodes of a quiescent plasma. Through the excitation of nonlinear eigenmodes, which have no conceptual counterpart in fluid turbulence, the cascading of the primary bump-on-tail mode to shorter wavelength mode seemed to be established without suffering heavy Landau damping, and thus the power-law turbulence spectrum is established. Observations show that the peaks of the harmonic spectra form a steep power-law with an index ~ —5 or —6. According to the laboratory and simulation results, the harmonic waves are excited at multiples of the plasma frequency, u ~ nujpe, n = 2, 3, 4, • • • (here, u>pe = 47rne2/me is the square of the electron plasma frequency, n, e, and me being the ambient density, unit electric charge, and electron rest mass, respectively), and they all propagate with phase speeds roughly equal to the beam propagation speed, ui/k ~ VQ, where VQ is the average electron beam speed. This is the reason why these modes are not Landau damped, despite the fact that the wavelengths are much shorter than the primary Langmuir mode. Theoretical understanding of the harmonic generation was very limited, but according to the early theories developed in the 1970s, the harmonics were viewed as forced electrostatic perturbations generated by the interaction of trapped electrons and large-amplitude coherent Langmuir waves (Joyce et al., 1971; Klimas, 1983, 1990; Manheimer, 1971; O'Neil et al., 1971; O'Neil and Winfrey, 1972). Recently, however, an alternative view was developed by the author and his colleagues (Yoon, 2000; Gaelzer et al., 2002, 2003; Yoon et al., 2003; Umeda et al., 2003), in which these modes are considered as eigenmodes of nonlinear plasma system. This approach was prompted by recent simulations (Schriver et al., 2000; Kasaba et al., 2001) which show that harmonic modes persist even in late nonlinear phase when the coherent phase space structure is no longer apparent, and when the plasma has entered a stage which can be genuinely characterized by random-phases.
NONLINEAR DISPERSION RELATION FOR HARMONIC LANGMUIR MODES The detailed derivation of the nonlinear eigenmode solution, i.e., the harmonic mode dispersion relation, is given in the paper by Yoon et al. (2003), and thus shall not be repeated here. Instead, we briefly outline the general procedure. The starting point of the present analysis is the formal nonlinear spectral balance equation, given by Eq. (3) of the paper by Yoon (2000),
.-(i^!|+*>)/(..-./*-[«^C/M/(.-^ - {X(2>(K'|K - K')} 2 f ^ 7 = 7 ^ + T ^ T T ) J («) + *(3)(«'l - *'|«) '(«') H*)],
where K= (k,w), « ' = (k',o/), K-K' = (k - k', u - u'), and J da' = fdk'Jdw'.
In Eq. (1),
x ( 2 Vl« - «') = ^ £ - ^ - j ^ z v \ j dv{(k''gK) [(k "k/)'&t-"'fa] (3) X
( K ' | - K'\K) = \ £
^
^
+{K> K
" ~ K>)}'
J dv (k' • g K ) [(k' • g . - ^ )(k • g K / . ) + (K ~
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-K')],
(1)
are the various linear and nonlinear plasma response functions, / o (v) is the velocity distribution function for species a [normalized to unity, J dv / o (v) = 1], 1 d g " ~ w - k - v + iOdv' and the summation Yla ' s o v e r t n e particle species, with o;2a = 47rne2/mo representing the square of the plasma frequency for species a — e, % (e and i stand for the electrons and ions). The quantity I(K) is the phase-average over the square of the wave electric field, The derivation of Eq. (1) and approximate forms of the various response functions can be found in the paper by Yoon (2000). We simply mention that the desired dispersion equation is obtained from the real part of Eq. (1), while the imaginary part leads to the wave kinetic equation. The total spectral wave intensity, / ( K ) , is given by the sum of individual wave intensity for each normal mode, designated by a,
/(*) = £ £W*(<"-<""£)• a
(2)
Specifically, we shall adopt a = Ln to designate the rath-harmonic Langmuir wave. For n = 1 (the fundamental Langmuir mode), the customary Bohm-Gross dispersion relation, uifc1 = ai pe (l + 3k2Xjje/2), is well-known, where X2Dt, = Te/(Aithe2) is the square of the debye length, Te being the electron temperature. By inserting Eq. (2) into the real part of Eq. (1), we obtain a nonlinear dispersion equation given by 1 00-Re(^ - Re (e(* a^) 4 V f dk' Ix^kV'-ff"" ^ ~ k',™£" ~ ^ ( " ' 1 ) ) | 2 V \ r '- D O (W) O Jr Wfk) O. (k,™k )-A^Jdk e ( k _ k > w L»_ f f ^(-i))
(3)^
where n > 2. After some suitable approximations are made, Eq. (3) can be shown to reduce to
4 " k ' = { [ ( " - 1)k2 k ' + nfc'2 k] • (k - k') + n (n - 1) |k - k'| 2 (k • k ' ) } 2 {n4 (n - I) 4 k2 ka |k - k f } " 1 ,
(4)
where u>^ik, is the fundamental Langmuir mode dispersion relation. The detailed deductive analysis of the above equation leads to the following desired dispersion relation for the nth-harmonic Langmuir mode:
^
= ^ (n + 4 " ' + \ k2X2De + ^
Ak),
(5)
where
^-2i^^/ d 3 k '^'^-')( k ')( f c / 2 - k - k ' + !p)'
(6)
with #k = 0. For the purpose of illustration, let us consider a specific model for the harmonic Langmuir-wave spectra. As noted already, on the basis of simulations and experiments, the nth-harmonic Langmuir mode (Ln) can be modeled with a spectrum with average wave number located at roughly nko ~ nu>pe/Vo. On the basis of this consideration, we model the one-dimensional harmonic Langmuir mode spectra by ILn(k) =
In(*1/2S)-1e-(k-nk°f'S2,
where In = J dk /z,n(/c), and S represents the spread associated with the spectra. This leads to
u!pe
2
2
\k0
2) ui£e
where k0 ss uJpe/V0. To test the idea of the present nonlinear eigen-mode theory of harmonic generation, we have also performed onedimensional electrostatic Vlasov simulation. The full details of the simulation technique and the in-depth analysis of
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the results are the focus of the paper by Umeda et al. (2003). The simulation result shown in Fig. 1 corresponds to the intensity of the waves plotted in grayscale format against normalized frequency and wavenumber, u)/u:pe and kV0/ujpe. The input parameters for the simulation are nt/n0 = 10~3, Vo = 3.5 ve, and Te/Tb = 4. The Fourier transformation is performed over the simulated data in both space and time, and the result is a simulated ui-k dispersion diagram as shown in Fig. 1. We have supperposed the theoretical dispersion relation curves, given by Eq. (7) on top of the numerically generated wave intensity versus u> and k. The result is the comparison between the theory and simulation. The result is an excellent agreement between the simulation result and theory.
- Fig. 1. Simulated dispersion diagram for harmonic Langmuir waves. The intensity of the waves are Fourier analyzed in both space and time, and the result is plotted in grayscale format against normalized frequency and wavenumber, u/wpe and kV0/ujpe. We have superposed the theoretical dispersion relation curves (7).
WAVE KINETIC EQUATION AND SATURATED WAVE SPECTRUM Now let us consider the imaginary part of the nonlinear spectral balance equation (1). For the fundamental Langmuir mode (LI), the complete nonlinear wave kinetic equation, which includes nonlinear decay and induced scattering processes, have already been derived by Yoon (2000) and numerically solved by Ziebell et al. (2001). For the present purpose, however, the nonlinear wave couplings do not play a significant role, and since the time domain of interest is sufficiently short, the effective wave kinetic equations for all harmonic modes are the quasilinear wave kinetic equation,
«&« = » W 4 / dvS^ -k • V) k • ^ W).
W
This equation needs to be closed by the particle kinetic equation for the electrons,
Equation (8) has the same structure as the conventional quasilinear wave kinetic equation, except for the overall coefficient proportional to n2. Therefore, even without solving this equation, one can readily see that the harmonic modes will start to grow in the linear regime, provided that a minimum level of spectral intensity exists for these modes. The initial electron distribution function is given by a Maxwellian (thermal) core plus an energetic beam component, while the ions are treated as quasi-stationary. We have numerically solved the complete set of wave and particle kinetic equations, Eqs. (8) and (9) in one-dimensional limit. In the present scheme, we employ the following ad hoc procedure to define the initial state for all eigenmodes: For a given n-th harmonic, the small level of initial spectrum is modeled by a gaussian form,
W*) = ^ e x p ( - ^ M 2 ! ) ,
(10)
where kn is the normalized wave number associated with the primary LI mode. This choice is guided by the linear growth property which dictates that the n-th harmonic mode should grow in the vicinity of k ~ nkL1, with some spread in wave number, characterized by A. The precise functional form is not crucial in this regard. The quantity /„ is determined by the expression In = he-
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(11)
where / i , a, and /? are all constants that can be arbitrarily chosen. It turns out that this particular profile allows for a combination of exponential and power-law dependence among the peaks of saturated harmonic mode spectra. Since the dynamical evolution of the harmonic modes is dictated by quasilinear equation, the choice of input will be directly reflected in the saturation spectra. Our choice (11) gives us sufficient freedom to adjust our theory to match the simulation result to be shown later.
- Fig. 2. Time evolution of the normalized wave intensity. max/2n(£)/(87rn'ie), versus uipf,t, in logarithmic vertical scale, for the first ten harmonics.
Figure 2 shows the time evolution for the peak of each harmonic mode spectrum, max I^n(t)/(87rnTe). The parameters relevant to determine the initial spectra are I\ = 10~3, a = 5 and /? = 2.236. Note that the higher the harmonic mode number, the faster the mode grows initially, in agreement with the simulations (Klimas, 1983, 1990). Note also that all eigenmodes reach saturation at about the same time. The evolution of the wave-number spectrum for weak harmonic Langmuir turbulence from the linear phase until quasi-saturation stage can be seen in Fig. 3. In this figure, where we have plotted the superposition of all the harmonic wave intensities,
l + (k)=
£
7 + ,(fc),
n=l,2,3,...
the total time interval ranges from u>pet = 0 to uipet = 3400. Note that the initial form of superposed spectra (dashed line) is not in the power-law form, but it achieves a power-law spectral shape at the saturation stage. This is owing to the fact that the higher harmonics grow faster than the lower harmonics. Our choice of spectral shape parameters, a and j3, and the specific form of initial spectra were partly designed to produce a power-law form at quasi-saturation stage. If we connect the peaks of the individual harmonics, then one obtains an overall power-law, /+(fc)ocAT5. The spectrum shown in Fig. 3 bears a qualitative resemblance with some measurements made on weak beam-plasma systems [see, for instance, Fig. 4 of the paper by Apel (1969)]. The spectral property in terms of frequency, instead of wave number, follows the same power-law pattern and is not shown here. - Fig. 3. Evolution of the total wavenumber spectrum, 7+(fe)/(87rnTe) = £™ =? I+n(k)/(SnhTe), versus kve/uipe (in log-log scale), showing up to 10 harmonics. Since the results presented here practically correspond to the quasilinear stage of the kinetic evolution, the electron distribution function is mostly affected by the linear wave -particle interaction with the combined fundamental and nonlinear harmonic Langmuir modes. However, the energy content of the harmonic modes is very low compared to the fundamental mode. As a consequence, the temporal evolution of the distribution function is very similar to the customary quasilinear theory which does not include the harmonics.
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- Fig. 4. Simulated spectrum at saturation. The straight diagonal line represents the power-law intensity distribution, \Ex\2(kx) <x k~5. To compare the theoretical spectrum with the simulation, we present the simulated spectrum in Fig. 4, where the normalized wave intensity near the saturation is plotted versus the wave number, in log-log scale. The powerlaw index of ~ —5 to —6 associated with the initial noise spectrum is indicated by the diagonal line. The wave spectrum shown in the figure corresponds to the near-saturation time of upet = 819.2. As the readers can appreciate, the simulation shows that the final wave intensity features a quasi power-law-like spectra with an index of roughly between —5 and —6. From this, it appears that the turbulence spectrum of ~ —5 to —6 seems to be some sort of quasi universal constant, characterizing the Langmuir turbulence. GENERATION O F S U P E R T H E R M A L ELECTRONS The foregoing discussion deals with physical processes which take place relatively early in the evolution of beamplasma system. That is, according to our theory, the excitation of harmonics of Langmuir mode and the formation of quasi power-law spectral distribution of the waves occur during the quasilinear evolution stage of the electrons. During this phase, the system can be described essentially by the quasilinear particle diffusion equation (9) as well as the quasilinear wave kinetic equation (8). The only place where the nonlinear physics becomes important is in the computation of the nonlinear dispersion relation (5) for the harmonic modes, Ln. The formation of superthermal electron tail is, on the other hand, a process which occurs later in time, and it is intimatedly associated with the fully nonlinear wave-particle interaction. The superthermal electron tail has been known in relation to the nonlinear beam-plasma interaction process, but to this date, no definitive physical mechanism has been put forward. In the generalized weak turbulence theory by Yoon (2000), the particle kinetic equation has extra terms which represent hyper-diffusion process resulting from the fully nonlinear wave-particle interactions. For the electrons, the generalized particle kinetic equation is given by 8Fe
d f
dFe\
d* f
dF
d*Fe \
where
* W ^ / * f E «(™fc - k • v) i£L, Mtjk^) = ~-^jdkjdkJ
h k', (k - k')* - ^ p Y. ^
«««(*) = 2 ^ / * / * ' H ^
- "'"*') S ^ - *'"*' - (k - k') • vl &' ^
5 (a^-a'4,)2%4-a'^,-(k-k')-v]/k
(13)
In the above, L mode is the primary Langmuir mode, LI. The first term on the right-hand side of Eq. (12) corresponds to the quasilinear diffusion (plus the collisinal drag) term. The extra terms on the right-hand side of Eq. (12) associated with higher-order derivatives become important only after the quasilinear plateau-formation stage. During the stage where power-law spectrum is generated, discussed in the previous section, it is safe to disregard the higher-order derivative terms. By the time the full particle kinetic equation (12) comes into play, however, the simple quasilinear wave kinetic equation is no longer valid, but a fully nonlinear particle kinetic equation (12) as well as the fully nonlinear wave kinetic equation (Yoon, 2000; Ziebell et al., 2001) must be adopted instead. It turns out that the harmonic modes do not participate in the generation process of the superthermal electron tail. Therefore, in the -257-
discussion of the superthermal tails, the harmonic modes can be ignored. Thus, the equations for waves (L = LI and S, which stands for ion-sound mode) are
+2
2^
" i
w
k
I *
^k,k' ''(""k - " wk, - (r u k _ k ,) I
-
^k I
£ /dk'/ dv ^l? <5l<7 ^- a '^'-( k - k ' ) - v l ^ - ^ ( ^ ' ^ - ^ ^ ( ^ e + FO + l£,h /£L (k - k') • ± (jo^ - a'ub)F9 - ^ (a^) Fi)j,
(14)
and
+
2^ tr',
dk
k
v
kM'°yau;k~
a u
k' ~a
u^v)
lauiklw
i k _ k / - o- ujy., i k _ k , J k
-CT o;k_k, _Zk, J k
I, (15)
^
where the decay coefficients, V ^ , and V ^ , , are defined by _ ^ e 2 M k _ k >(k-k') 2 ' ' 4T e 2 fc 2 fc' 2 |k-k'| 2 '
L k k
s k k
' '
Tre^^lk'-Ck-k')]2 4Te2 A; 2 fc' 2 |k-k'| 2 '
l
'
and u»k = a;pefcADe('"e/mj)1/2(l + 3Ti/Te)1^2(l + fe2AQe)^1^2 is the familiar ion-sound mode dispersion relation, and the quantity yUk is defined by Mk = ^ L ( ™ e / m i ) 1 / 2 (1 + ZTjTey2. The numerical solution of the nonlinear particle kinetic equation (12) together with the fully nonlinear wave kinetic equations for LI and 5 modes (14) and (15) in one-dimensional result is shown in Fig. 5. The numerical result shows that the initial distribution is substantially broadened in velocity space as a result of the nonlinear hyper-diffusion process. To convince ourselves that the generation of superthermal tail is indeed owing to the hyper-diffusion term, we arbitrarily turned the hyper-diffusion term off in our code. We found that the final distribution is virtually identical to the initial distribution except for the plateau region. Therefore, it became clear that the generation of superthermal tail cannot be discussed on the basis of diffusion approximation.
- Fig. 5. Generation of superthermal tail in the electron distribution by nonlinear hyper diffusion terms. The initial distribution and the final distribution are shown. The superthermal tail formation cannot be discussed on the basis of quasilinear diffusion equation, but is a result of the hyper-diffusion terms which appears in the generalized turbulence theory by Yoon (2000).
CONCLUSIONS A N D DISCUSSION In this paper, we have briefly reviewed the history of Langmuir wave turbulence as developed over the past four decades or so. We have then presented a new twist on the Langmuir turbulence scenario which involves the generation of harmonics of Langmuir waves, which is essentially a quasilinear process, although the existence of the harmonic modes themselves cannot be discussed on the basis of linear theory. These modes are shown to form a quasi power-law
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distribution during the beam plateau formation stage, and saturate early. Then, the fully nonlinear wave-coupling processes take over, which lead to the formation of superthermal electron tail. Superthermal electron tail cannot be discussed on the basis of traditional quasilinear diffusion equation plus fully nonlinear wave kinetic equations, but in order to discuss this effect, one must include the new hyper-diffusion term in the generalized particle kinetic equation, as first derived in the paper by Yoon (2000). The combination of the harmonic mode excitation and hyper-diffusion thus seems to be capable of accounting for two basic features which characterize the Langmuir turbulence, namely, the quasi power-law wave spectrum and the generation of superthermal electron population in the tail of the distribution. In light of these findings, it is therefore, proposed that a number of outstanding issues associated with the Langmuir turbulence can be explained by the author's generalized weak turbulence theory. ACKNOWLEDGMENTS The author gratefully acknowledges the hospitality of the Radio Atmospheric Sience Center (RASC), Kyoto University, Japan, during his visit in the Summer of 2002, as well as technical support and fruitful discussions with Y. Omura, R. Gaelzer, and T. Umeda. He is particularly appreciative of T. Umeda for providing his numerical simulation results, and for R. Gaelzer for carrying out some of the numerical integrations of theoretical equations. This research was supported by the Department of Energy (DOE) grant DE-FG02-00ER54584. REFERENCES Adam, J. C , G. Laval, and D. Pesme, Phys. Rev. Lett, 43, 1671, 1979. Akhiezer, A. I., I. A. Akhiezer, R. V. Polovin, A. G. Sitenko, and K. N. Stepanov, Plasma Electrodynamics, Pergamon, New York (1975). Akimoto, K., H. L. Rowland, and K. Papadopoulos, Phys. Fluids, 31, 2185, 1988. Apel, J. R., Phys. Rev. Lett, 19, 744, 1967. Apel, J. R., Phys. Fluids, 12, 640, 1969. Boswell, R. W., and P. J. Kellogg, Geophys. Res. Lett., 10, 565, 1983. Davidson, R. C., Methods in Nonlinear Plasma Theory, Academic, New York, 1972. Dawson, J. M., and R. Shanny, Phys. Fluids, 11, 1506 (1968). Dreicer, H., I. C. Ingraham, and D. B. Henderson, Phys. Rev. Lett, 26, 1616, 1971. Dum, C. T., J. Geophys. Res., 95, 8095, 1990a. Dum, C. T., J. Geophys. Res., 95, 8111, 1990b. Dum, C. T., J. Geophys. Res., 95, 8123, 1990c. Dupree, T. H., Phys. Fluids, 9, 1773, 1966. Dupree, T. H. Phys. Fluids, 15, 334, 1972. Gaelzer, R., L. F. Ziebell, and P. H. Yoon, Phys. Plasmas 9, 96, 2002. Gaelzer, R., P. H. Yoon, T. Umeda, Y. Omura, and H. Matsumoto, Phys. Plasmas, 10, 373, 2003. Galeev, A. A., R. Z. Sagdeev, Yu. S. Sigov, V. D. Shapiro, and V. I. Shevchneko, Sov. J. Plasma Phys., 1, 5, 1975. Gentle, K. W., and J. Lohr, Phys. Fluids, 16, 1464, 1973. Goldman, M. V., Rev. Mod. Phys., 56, 709, 1984. Goldstein, B., W. Carr, B. Rosen, and M. Seidl, Phys. Fluids, 21, 1569, 1978. Hasegawa, A., Plasma Instabilities and Nonlinear Effects, Springer, New York, 1975. Iroshnikov, P. S., Sov. Astron., 7, 566, 1964. Joyce, G., G. Knorr, and T. Burns, Phys. Fluids, 14, 797, 1971. Kadomtsev, B. B., Plasma Turbulence, Academic, New York, 1965. Kaplan, S. A., and V. N. Tsytovich, Plasma Astrophysics, Pergamon, Oxford, 1973. Kasaba, Y., H. Matsumoto, and Y. Omura, J. Geophys. Res., 106, 18693, 2001. Kellogg, P. J., S. J. Monson, W. Bernstein, and B. A. Whalen, J. Geophys. Res., 91, 12065, 1986. Klimas, A. J., J. Geophys. Res., 88, 9081, 1983. Klimas, A. J., J. Geophys. Res., 95, 14905, 1990. Kolmogorov, A. N., Dokl. Akad. Nauk SSSR, 30, 299, 1941. Kraichnan, R. H., Phys. Fluids, 8, 1385, 1965. Kruer, W. L., P. K. Kaw, J. M. Dawson, and C. Oberman, Phys. Rev. Lett, 24, 987, 1970. Laval, G., and D. Pesme, Phys. Fluids, 26, 52, 1983. Levitskii, S. M., and I. P. Shashurin, Sov. Phys. JETP 25, 227, 1967. Liang, Y.-M., and P. H. Diamond, Phys. Fluids B, 5, 4333 (1993). Llobet, X., W. Bernstein, and A. Konradi, J. Geophys. Res., 90, 5187, 1985. Malmberg, J. H , and C. B. Wharton, Phys. Fluids, 12, 2600, 1969.
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Manheimer, W. M, Phys. Fluids, 14, 579, 1971. McComb, W. D., The Physics of Fluid Turbulence, Oxford University Press, Oxford, 1990. Melrose, D. B., Plasma Astrophysics, Gordon and Breach, New York, 1980. Mizuno, K., and S. Tanaka, Phys. Rev. Lett, 29, 45, 1972. Mori, H., J. Phys. Soc. Japan, 35, 592, 1973. Morse, R. L., and C. W. Nielson, Phys. Fluids, 12, 2418 (1969). Nishikawa, K.-I., and I. H. Cairns, J. Geophys. Res., 96, 19343, 1991. O'Neil, T. M., J. H. Winfrey, and J. H. Malmberg, Phys. Fluids, 14, 1204, 1971. O'Neil, T. M., and J. H. Winfrey, Phys. Fluids, 15, 1514, 1972. Omura, Y., H. Matsumoto, T. Miyake, and H. Kojima, J. Geophys. Res., 101, 2685, 1996. Pritchett, P. L., and J. M. Dawson, Phys. Fluids, 26, 1114, 1983. Robinson, P. A., Rev. Mod. Phys., 68, 507, 1997. Robinson, P. A., and D. L. Newman, Phys. Fluids B, 2, 2999, 1990. Sagdeev, R. Z., and A. A. Galeev, Nonlinear Plasma Theory, Benjamin, New York, 1969. Schriver, D., M. Ashour-Abdalla, V. Sotnikov, P. Hellinger, V. Fiala, and A. Mangeney, J. Geophys. Res., 105, 12919, 2000. Seidl, M., W. Carr, D. Boyd, and R. Jones, Phys. Fluids, 19, 78, 1976. Sitenko, A. G., Electromagentic Fluctuations in Plasmas, Academic, New York, 1967. Sitenko, A. G., Fluctuations and Nonlinear Wave Interactions in Plasmas, Pergamon, New York, 1982. Theilhaber, K., G. Laval, and D. Pesme, Phys. Fluids, 30, 3129, 1987. Tsunoda, S. I., F. Doveil, and J. H. Malmberg, Phys. Rev. Lett, 58, 1112, 1987. Tsytovich, V. N., Nonlinear Effects in a Plasma, Plenum, New York, 1970. Tsytovich, V. N., An Introduction to the Theory of Plasma Turbulence, Pergamon, New York, 1977a. Tsytovich, V. N., Theory of Plasma Turbulence, Consultants Bureau, New York, 1977b. Umeda, T., Y. Omura, P. H. Yoon, R. Gaelzer, and H. Matsumoto, Phys. Plasmas, 10, 382, 2003. Utlaut, W. F., and R. Cohen, Science, 174, 245, 1971. Vedenov, A. A., Theory of Turbulent Plasma, Elsevier, New York, 1968. Vifias, A. F., H. K. Wong, and A. J. Klimas, Astrophys J., 528, 509, 2000. Vyacheslavov, L. N., V. S. Burmasov, I. V. Kandaurov, E. P. Kruglyakov, O. I. Meshkov, and A. L. Sanin, JETP Lett. 75, 41, 2002. Weinstock, J., Phys. Fluids, 12, 1045, 1969. Whelan, D. A., and R. L. Stenzel, Phys. Fluids 28, 958, 1985. Yin, L., M. Ashour-Abdalla, M. El-Alaoui, J. M. Bosqued, and J. L. Bougeret, J. Geophys. Res., 103, 29619, 1998. Yoon, P. H., Phys. Plasmas, 7, 4858, 2000. Yoon, P. H., R. Gaelzer, T. Umeda, Y. Omura, and H. Matsumoto, Phys. Plasmas, 10, 364, 2003. Zakharov, V. E , Sov. Phys. JETP, 35, 908, 1972. Ziebell, L. F., R. Gaelzer, and P. H. Yoon, Phys. Plasmas, 8, 3982, 2001.
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NEW NON-STOCHASTIC ACCELERATION IN MULTI-COMPONENT PLASMAS T. Mizuta1 and M. Hoshino1 'The University of Tokyo, Hongo 7-3-1 Bunkyo-ku Tokyo, JAPAN
ABSTRACT We discuss a non-stochastic particle acceleration mechanism under the cyclotron resonance with two electromagnetic ion-cyclotron (EMIC) waves. We find that the strong ion acceleration occurs in multi-component plasma if a heavy ion can resonate with both the EMIC wave near the cutoff frequency and the other EMIC wave with the Alfven velocity. The particle is preferentially accelerated perpendicular to the ambient magnetic field, and the acceleration rate is simply expressed by VL/VA = (EL I Bo )(c / VA )O.t, where Bo is the ambient magnetic field, V± is a velocity perpendicular to the ambient field, V\ is Alfven velocity, c is the light speed, Q is a cyclotron frequency, t is time and E± is electric field of the EMIC wave near the cutoff frequency.
INTRODUCTION Particle acceleration in collisionless plasmas is one of the most important problems in space plasma physics. The most popular mechanism is a "stochastic" acceleration in which particles interact with turbulent waves in random phases, for example, Fermi acceleration. Since the acceleration time scale of the stochastic process is in general much slower than the cyclotron period, it is believed that the other fast acceleration process is probably occurring in many active acceleration regions such as the solar corona and the terrestrial aurora etc. We discuss a resonant acceleration of particles with two EMIC waves, which provides the "non-stochastic" acceleration process. Although the particle cannot be effectively accelerated under the standard two EMIC waves in ion-electron plasma, we find that the strong acceleration occurs in a multi-component plasma if the particle can resonate with both the EMIC wave near the cutoff frequency and the other EMIC wave with the Alfven velocity (Mizuta and Hoshino, 2001). In this paper, we discuss the phase bunching of the resonant particles in this non-stochastic acceleration process, and we show the phase bunching plays an important role on this acceleration. IN ONE EMIC WAVE We make a brief review of the particle behavior in one coherent EMIC wave propagating parallel to the ambient magnetic field (e.g. Matsumoto, 1979; Terasawa, 1989). In this case there is a constant ^ , which is
2
Bok
^
(1)
where FJ is velocity component parallel to the magnetic field, BL is wave magnetic field, k is a wave number, y/ is the phase angle between VL and B±, and VR is the resonant velocity, VR = (co-Q.)/k (CO is a wave frequency). Figure la shows a typical particle motion in (i//,VA space. We write y/ = 0 as " f i ± " because V, has the same direction as BL • In the similar way we call y/-3127c as "EL", etc. The particles having V =VR are initially moved toward y/ = EL where they are "bunched" to the EL direction, and they can be accelerated by the wave
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electric field (solid arrow). After that, particles are transported to - El along the constant % curve (dashed arrow). We call the circle in figure la "trapping circle". The radius of the trapping circle Vt is
Note that as the wavelength is longer, the trapping circle becomes larger as well. We study a test particle motion of ion by integrating the Lorentz equation in time for a given wave. The simulation parameters is listed in the table 1 as the case of RUN 1. Figure lb shows the time evolution of VL for RUN1, the horizontal axis is time normalized by the inverse of the proton cyclotron frequency (fi" 1 ), the vertical axis is F± normalized by the Alfven velocity ( VA) and the gray scale indicates the number density in a log scale. One can find that the particles are accelerated from y[ == o to 1.5 within t = 20Q.p', and after that, they are deceleration phases by turn. This mechanism is known as the pitch-angle scattering process of particle trapped by a finite amplitude wave.
Fig. 1. (a): A typical particle motion in (y/,V,,) space. Thin solid curves are constant _£•. At first the particle is bunched to EL (solid arrow) and after that, the particle is trapped and transported to EL (dashed arrow), (b): the time evolution of VL for rani in table 1.
Table 1. Parameters of the test particle simulations Runs RUN1: n-EMIC RUN2: n-EMIC and 1-EMIC Common parameters Particles Initial bulk velocity
: He2 ions ( £ ^ / ; ?+ = 0.5fi p ) : (V^, I\)
=
Initial thermal velocity: 0.021^ for any direction n-EMIC: k = -0.23Q ; ) F;',co = 0.23Qp ,B±=0.\0BQ,VR
1-EMIC: k = 0.030QpVA[,a ; = 0.53Q;j,
=VA
Fig. 2. a dispersion relation in a multicomponent plasma which consists protons, He21 and electrons.
IN TWO EMIC WAVES Next we discuss the particle behavior in the two EMIC waves. Figure 2 shows the dispersion relation in a multi-component plasma which consists of protons, He21 ions and electrons (Smith and Brice, 1964). We define the sign of CO and the phase velocity (col k) as follows: when a wave is left-hand (right-hand) circularly polarized, we take co>0 (a> < 0), while when the wave having phase velocity parallel (antiparallel) to the ambient field, we
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take wl k > 0 (col k < 0). The He ion can resonate with both n-EMIC (normal wavelength EMIC wave) and 1EMIC (long wavelength EMIC wave). We study the test particle motion in these waves. The simulation parameter is given in table 1 as the case of RUN2. We assumed that two waves have the almost same Poynting flux. Figure 3a shows the time evolution of VL in the same format as figure lb. It is quite different from the case of RUN1, namely, the most of particles are accelerated without deceleration. We confirmed that this continuous acceleration occurs up to / = 10000Q~' (figure3b). Figure 4 is the same format as figure la, the horizontal axis is the phase of nEMIC (if/n) and 1-EMIC (i//l), the vertical axis is V^ and the gray scale is integrated number density during t =0-100fi~'. The black dashed circles are the trapping circles for each waves. From eq.(2) we know that the trapping circle of 1-EMIC is larger than that of n-EMIC. The He2 ions are trapped well inside the n-EMIC trapping region, while the most of ions are bunched to the strong resonance region of EL of 1-EMIC. At first, the particles are bunched to £ of both waves like as the solid arrows in figure lb. The idea of the constant % is not applicable to this case with two EMIC waves, but we find that the particles move along the constant % curves of both waves. During the trapping motion of He2' ions with n-EMIC wave, the parallel velocity pj must have same value, and the particles must oscillate near£ ; . Therefore the particles can resonate with a specific phase for 1-EMIC, EL l and gain energy through EL, directory, though the particles travel in various phases for n-EMIC and the forces to the particles are canceled. Based on the above discussion, we assume that the particles feel only EL, and from Fig. 3. The time evolution of Vx for run2 in table 1 (a) to momentum equation we obtain, t = 200Q; 1 , (b) to / = i oooon; 1 . ^ V 'A
= ^ • - ^ 1 t. R "0
V
(3)
'
"A
. This relationship is shown in the white dashed line in figure 3a. We find a small number of particles (about 7%) do not "unluckily" participate in the non-stochastic acceleration process and they are pitch-angle scattered. Since the frequency of the pitch-angle scattering shows that the particles are scattered by 1-EMIC wave, they are trapped by 1-EMIC (figure 4b). After one period of the pitchangle scattering (/ = 120n~'), the most of unlucky particles start to participate in the non-stochastic acceleration process. In this way, the number of particles that enter into this acceleration process increases with time. DISCUSSION We discussed that He21 ions can preferentially resonate with a long wave having a long wavelength around the cutoff frequency (figure 2). This acceleration process is important for many space plasma applications in the solar corona, ionosphere, etc. Cranmer et. al. (1999) shows that, in the solar polar coronal holes, O" ions are much hotter than protons and the perpendicular temperature is hotter than the parallel one. The ratio of the temperature between O3 ions and protons is much larger than the ratio of the mass between those, which implies that a preferential acceleration process works to O3' ions. We have not discussed how a coherent long EMIC wave is excited. One of the possible processes is an anisotropic thermal He2 ion beam instability (i.e. Killen et. al., 1995; Gary et. al. 2000), other possible process is a parametric instability (i.e. Hoshino and Goldstein, 1989; Kubo, 1986). The excitation mechanisms of this mode should be studied further. Roughly speaking, one of the necessary conditions for this non-stochastic acceleration process is that two waves have the almost same Poynting flux. When one of two waves has much smaller Poynting flux than the other wave, the particles move in such a way that they neglect the constant % of the wave having much smaller Poynting flux, and the particles are trapped by the wave having much larger Poynting flux. Therefore, the particles are pitch-
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angle scattered by the wave of the larger Poynting flux and our non-stochastic acceleration process does not occurs. In more details, we found that the conditions imposed on the amplitudes of the two EMIC waves for this acceleration process are Blf < BL and AtBLI > ZnBLn , where X- are the wavelength of the 1-EMIC and the n-EMIC, respectively (Mizuta and Hoshino, 2003). The non-linear effect of the accelerated particle to the EMIC wave is an important issue. Tanaka (1985) has performed self-consistent numerical simulations for the He2 acceleration problem, and he found some of He2+ ions can be accelerated perpendicular to the magnetic field. He found that the velocity of He21 has the same direction with EL of EMIC waves. We consider that these self-consistent simulation results have close relation to our non-stochastic acceleration model, and it is further studying. number density number density Fig 4. Integrated number density during / = 0-100fi~' Horizontal axes are the phase for (a) n-EMIC ( y/n) and (b) 1-EMIC ( \f/t ), vertical axes are V,, • Black dashed circles are typical trapping circles.
REFERENCES Cranmer, S. R., et. al., An empirical model of a polar coronal hole at solar minimum, Astrophys. J., 511, 481 (1999). Fermi, E., On the origin of a cosmic radiation, Phys. Rev., 75, 1169 (1949). Gary, S. P., et. al., Alpha/proton magnetosonic instability in the solar wind, J. Geophys. Res., 105, 20989 (2000). Hoshino, M. and M. L. Goldstein, Time evolution from linear to nonlinear stages in magnetohydrodynamic parametric instabilities, Phys. Fluids, Bl(7), 1405 (1989). Kennel, C. F. and F. Engelmann, Velocity space diffusion from weak plasma turbulence in a magnetic field, Phys. Fluids, 9, 2377 (1966). Killen K., et. al., Linear and nonlinear properties of ULF waves driven by ring-beam distribution functions, J. Geophys. Res., 100, 5835 (1995). Kubo, The study of the parametric instability of Alfven wave in the 3-component system, master thesis, University of Tokyo (1986). Matsumoto, H., Non linear whistler-mode interaction and triggered emissions in the magnetosphere : a review, Wave Instabilities in Space Plasmas, D. Reidel Pub. Co., 74, 163 (1979). Mizuta, T. and M. Hoshino, Preferential acceleration of heavy ions multi-component plasmas, Geophys. Res. Let., 2S,3099(2001). Mizuta, T. and M. Hoshino, A new acceleration process in two electromagnetic ion-cyclotron waves having different wavelength, to be submitted Phys. of Plasma (2003). Smith, L. and N. Brice, Propagation in multicomponent plasma, J. Geophys. Res., 69, 5029 (1964). Tanaka, M., Simulations of heavy-ion heating by electromagnetic ion-cyclotron waves driven by proton temperature anisotropies, J. Geophys. Res., 90, 6459 (1985). Terasawa, T., Particle scattering and acceleration in a turbulent plasma around comets, Geophysical Monograph, American Geophysical Union, 53, 41 (1989). E-mail address of T. Mizuta mizuta(5),space,eps.s.u-tokvo.ac,ip Web address of T. Mizuta http://stp-www.eps.s.u-tokvo.ac.jp/~mizuta/
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SECTION 5: Shocks and Interplanetary Phenomena
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GEOTAIL OBSERVATIONS OF SOLAR WIND AND INTERPLANETARY PHENOMENA T. Terasawa 1 1
Department of Earth and Planetary Science, University of Tokyo, 7-3-1 Kongo, Bunkyo-ku, Tokyo 113-0033, JAPAN ABSTRACT
In this review, based on the recent GEOTAIL observation I will cover three broad topics, bow shock and foreshock phenomena, interplanetary phenomena, interstellar medium, and solar flare effects. (1) I will concentrate on the physics of energetic diffuse ions found in the bow shock and foreshock regions, and discuss three subtopics, namely, the origin of these ions, their nonlinear reaction to the shock structure, and their transport process from the region upstream of the nose bow shock to the predawn foreshock region. (2) I will discuss two subtopics about propagating interplanetary shocks ahead of coronal mass ejecta: It is shown that in addition to the well-known diffusive ion acceleration at these shocks there sometimes occurs diffusive electron acceleration. The nonlinear reaction of the accelerated particles, the same phenomenon as found in the bow shock case, is identified at least once in an moderately strong interplanetary shock. (3) In addition to the above 'normal' bow shock/interplanetary phenomena, GEOTAIL has provided unique 3D phase space information on a 'stranger', namely pickup He+ ions of local interstellar medium origin. (4) I will then describe unexpected direct observations of solar flare signals on GEOTAIL: the spiky enhancements (duration ~ several minutes) of the sunward electric fields and background counts of plasma detectors, both of which occur concurrently with the peaks of large solar X-ray flares.
INTRODUCTION The number of GEOTAIL papers published in years 1991-2003 treating topics of solar wind and interplanetary phenomena is found to be ~ 50*, which is slightly less than 10 % of all GEOTAIL papers published in the same interval (~600). This relatively small percentage reflects the fact that GEOTAIL is basically designed for the studies of the magnetosphere — especially the magnetotail -—, to which the main interest of the GEOTAIL scientific community is directed. Nevertheless, as briefly reviewed in this article, GEOTAIL has been providing unique and important datasets for the study of the regions surrounding the earth's magnetosphere. To get a measure how the interplanetary phenomena are covered by the GEOTAIL observations, let us start counting the number of observed interplanetary shocks (IPSs). During years 1999-2001 GEOTAIL recorded ~ 40 IPSs in the solar wind region upstream of the earth's bow shock. This number of IPSs amounts to 36% of all IPSs detected by SOHO during the same interval. This percentage, being consistent with the fact that GEOTAIL has been exposed to the solar wind for ~ 40% of time during its orbital motion (Figure la). Figure la, at the same time, shows how the bow shock and its upstream region are covered by the GEOTAIL observation. Earlier in its mission phase, GEOTAIL covered more distant-tail part of the magnetosphere (Figure lb). During this orbital phase the orbital percentage for the solar wind region was much limited, but data for several important events were obtained (see the following discussion). In this review, I will choose a limited number of topics from the solar wind/interplanetary/upstream/flare studies based on GEOTAIL observations consisting of the low energy plasma experiment (LEP; Mukai et * We have not included papers simply using GEOTAIL as a solar wind monitor.
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al., 1994), the comprehensive plasma instrument (CPI; Frank et ah, 1994), the energetic particles and ion composition instrument (EPIC; Williams et al., 1994), the high energy particle experiment (HEP; Doke et al., 1994), the plasma wave instrument (PWI; Matsumoto et al., 1994), the electric field experiment (EFD: Tsuruda et al., 1994), and the magnetic field experiment (MGF; Kokubun et al., 1994).
Fig. 1. (a) Sample orbits of GEOTAIL for the near-earth phase (1 March - 10 November 1995). (b) The orbit of GEOTAIL during the distant-tail phase (1 January - 31 December 1994). The nominal positions of the bow shock (a and b) and the magnetopause (a) are also shown. The unit of spatial scale is the earth radius ( R E ) .
BOW SHOCK AND FORESHOCK PHENOMENA In this section I concentrate on the topics of upstream energetic ions (of several tens of keV). I refer the reader to the original papers for other topics such as motions/positions of the bow shock (Tsubouchi et al., 2000; Dmitriev et al., 2002), whistler wave observations (Hayashi et al., 1994; Matsui et al., 1997; Zhang et al., 1998, 1999), 2fpe radio emissions (Reiner et al., 1997; Kasaba et al., 2000a), the hot flow anomaly (Sibeck et al., 1999), waves/particles and disturbance in the magnetosheath (Petrinec et al., 1997a, 1997b; Sibeck et al., 1997, 2000; Nemecek et al., 1998; Zong et al., 1998, 1999; Seon et al., 1999; Matsuoka et al., 2000; Safrankova et al., 2000; Zastenker et al., 2002; Nagano et al., 2003), the observations of anomalously low density solar wind in the magnetosheath (Kasaba et al., 2000b; Terasawa et al., 2000), and ULF waves associated with lunar wake (Nakagawa et al., 2003). It is also noted that Matsumoto et al. (1997) presented an excellent review of the PWI measurement of waves (frequency ~ 10 Hz) in the upstream and bow shock regions. Origin of diffuse ions Energetic ions up to ~100 keV (and sometimes ~several hundreds keV) have been observed both upstream and downstream from the earth's bow shock (e.g., Lin et al., 1974; Fuselier, 1995; Anagnostopoulos, 1994). These ions have a wide angular distribution, and are called 'diffuse ions'. It is commonly observed that large amplitude low frequency electromagnetic waves (0.01-0.1 Hz) accompany these diffuse ions. There are at least two possible origins for these diffuse ions: One is the diffusive shock acceleration (DSA, hereafter) process at the bow shock, and the other is the leakage of magnetospheric ions. One evidence of the DSA process, for example, is a good correlation between the diffuse ion density and the solar wind density (Trattner et al., 1994). There is, on the other hand, evidence of magnetospheric leakage: During the upstream diffuse ion events magentospheric ions such as O + , N +1 , and O +2 have been identified (e.g. Mobius et al., 1986; Sarris and Krimigis, 1988; Desai et al., 1999; Christon et al., 2000; Keika et al., 2003). Therefore it is quite important to determine to what extent each of these mechanisms contributes to form the upstream population of energetic ions. In the nominal foreshock region, namely the region of |YGSE| ~ 20 — 30 RE (as covered by the previous spacecraft, such as IMP-6, ISEE-1/2 and AMPTE), it turns out to be a difficult task to discriminate the above two possibilities: When the observer is connected to the quasi-parallel part of the bow shock by -268-
Fig. 2. (a) The GEOTAIL position around 08:30 on 19 February 1994 relative to the bow shock and the magnetopause. The dashed and solid arrows show the IMF of nearly Parker's spiral direction and nearly radial direction, respectively, (b) The E-t (energy versus time) plot of count rates of energetic ions (5-40 keV/q) in the directions sunward ±90°. (c) The magnetic field intensity, (d) The IMF longitudinal angle (
the IMF (interplanetary magnetic field), where the efficiency of DSA process is expected to be maximized, he/she is also connected to the magnetosheath region adjacent to the magnetopause, where the leakage of the magnetospheric particles is taking place. If we can observe diffuse ions in the foreshock region with a large cross-field separation from the magnetopause, we could circumvent this difficulty. Figure 2 shows such an event (Sugiyama et al., 1995b): Before ~08:28 the IMF was nearly in the Parker's spiral direction except for two short intervals around ~08:16 and ~08:19. Under such a spiral geometry (a dashed arrow in Figure 2a), GEOTAIL was magnetically disconnected from the bow shock and the magnetopause. At ~08:28 the field lines passing the GEOTAIL position began to intersect the quasi-parallel bow shock (a solid arrow in Figure 2a), but was separated from the nearest magnetopause surface by ~20 RE- In the E-t (energy versus time) plot (Figure 2b), diffuse ions > 20 keV/q appeared within a few minutes of the beginning of the field intersection with the bow shock. If these ions should have come from the magnetosphere, they need to have crossed the field lines by > 20 RE within a few minutes. However this is impossible since the cross-field diffusion over 20 RE needs at least ~3 hours even in the Bohm limit. Therefore, these diffuse ions are most naturally explained in terms of the DSA process at the bow shock connected to the observer by the IMF. Recently, an important progress has been made through the charge-state measurement of energetic ions > 50 keV in the foreshock and magnetosheath regions by the EPIC aboard GEOTAIL (Keika et al., 2003). Firstly these authors classified energetic ions into two categories, namely, those with low charge state (LCS, e.g. O + and N + ) presumably of the magnetospheric origin, and those with high charge state (HCS, e.g. O"1"6 and O +7 ) presumably of the solar wind origin. Secondly they have shown that the LCS ions are found preferentially in the duskside magnetosheath while the HCS ions are found preferentially in the dawnside foreshock regions. This spatial feature of the HCS ions provides a strong evidence that they are accelerated at the bow shock and its vicinity. Nonlinear effect It has been known that well ahead of the main shock ramp the upstream solar wind flow is decelerated by a few km/s ~ few tens of km/s in the foreshock region where the diffuse suprathermal ions (10-102 -269-
keV) have non-negligible pressure (Formisano and Amata, 1976; Diodato and Moreno, 1977; Gosling et al., 1978; Bonifazi et al, 1980a, 1980b, 1983; Zhang, Schwingenschuh, and Russell, 1995). In the general context of DSA process, shocks affected by this nonlinearity have been studied in terms of 'cosmic ray mediated' shocks (CRMSs, hereafter). The CRMS effect is considered essential in source regions of cosmic rays, such as supernova shocks and shocks in gamma ray bursters, for which only remote astronomical observations are possible. Therefore, the studies of CRMS are done mainly from the theoretical side (e.g., Drury and Volk, 1981; Ko, 1995; Ellison et al., 2000), and the bow shock provides a unique opportunity for in situ observations of CRMS. However, detailed studies of the CRMS nature of the bow shock have been prevented by the transient variations of the solar wind: It has not been clear how the deceleration of the solar wind flow correlates with the diffuse ion energy density. Terasawa et al. (2001), on the other hand, has presented a case study of an exceptionally clear example of the bow shock observation as a CRMS (Figure 3).
Fig. 3. Data are shown for the period from from 22:30 UT on 8 October (t = 22.5 h) till 02:12 UT on 9 October (t = 26.2 h): (a) Energy-versus-time plots of energetic ions (5-40 keV/q) and (b) for solar wind ions (0.3-8 keV/q) where counts per 12-sec sample are shown in a pseudo-color scale. (White stripes at t ~ 24 in both panels were due to a data gap.) (c) Solar wind velocities observed by GEOTAIL (red curve) and WIND (black dotted line), (d) Energy density of diffuse ions (EDJ: red) and magnetic and thermal proton subpressures (PB' blue, and Psw: black). To avoid overlaps, Psw is shifted down by one decade, (e) The amount of the solar wind deceleration (AVSW) is plotted against the energy density of diffuse ions (Ex>j), where AVSW is defined as V™,GEOTAIL - V^.wiND- (f) The GEOTAIL orbit from 0 UT on 8 October 1995 to 0 UT on 10 October 1995. Nominal shapes of the bow shock and magnetopause are also drawn. The GEOTAIL crossing of the bow shock at 01:34 UT on 9 October 1995 was near the subsolar point at (X, Y, Z ) G S E = ( 1 3 . 8 , 1.0, 1.1) RE-
Figure 3 (f) shows the orbit of GEOTAIL for 48 hours: From 22:00 UT on 8 October 1995, the GEOTAIL spacecraft traversed the foreshock region of the nose bow shock over 3 hours. During this interval the solar -270-
wind was more or less steady and continuously monitored by the WIND spacecraft cruising about 130 RE upstream of the GEOTAIL position (The estimated convection delay time of the solar wind from WIND to GEOTAIL was ~22 minutes.) Figure 3 (a) and (b) respectively show the energy-versus-time (E-t) plots of diffuse ions (propagating sunward) and solar wind ions observed aboard GEOTAIL. Figure 3 (c) shows changes of the solar wind velocity VJ^GEOTAIL at GEOTAIL (red curve) and V5WIWIND at WIND (black dotted line). The GEOTAIL crossing of the bow shock at t = 25.57 h (01:34 UT) is evidenced by a sudden widening of the solar wind E-t plot (Figure 3 (b)) as well as a sharp drop in the solar wind velocity (Figure 3 (c)). Toward the bow shock, the upstream diffuse ions showed a gradual increase of their intensity, which is seen as the change of the color from yellow to dark red in Figure 3 (a). Figure 3 (c) shows that V^GEOTAIL decreased by 10-150 km/s from the VSWJWIND along with the increase of the diffuse ion intensity. To see the change of the diffuse ion intensity quantitatively, their energy density Em in the shock rest frame is calculated as,
EDI = J I j\m(v-vBS)2f(v)dv
(1)
where f(v) is the observed velocity space distribution function, and UBS is the bow shock velocity in the observers' frame. An assumption here is that the bow shock rest frame is nearly identical to the observers' rest frame, namely, that VQS is negligibly small and can be set to zero. That no multiple bow shock crossing was observed during this event is consistent with this assumption. The subpressure exerted by the diffuse ions, PJJI is calculated as (JDI — ^)EDI assuming JDI—5/3. Figure 3 (d) shows the change of Eoi (a solid line) as well as changes of the magnetic subpressure (PB, a dashed line), and the thermal proton subpressure of the solar wind {Psw, dots). EDJ and PB (and Psw weakly) showed gradual increases before their final jump at the bow shock. Figure 3 (e) shows the scatter plot of AVsm plotted against Em, which shows a clear negative correlation with the correlation coefficient 0.74. From the observed solar wind parameters, the change of the ram pressure of the solar wind, A(pswV^w) in the bow shock rest frame, is calculated to be ~ —(1-2) xlO~ 10 Pa just before the bow shock crossing. On the other hand, subpressure increases, APDI: APB, and APSW, are obtained as +0.8 X 1O~10 Pa, +0.2 x 10~10 Pa, and +0.1 X 10~10 Pa, respectively. (To calculate these increases we took values of PDI, PB, and Psw at t = 22.5 h as their 'base' values.) Thus their summation, APDI+APB+APSW ~ 1.1 x 10~10 Pa, roughly compensated the ram pressure change. This is what is expected for CRMSs. Predawn foreshock region Extensive studies of the nature of diffuse ions and related low frequency waves have been made mainly in the upstream region ahead of the terminator (X ~ 0 RE), and there have been only a few studies of upstream and bow shock regions well behind the terminator (X -C 0 RE) (e.g., Terasawa et al, 1985; Greenstadt et at., 1990; Sugiyama et al., 1995a; Bennett et al., 1997). Utilizing the GEOTAIL observations during the interval shown in Figure 1 (b), one can make comprehensive studies of this 'last frontier' of the upstream region. Figure 4 (a) shows the GEOTAIL orbit during the interval of 25 June - 2 July 1994 covering the predawn region (X, F)GSE ~ (—60,-70) RE to the nose upstream region (lygeotaiil ~ 10 RE, say). Figure 4 (b) shows the variation of the X component of the solar wind velocity throughout the interval. Figure 4 (c) shows the longitudinal angle of the IMF (ips) which was mainly within ±30° of the spiral direction (—45° or 135°). Figure 4 (d) shows that the intensity of protons (~ 36 keV) was maximized around the nose upstream region and decreased toward the predawn direction (Terasawa et al., 2003). Figures 5 (a)-(c) shows the ecliptic projection of the phase space distributions of diffuse protons. The corresponding anisotropy plots are shown in Figure 5 (d)-(f). These protons are more or less isotropic in the nose upstream region ((a) and (d)), but they have anisotropic pancake distributions around the IMF (perpendicular > parallel) in the predawn upstream region ((b)-(c) and (e)-(f)). In the predawn upstream region the proton intensity was maximized when the IMF was close to the spiral direction (not shown). These observations are consistent with the conclusions obtained from an ISEE-3 crossing of this region in 1983 (Terasawa et al., 1985), which had remained unconfirmed because of the incomplete energy and angular coverage of ISEE-3 observations. The new GEOTAIL observation gives a support for the early interpretation that the diffuse ions are predominantly produced in the nose upstream region and transported by the solar wind flow mainly
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Fig. 4. (a) The GEOTAIL orbit (25 June - 2 July 1994). Throughout the interval, the satellite was near the ecliptic plane (|^gse| < 5 R E ) . The nominal positions of the bow shock (BS) and the magnetopause (MP) are also shown. In panels (b)-(d), the following quantities are plotted against the Y coordinate of the GEOTAIL position (Ysse, R E ) : (b) The X component of the solar wind velocity (Vsv,x, km/s); (c) The longitudinal angle of the IMF (B, deg); and (d) Counting rates of sunward-flowing protons.
in the direction perpendicular to the IMF (namely, transport with the ExB drift motion). INTERPLANETARY PHENOMENA From the GEOTAIL observations in the solar wind, I discuss selected topics about interplanetary CME shocks and pickup He+ ions of the LISM (local interstellar matter) origin. For other topics such as galactic/anomalous cosmic ray particles and solar/corotaing energetic particles (e.g., Hasebe et al., 1994, 1995, 1997, 2001; Takashima et al., 1997; Kobayashi et al., 1998; Doke et al., 1999), type III solar radio bursts (Kasahara et al., 2001), and kinks or discontinuities (Whang et al., 1998; Kessel, Quintana, and Peredo, 1999; Nakagawa et al., 2000), I refer the reader to the original papers. Diffusive acceleration of electrons at interplanetary CME shocks During the period from late '70s to early '80s, when the DSA model was being established as the 'standard theory' of cosmic ray origin, there were close collaborations between astrophysical theorists and space experimentalists working on data from heliospheric shock environments (see reviews, e.g., Tsurutani et al., 1985a, 1985b; Terasawa and Scholer, 1989; Terasawa, 2001). However, not all aspects of DSA were covered by these early studies. One topic not well explored was the difference between behaviors of ions and electrons: While ion events have been extensively studied observationally and theoretically since the middle '80s (e.g., Lee, 1983; Kennel et al.,1986; Gordon et al., 1999; Igarashi et al., 2003), detailed analyses for electron events have become available only recently (e.g., Shimada et al., 1999). It is noted that shock-accelerated ions can excite right-handed polarized MHD waves via cyclotron resonant condition and be self-scattered. (These right-handed waves belong to the lowest frequency branch of the whistler mode.) While it might seem possible for shock-accelerated electrons to excite left-handed polarized MHD waves and be self-scattered similarly, severe cyclotron damping by thermal ions for these left-handed waves practically prevents these excitation and scattering processes. To circumvent this difficulty, Levinson (1992) proposed to consider a cyclotron resonant interaction between electrons and obliquely propagating whistler waves. However, since obliqueness reduces the interaction efficiency, high Alfven Mach numbers (>~40) are needed for the Levinson's process to work. Therefore, if this Levinson's process is the only mechanism by which whistler waves are excited, electron DSA could not be effective for typical heliospheric shocks whose Alfven Mach numbers are less than 10-20. -272-
Fig. 5. Panels (a)-(c) show the ecliptic cuts of the phase space distributions (PSD) of diffuse protons,
f(V),
in the velocity range of 1000-2869 km/s (5.2-43 keV) for (a) 00:51-00:54 UT on 2 July 1994 at ( X , Y , Z ) G S E = (+16.6, +9.4, +0.3) RE, (b) 20:14-20:17 UT on 28 June 1994 (at (X,Y,Z)GSE = (-14.2, -52.6, -0.2) R E , and (c) 19:35-19:38 UT on 25 June 1994 at ( X , Y , Z ) G S E = (-52.2, -68.7, +2.0) R E , The pseudo-color scale and
directions are given in the left of the figure. Directions of the IMF are shown by straight lines (marked 'B'), which are drawn to pass through Vsw (solid circles). Panels (d)-(f) show the traditional anisotropy plots of 36 keV protons for the same intervals as (a)-(c). A dark red spot (or tongue) below the center of panel (a) represents a contamination of solar wind heavy ions.
Figure 6 shows intensity variations of nonthermal (a) protons and (b) electrons around the passage of a moderately-strong (the Alfven Mach number MA ~ 6) interplanetary shock on 21 February 1994 (Baker et al., 1995; Shimada et al., 1999). This interplanetary shock was created in a CME event associating with a near-disk-center (N09 W02) flare at 0138 UT on 20 February 1994, about 32 hours before the shock arrival. The arrival of the shock front at GEOTAIL was identified at ~09:03 UT from the observations of thermal plasmas and magnetic field. Both nonthermal protons and electrons showed gradual increases between 06 UT and 09 UT, and reached their maxima at the shock arrival. These behaviors are consistent with what the standard DSA theory describes. As seen Figure 6 (a), ions showed a text-book like behavior for the diffusive shock acceleration events: exponential increases in the upstream region (before 0903 UT) and flat tops in the downstream region (after 0903 UT). In Figure 6 (b), electrons also had exponentially increasing upstream profiles+.The spatial diffusion coefficients corresponding to the observed upstream slopes were ~ severalxlO18 cm2 sec"1 both for electrons and protons. Since this value of the diffusion coefficient is by 2-3 orders of magnitude smaller than the typical interplanetary values (Palmer, 1982), local excitations of waves resonating with these particles are needed. Shimada et al. (1999) have confirmed that there was an enhancement level of low frequency (~ 0.01 Hz) MHD waves which are responsible for the scattering of ions. For electrons, they have also identified whistler wave enhancement in the region around the interplanetary shock. Figure 6 (c) shows the power spectrum of right-hand component of the upstream magnetic turbulence in the frequency range of 0.01-8 Hz. Above ~1 Hz, there were intermittent bursts of whistler waves which were created at the shock front or in the shock downstream region, and propagated upstream (Shimada et Note that the behaviors of protons and electrons at the shock front and in the downstream region were different: Below several keV there were jumps of the electron phase space densities at the shock. Further electrons showed general decreasing trends till 1400 UT. Shimada et al. (1999) explained these differences in terms of heating of suprathermal electrons at the shock front as well as the adiabatic cooling accompanying with the downstream plasma expansion.
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Fig. 6. The GEOTAIL observation in the interplanetary space during the interval of 5-14 UT on 21 February 1994 including an interplanetary shock passage at 09:03 UT. (a): variations of nonthermal proton flux intensities (arbitrary unit), at the energies of 0.055-0.11, 0.11-0.65, 0.65-1.5, and 1.5-4.0 MeV. (b): variations of nonthermal electron flux intensities (arbitrary unit), at the energies of 0.42-0.49, 0.85-1.0, 1.7-2.1, 4.2-5.0, and 8.5-10.2keV. (c): the right-hand component of the upstream magnetic turbulence in the frequency range of 0.01-8 Hz (6-9UT).
al., 1999). In the lower frequency range (~1 Hz), on the other hand, there was a more-or-less continuous component whose intensity showed a gradual increase toward the shock (the right edge of the figure). Recently, Nakata et al. (2003a, 2003b) has pointed out the importance of this gradually increasing component as the scatterer of low energy electrons, and argued that they were created in the nonlinear cascading process from the low frequency (~0.01 Hz) MHD waves excited by the shock-accelerated protons. This observation suggests that there is a 'proton-assisted' DSA process for nonrelativistic electrons by which we can avoid the Levinson's limitation. 'Cosmic-ray-mediated' interplanetary shock In the previous subsection, we have discussed the nonlinear effect of the DSA process at the bow shock (production of upstream diffuse ions) in terms of the 'cosmic-ray-mediated' shock (CRMS) property. It is of interest how we can find the similar property in interplanetary CME shocks. However, after the search of the CRMS property among several tens of interplanetary shocks observed in the interval of 1994-2001 by GEOTAIL, only one example has been so far identified (Terasawa et al., 1999). This rarity of the CRMS property makes a marked contrast with the bow shock observation: Essentially in all the quasi-parallel bow shock crossings where diffusive ions appear, the CRMS property, namely the partial deceleration of the incoming solar wind flow, can be identified. This difference between the bow shock and interplanetary shocks certainly comes from the fact that interplanetary shocks usually have smaller Alfen Mach numbers {MA <2-3, see e.g., Berdichevsky et al., 2000) than the bow shock (MA ~ 5-10). Only in some exceptionally large IPSs, MA can exceed 5-10. The unique example of the interplanetary CRMS was obtained on 21 February 1994, during which the evidence of electron DSA was obtained (previous subsection, Figure 6). Figure 7 shows an enlarged time profile over 4 hours: From 8 UT to 9 UT, just before the shock arrival (09:03 UT), gradual increases were seen in the magnetic field magnitude, Babs (the top panel), in the solar wind density, Nsw (the second panel), in the thermal proton temperature, Tsw,p (the third panel), and in the solar wind velocity, Vsw (the fourth panel). For example, Vsw showed a gradual increase of ~50 km/s ahead of the main velocity jump of ~400 km/s. (Note that the velocity increases in the observers' rest frame. In the shock rest frame, the gradual velocity change corresponds to the gradual deceleration of the upstream plasma flow.) The changes in these quantities are what are expected if this shock was a 'cosmic-ray-mediated' shock (CRMS). For quantitative discussion, Figure 7 (e) shows subpressures exerted by nonthermal electrons (250 eV-40 keV; blue), by nonthermal ions (70 keV-10 MeV; red), by magnetic field (black), and by thermal protons (blue, -274-
Fig. 7. Observational evidence of the CRMS nature at the interplanetary shock arrived at GEOTAIL on 21 February 1994; (a) the magnetic field magnitude, Babs, (b) the solar wind density N sw , (c) the proton temperature Tsw,p. (d) t n e s ° l a r w i n d velocity V s w x , and (e) subpressures (magnetic Pb, electron e, thermal proton pressure Psw,p, energetic protons HEP, and their sum).
dotted) as well as their sum (thick black). Summing up the subpressure increases from 8 UT to 9 UT, one obtains the value of ~1.0 xl0~ 1 0 Pa. On the other hand, the ram pressure of the solar flow (in the shock rest frame) was decreased by ~1.3 xl0~ 1 0 Pa during the same interval, which is fairly close to the above sum of subpressures. Taking into account of the uncertainty of the subpressure determination (~30 %), one can conclude that this interplanetary shock belonged to the category of CRMS. (Of course, this conclusion is based on the presumption that the observed changes of Babs, Nsw, Tsw,p and Vsw between 8 UT and 9 UT were not accidental, and the change of the shock propagation speed during the same interval was less than ~ 50 km/s. We need more CRMS samples of interplanetary shocks to make further quantitative discussion.) LOCAL INTERSTELLAR MEDIUM It is widely known that the heliosphere is filled with the pickup ions (PUIs) of the local interstellar medium (LISM) origin, which have been believed to be well pitch-angle scattered in the solar wind and have a spherical distribution function in the velocity space (e.g., Vasyliunas and Siscoe, 1976). Mobius et al. (1985) first reported the detection of He+ PUIs of the LISM origin using information from the AMPTE
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spacecraft (for the recent review, see, e.g., Gloeckler and Geiss, 2001). While previous observations have been made by mass-spectrometer-type sensors which resolve ion species (M/q), Noda et al. (2001) showed that simple electrostatic analyzers only with energy-per-charge separation, such as the LEP on GEOTAIL and the ion analyzer on the NOZOMI spacecraft, can be utilized for the study of He+ PUIs. A weak point of the observation by electrostatic analyzers, of course, is that there are contaminating solar wind ions (protons, alphas, and heavier ions) which often make unambiguous identification of He+ PUIs not easy. However, while the data from the mass spectrometers so far are limited to the 2D plane in the velocity space, the data from the electrostatic analyzers can provide information of 3D distribution function of He + PUIs.
Fig. 8. (a) A schematic illustration of the torus distribution of He+ PUIs of the interstellar origin, which is being viewed from the (+Vx, +Vy, +Vz) direction. In the part inside the torus there are contaminating solar wind ions (protons, alphas, and heavier ions), (b) An ecliptic projection of the observed distribution function, which is a cut along the conical surface +22.5° above the ecliptic plane. The observation of these torus He+ PUIs was made 16:01-16:47 UT on 4 December 2000 at (X,Y,Z)GSE ~ (4, -30, 3) RE- The green arrows show the average magnetic field direction during the observation interval, (\B\,
Following the Noda's works, Oka et al. (2002a) have presented a detailed study on the shape of the 3D distribution function of He+ PUIs with the GEOTAIL LEP dataset: In the solar wind intervals, which were selected as being free from the contamination of foreshock diffuse ions, unambiguous identification of He+ PUIs of the LISM origin was possible for ~ 57 % of the observation time. In these 57 %, the 3D spherical (isotropic) distribution functions are found for ~34 %, and 'torus-like' distribution functions (gyrotropic but not isotropic) are identified for ~23 % (Figure 8). Oka et al. (2002a) have further shown the status of the He+ PUIs, spherical or 'torus-like', depends on the level of the magnetic turbulence of the solar wind, high or low. Oka and Terasawa (2003) have further investigated the origin of these turbulence, and concluded that they are intrinsic to the solar wind and excluded the possibility of self-excitation by the PUIs themselves. This observation that the He + PUIs have a 'torus-like' shape in the velocity space for significant amount of times indicates that the pitch angle scattering process for these ions is not as much effective as previously thought, but is consistent with the recent observations of long scattering mean free path of the He+ PUIs as large as ~ 1 AU (Gloeckler et al., 1995; Mobius et al, 1998; Fisk et al., 1997; Schwadron et al., 1999). The PUIs have been considered as an effective source of the accelerated particles produced at various heliospheric shocks, such as the termination shock, CIR shocks, and planetary bow shocks. The direct evidence of the acceleration of He"1" PUIs at the earth's bow shock has been obtained by Wang et al. (from AMPTE observations, 1995) and Oka et al. (from GEOTAIL observations, 2002b). -276-
Fig. 9. (a)-(d) show the observation from 11:40 to 12:10 UT on 6 November 1997: (a), (b), and (c) are E-t plots for the electron sensor, the ion sensor, and the solar wind ion sensor of the LEP instrument, respectively, (d) shows the soft X-ray intensities at two wavelengths (1-8 A and 0.5-4 A) monitored by the GOES spacecraft (one minute average), (e) shows the Yokhok hard X-ray counts at three energy bands, 23-33 keV (black), 33-53 keV (green), and 53-93 keV (red) with arbitrarily scaled background counts of LEP ion sensor (blue), (f) shows the Yohkoh gamma ray counts at two energy bands, 2.1-2.4 MeV (red) and 4.0-7.2 MeV (dashed red) with arbitrarily scaled background counts of LEP ion sensor (blue). Note that the LEP 'data' are affected by the plasma sheet ions when the scaled counts became less than ~300 (panel (e)) or ~40 (panel (f)).
SOLAR FLARE EFFECTS During solar flares photons in various wavelength are emitted. Unexpectedly it is found that the electric field measurement (EFD) of GEOTAIL is affected most likely by extreme ultraviolet (EUV) photons enhanced during solar flares (Takei et al., 2003a). It is further found that the particle counters (microchannel plates for ions and channeltrons for electrons) of the LEP instrument are sensitive to the hard X-ray photons (Takei et al., 2003b) as shown below. On 6 November 1997 GEOTAIL was in the tail plasma sheet and counting rates of particles were moreor-less steady for the interval plotted in Figure 9 (a)-(c) (11:40-12:10 UT), except, during the ~2 minute interval around 11:53-54 UT, when energy independent count increases were seen in the LEP data from the ion sensor (Figure 9 (b)) and the solar wind ion sensor (Figure 9 (c)). Figure 9 (d) shows the simultaneous GOES soft X-ray observation, from which we have found that this event occurred during the initial phase of a large flare (X9.4/2B) at S18 W63 on the solar disk. The soft X-ray intensities reached their maxima around 11:54 UT. It is known that photon emission profiles at solar flares are energy dependent and are sharper for higher energy. In panels (e) and (f) of Figure 9 the LEP data (blue) are shown in the subinterval 11:50-12:00 UT, during which the Yohkoh hard X-ray and gamma-ray detectors showed bursty increases starting from ~11:52 UT with durations of several minutes (Yoshimori et al., 1999). In these panels it is seen that the LEP 'data' closely followed the hard X-ray light curve at 53-93 keV. Since the particle counters of the LEP are put inside the GEOTAIL satellite structure and separated from external radiations by Aluminum walls of thickness ~ several mm, the solar photons penetrating to these counters should have energy above several tens of keV. Yoshikawa (2003) experimentally confirmed that the particle counters of the same type as the LEP's are sensitive to hard X-ray to gamma ray photons with a quantum efficiency of ~ several %. Thus the detection of solar hard X-rays by the LEP sensors is not unreasonable happening. In addition to a case study for the 6 November 1997 event, Takei et al. (2003b) have made a statistical survey of solar flare events to see whether similar count enhancements are seen in the LEP data at other
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intense flares. It turns out that among 30 X-class flares occurring in the 1997-2000 season 11 flares gave detectable effects on LEP. Note that these 11 flares except the 6 November 1997 flare occurred when GEOTAIL was in the solar wind or magnetosheath. In these regions the ion temperature was relatively low so that there is an energy 'window' in which only solar hard X-ray photons could give detectable counts. For other 19 flares GEOTAIL was in the plasma sheet where the high temperature ions mask the 'signals', if any, of solar photons. (The solar photons on the 6 November 1997 event had quite high intensities so that they gave significant count enhancements to LEP even in the plasma sheet condition.) The lessons from the above study of the solar flare effect on the particle sensors are two-fold. One is experimental: For future missions coming closer to the sun, like the Bepi-Colombo mission to Mercury, solar hard X-rays have intensities inversely proportional to the distance from the sun, and can give significant noise levels to plasma measurements. Second is scientific: Since the 'duty cycle' for the solar EUV/hardX ray photon 'detection' by the EFD/LEP sensors or by similar plasma sensors is nearly 100 % without interruption by, e.g., shadowing by the earth, there arises possibility to utilize these solar 'signals' for scientific purpose. Recently, we have obtained an illustrative example: For three homologous flares on 24 November 2000 the data coverage of the Yohkoh hard X ray detector was incomplete because of the earth shadow. Since the three flares were continuously 'monitored' by GEOTAIL, the EFD 'data' of the solar EUV photons was found useful in analyzing these events (Takasaki et al., 2003). SUMMARY AND COMMENT While the GEOTAIL spacecraft project was originally planned for the detailed study of the magnetosphere and magnetotail physics, it has observed many interesting phenomena in the solar wind and the magnetosheath. In this review, we have focused on several selected topics. Firstly, we have reviewed GEOTAIL contributions to the understanding of the physical process in the foreshock and bow shock regions: bow shock origin of diffuse ions, quantitative assessment of the nonlinear back reaction of diffuse ion production ('Cosmic-Ray-Mediated' shock (CRMS) effect), and the transport process of diffuse ions toward the predawn foreshock region. Secondly, we have seen GEOTAIL observations of propagating interplanetary shocks (IPSs) ahead of CMEs. GEOTAIL have provided the first clear evidence of diffusive electron acceleration at an IPS, as well as the first observation of the CRMS effect at the same IPS. Thirdly, we have described observations of pickup interstellar He+ ions by GETAIL. Unique contribution of GEOTAIL is the first identification of phase space torus of these ions. Such a torus shape has been theoretically expected at the initial stage of the assimilation process of pickup ions to the solar wind, but never been identified observationally. Finally, we have seen 'direct observations' of solar flare extreme ultraviolet (EUV) and hard X-ray photons by GEOTAIL sensors. The electric field sensor identifies transient increases (duration ~ several minutes) of the sunward electric field near the peaks of solar flare X ray emission. It has been further shown that since the particle counters of the plasma instrument are sensitive to hard X-ray photons some of big solar events have been 'observed' as significant increases of the background counts of the plasma particle observations. It is hoped that the GEOTAIL mission continues to cover the declining phase of the current solar cycle 23 (May 1996 - 2006?) and further into the next solar cycle 24. Particularly, further detection of 'cosmic ray mediated' interplanetary shocks is highly desired. In this respect I note that a quite strong interplanetary shocks deteced on 29 October 2003 (average propagation speed ~ > 2000 km/s) was found to have 'cosmic ray mediated' shock features. Detailed studies of this interesting shock are now under way. ACKNOWLEDGEMENTS I am grateful to all the members of the GEOTAIL project for their contributions to the detailed studies of solar wind and interplanetary phenomena. I thank Drs. T. Kosugi, K. Shibata, and K. Watanabe for their assistance in comparing the LEP data with Yohkoh observations. Special thanks are also due to Drs. A. Nishida, T. Mukai, and M. Hoshino for various discussions and comments. This work is partially supported by Grant-in-Aid for Scientific Research, No. 13874055 and 14340066, from the MEXT (Ministry of Education, Culture, Sports, Science, and Technology of Japan).
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REFERENCES Anagnostopoulos, G. C , Phys. Scripta T52, 142, 1994. Baker, D. N., et al., J. Spacecraft Rockets 32, 1060, 1995. Bennett, I., et al., J. Geophys. Res. 102, 26927, 1997. Berdichevsky, D. B., et al., J. Geophys. Res. 105, 27289, 2000. Bonifazi, C , et al., J. Geophys. Res. 85, 3461, 1980a. Bonifazi, C , et al., J. Geophys. Res. 85, 6031, 1980b. Bonifazi, C , et al., J. Geophys. Res. 88, 2029, 1983. Christon, S. P., et al., Geophys. Res. Lett. 27 2433, 2000. Desai, M. I., et al., J. Geophys. Res. 105, 61, 2000. Diodato, L., and G. Moreno, J. Geophys. Res. 82, 3615, 1977. Dmitriev, A. V., et al., Terr. Atmos. Oceanic Sci. 13, 499, 2002. Doke, T., et al., J. Geomag. Geoelectr. 46, 713, 1994. Doke, T., et al. Adv. Space Res. 23, 487, 1999. Drury, L. O'C, and H. J. Volk, Astrophys. J. 248, 344, 1981. Ellison, D. C , E. G. Berezhko, and M. G. Baring, Astrophys. J. 540, 292, 2000. Fisk, L. A., et al., Geophys. Res. Lett. 24, 93, 1997. Formisano, V. S., and E. Amata, J. Geophys. Res. 81, 3907, 1976. Frank, L. A., J. Geomag. Geoelectr. 46, 23, 1994. Fuselier, S. A., Adv. Space Res. 15 43, 1995. Gloeckler, G., et al., Geophys. Res. Lett. 22, 2665, 1995. Gloeckler, G., and J. Geiss, Space Sci. Rev. 97, 169, 2001. Gordon, B. E., et al., J. Geophys. Res. 104, 28263, 1999. Gosling, J. T., et al., Geophys. Res. Lett. 5, 957, 1978. Greenstadt, E. W., et al., Geophys. Res. Lett. 17, 753, 1990. Hada, T., D. Koga, and E. Yamamoto, Space Sci. Rev. 107, 463, 2003. Hasebe, N., et al., Geophys. Res. Lett. 21, 3027, 1994. Hasebe, N., et al., J. Geomag. Geoelectr. 47, 1333, 1995. Hasebe, N., et al., Adv. Space Res. 19, 813, 1997. Hasebe, N., et al., J. Phys. Soc. Jpn. 70, 3167, 2001. Hayashi, K., et al., Geophys. Res. Lett. 21, 2907, 1994. Igarashi, K., et al., Adv. Space Res. this issue, 2003. Kasaba, Y., et al., J. Geophys. Res. 105, 79, 2000a. Kasaba, Y., et al., Geophys. Res. Lett. 27, 3253, 2000b. Kasahara, Y., H. Matsumoto, and H. Kojima, Radio Sci. 36, 1701, 2001. Keika, K., et al., J. Geophys. Res. in press, 2003. Kennel, C. F., et al., J. Geophys. Res. 91, 11917, 1986. Kessel, R. L., E. Quitana, and M. Peredo, J. Geophys. Res. 104, 24869, 1999. Ko, C.-M., Adv. Space Res. 15, 149, 1995. Kobayashi, M., et al., J. Phys. Soc. Jpn. 67, 3991, 1998. Kokubun, S., et al., J. Geomag. Geoelectr. 46, 7, 1994. Lee, M. A., J. Geophys. Res. 88, 6109, 1983. Levinson, A., Astrophys. J. 401, 73, 1992. Lin, R. P., C.-I. Meng, and K. A. Anderson, J. Geophys. Res. 79, 489, 1974. Matsui, H., et al., J. Geophys. Res. 102, 17583, 1997. Matsumoto, H., et al., J. Geomag. Geoelectr. 46, 59, 1994. Matsumoto, H., et al., Adv. Space Res. 20, 683, 1997. Matsuoka, A., et al., J. Geophys. Res. 105, 18361, 2000. Mobius, E., et al., Nature 318, 426, 1985. Mobius, E., et al., Geophys. Res. Lett. 13 1362, 1986. Mobius, E., et al., J. Geophys. Res. 103 257, 1998. Mukai, T., et al., J. Geomag. Geoelectr. 46, 669, 1994. -279-
Nagano, I., et. al, J. Geophys. Res. 108, 1224.1, 2003. Nakagawa, T., S. Kokubun, and T. Mukai, Adv. Space Res. 26, 811, 2000. Nakagawa, T., Y. Takahashi, and M. Iijima, Earth Planets Space 55, 569, 2003. Nakata, K., et al., Adv. Space Res. this issue, 2003a. Nakata, K., et al., Proc. 28th International Cosmic Ray Conf. (Tsukuba) 6, 3697, 2003b. Nemecek, Z., et al., Geophys. Res. Lett. 25, 1273, 1998. Noda, H., et al., Space Sci. Rev. 97, 423, 2001. Oka, M., et al., Geophys. Res. Lett. 29, 54-1, 2002a. Oka, M., et al., Geophys. Res. Lett. 29, 33-1, 2002b. Oka, M., and T. Terasawa, Adv. Space Res. this issue, 2003. Palmer, I. D., Rev. Geophys. 20, 335, 1982. Petrinec, S. M., et al., J. Geophys. Res. 102, 26943, 1997a. Petrinec, S. M., et al., Adv. Space Res. 20, 767, 1997b. Reiner M. J., et al., Geophys. Res. Lett. 24, 919, 1997. Safrankova, J., et al., J. Geophys. Res. 105, 25113, 2000. Sarris, E. T., and S. M. Krimigis, Geophys. Res. Lett. 15, 233, 1988. Schwadron, N. A., et al., J. Geophys. Res. 104, 535, 1999. Seon, J., et al., Geophys. Res. Lett. 26, 959, 1999. Shimada, N., Astrophys. Space Sci. 264, 481, 1999. Sibeck, D. G., et al., Geophys. Res. Lett. 24, 3133, 1997. Sibeck, D. G., et al., J. Geophys. Res. 104, 4577, 1999. Sibeck, D. G., et al., J. Geophys. Res. 105, 129, 2000. Sugiyama, T., et al., Geophys. Res. Lett. 22, 81, 1995a. Sugiyama, T., et al., J. Geomag. Geoelectr. 47, 1141, 1995b. Takasaki, H., et al., Astrophys. J. in press, 2003. Takashima, T., et al., Astrophys. J. 477, L l l l , 1997. Takei, Y., et al., Adv. Space Res. this issue, 2003a. Takei, Y., et al., Proc. 28th International Cosmic Ray Conf. (Tsukuba) 6, 3223, 2003b. Terasawa, T., et al., Geophys. Res. Lett. 12, 373, 1985. Terasawa, T., and M. Scholer, Science 244, 1050, 1989. Terasawa, T., et al., Proc. 26th International Cosmic Ray Conf. (Salt Lake) 6, 528, 1999. Terasawa, T., et al., Geophys. Res. Lett. 27, 3781, 2000. Terasawa, T., et al., Proc. 27th International Cosmic Ray Conf. (Hamburg) 9, 3620, 2001. Terasawa, T., Sci. Tech. Advanced Materials 2 461, 2001. Terasawa, T., et al., Proc. 28th International Cosmic Ray Conf. (Tsukuba) 6, 3705, 2003. Trattner, K. J., et al., J. Geophys. Res. 99 13389, 1994. Tsubouchi, K., et al., J. Geophys. Res. 105 25097, 2000. Tsuruda, K., et al., J. Geomag. Geoelectr. 46, 963, 1994. Tsurutani, B. T., and R. G. Stone (Eds), Geophys. Monograph 34, 1985a. Tsurutani, B. T., and R. G. Stone (Eds), Geophys. Monograph 35, 1985b. Vasyliunas, V. M., and G. L. Siscoe, J. Geophys. Res. 81, 1247, 1976. Wang, K., et al., Eos Trans. AGU 76 F465, SH41A-9, 1995. Whang, Y. C , J. Geophys. Res. 103, 6513, 1998. Williams, D. J., et al., J. Geomag. Geoelectr. 46, 39, 1994. Yoshikawa, I., Rev. Sci. Instr. in press, 2003. Yoshimori, M., et al., Proc. 26th International Cosmic Ray Conf. (Salt Lake) 6, 5, 1999. Zastenker, G. N., et al., Planet. Space Sci. 50, 601, 2002. Zhang, T.-L., K. Schwingenschuh, and C. T. Russell, Adv. Space Res. 15, 137, 1995. Zhang, Y., et al., J. Geophys. Res. 103, 20529, 1998. Zhang, Y., et al., J. Geophys. Res. 104, 449, 1999. Zong, Q. G., et al., Geophys. Res. Lett. 25, 4121, 1998. Zong, Q. G., et al., Geophys. Res. Lett. 26, 3349, 1999; -280-
WHISTLER WAVES IN UPSTREAM REGION OF INTERPLANETARY SHOCKS K.Nakata1, T.Terasawa 1. N.Shimada 2 , I.Shinohara3, Y.Saito 3 and T.Mukai3 1
University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo, 118-0033, Japan Communication Research Laboratory, 4-2-1 Nukui-Kitamachi, Koganei, Tokyo, 184-8795, Japan 3 Institute of Space and Astronautical Science, 3-1-1 Yoshinodai, Sagamihara, Kanagawa, 229-8510, Japan 2
ABSTRACT Whistler waves are considered to play an important role in the electron dynamics for the collisionless shock formation process at interplanetary shocks (IPSs). In this report, we analyze IPS events observed by GEOTAIL on 21 February 1994 and 15 July 2000, focusing on whistler wave properties in their upstream region extending the previous work of Shimada et al. (1999). In both events, we have identified the existence of whistler mode waves in the upstream region of the IPS as well as the tendency of the intensity increase toward the shock front. At the same time we have found the detailed features differing between these events: While the intermittent but clear wave bursts were found on 21 February 1994 event, the waves were more or less continuous on 15 July 2000 event.
Introduction When we consider diffusive shock acceleration (DSA) process for the particle accelerations in collisionless shocks, the wave-particle interactions, through which particles are scattered and accelerated, are essential. Generally, the main scatterer of this process is thought to be Alfven waves, which can interact with protons or high energy electrons. However, the low energy electrons below ~100 keV cannot resonate with these Alfven waves because of their shorter gyroradii. So the waves with shorter wave length are needed and the most probable candidate for these electrons is whistler waves. Several authors [Fairfield, 1974; Greenstadt et al.. 1991; Tsurutani et al., 2001] investigated whistler waves in upstream region of earth's bowshock in detail. Whistlers have been also identified around the bowshock of Jupiter [Moses et al., 1984], near comet Giacobini-Zinner [Kennel et al.. 1986], and around propagating interplanetary shocks (IPSs) [Coroniti et al.. 1982; Tsurutani et al., 1983]. Shimada et al. (1999), explicitly focusing on both the existence of whistler waves and the acceleration of low energy electrons, proved that observed whistler waves contribute to scattering/accelerating electrons through the DSA process. Here we expand Shimada's work to an IPS newly observed by GEOTAIL on 15 July 2000: we reconfirm existence of whistler waves through the polarization study and validity of electron scattering process with whistler waves by using quasi-linear theory. Instrumentation In the solar wind around 1AU, whistler waves have the frequencies between ~1 Hz and ~100 Hz. To analyze this frequency range, we utilize the power spectra obtained through the fluxgate data (FX, 1/16 sec sampling) from the MGF instrument aboard GEOTAIL [Kokubun et al., 1994]. The Nyquist frequency of the FX is 8.0Hz, thus covering the lowest range of the frequency of whistler mode. In our spectrum analysis, we divide the wave power into three components of right-hand, left-hand and compressional ('R'.'L' and ; C) in spacecraft frame. When using the FX data, the spectra are calculated every 64 seconds interval using 1024 data samples. According to the standard theory of DSA, the phase space distribution function of accelerated particles shows exponential increases toward the shock front in the upstream region, as exp(—xvs/D), where -281-
x is the distance from the shock front, vs the shock speed in the upstream plasma rest frame, and D the spatial diffusion coefficients at corresponding particle energies. By fitting / = /oexp[(£ — to)/t\] + C to the observed electron distributions [Mukai et al.,1994], we obtain Dois. Direct Observations of Wave Power Spectra In this section, we show the magnetic field data and power spectra for two events. From the results of both events, there seems to exist two different types of waves, one with the frequency range of l-4Hz appearing intermittently , and the other with 0.1-l.OHz appearing more or less continuously. The minimum variance analysis showed that these waves had propagation directions within 20 deg of the averaged magnetic field, and have dominant wave powers in the right-hand polarized mode. We conclude that they belong to the whistler mode branch since whistler waves having frequencies ~>0.1 Hz in observers' frame keep the right-hand polarization sense irrespective to their propagation direction (upstream/downstream) for the observed plasma parameters and solar wind condition. 21 February 1994 event Figure la shows MGF magnetic data and power spectra of 21 February 1994 event. The shock reached GEOTAIL at 0903 UT, when GEOTAIL was at (-27,61,-2) Re in the GSE coordinate. Remarkably, intermittent broad bursts around the frequency of 1.0-4.0Hz are seen in the R-mode panel in the upstream region of 0600-0900 UT. Duration of each burst is about several minutes to ten minutes. As Shimada et al. (1999) confirmed, these are upstream whistler waves propagating away from the shock front. On the other hand, in the lower frequency range of 0.1-1.0 Hz, all of R, L and C components show gradual increase toward the shock. Figure 2a shows the spectra of R and L components after taking average in the interval of 0830-0900 UT as well as the mean spectrum in the interval of 0500-0530 UT (a thin dotted line, 'bg'). We regard the spectrum of 0500-0530 UT as the background spectrum here. In this figure, it can be seen that the R component is more enhanced than the L component in the frequency range of 0.1-1.0 Hz and 1.0-8.0 Hz. While the latter enhancement reflects the whistler wave bursts which are clearly seen in Figure la, the former 0.1-1.0 Hz enhancement means R-mode is more enhanced during the gradual increase toward the shock front. The higher 1.0-8.0 Hz enhancement reflects the whistler wave bursts clearly seen in Figure la. The lower 0.1-1.0 Hz enhancement means that, while the wave density is gradually increasing toward the shock, all components are not equally enhanced, and R-mode waves are enhanced most. 15 July 2000 event Figure lb shows the case of 15 July 2000 event, in which IPS arrived at 1435 UT when GEOTAIL was at (25,6.4,-1.7) Re in the GSE coordinate. The horizontal straight lines seen between 2.0 and 3.0Hz are artificial noise and not natural. In this event, wave bursts above lHz are not identified. Instead, gradual increase of waves in the lower frequency range of 0.1-1.0 Hz is much stronger and enhancement of right-hand component is more clearly seen than in the previous event. In Figure 2b, which shows the spectra of R and L components after taking average in the interval of 1400-1430 UT as well as the background spectrum between 0900-0930 UT, the more enhancement of R-mode waves is high-lighted. R component is apparently more enhanced than L component during the gradual increase before the shock front. Estimation of Power Spectra from Electron Observation To make sure that the electrons are actually scattered and accelerated by observed upstream whistler waves, Shimada et al. (1999) confirmed that observed radial diffusion coefficient (= Dof,s) is consistent with the whistler wave energy density. In this section, we also examine whether observed diffusion coefficient is consistent with the whistler energy density in 15 July 2000 event. Shimada (1998) discussed the relation between spatial diffusion and wave intensity by using quasi-linear theory, from which ohs/COS
8TT
fR
=
8 n*e*W(kR) V ^ - - 2lnXmax
(2TT)3
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-
+ Xmax
^
+ 2lnXmax
-
XmJ
Figure la : 21 February 1994 The first three panels show absolute value, longitude and latitude of magnetic field. Next three contours respectively show dynamic spectrum of right, left and compressional component in spacecraft frame.
Figure 2a : 21 February 1994 The average spectra of R and L components in the interval of 0830-0900 UT. The plus signs are wave power derived from spatial diffusion coefficients.
Figure lb : 15 July 2000 The format of this figure is same as fig. la
Figure 2b : 15 July 2000 The format is same in fig.2a. The interval time is 1400-1430. The vertical peaks above lHz are artificial noise.
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where \& and Xmax a r e shock angle and maximum angle between wave direction and ambient magnetic field. The wave powers calculated from diffusion coefficients in each electron energy range are the plus signs in Figure 2a. The wave power estimated from observed spatial diffusion matches observed wave power within one order, and therefore they concluded that observed wave power of whistler wave is large enough to scatters/accelerates electrons. The wave powers we calculated from Dot,s of 15 July 2000 event are shown in Figure 2b. As in 1994 event, spatial diffusion and wave power are consistent, thus showing that whistler waves in upstream region of the shock can scatter/accelerate the low energy electrons in this event, too. Discussion In both events analyzed here, we have identified the existence of whistler waves as the enhancement of right-hand component waves. We have also confirmed that scattering of the lower energy electrons is mainly caused by whistler waves observed in upstream regions of the IPSs. At the same time, there appears some difference between these events: In 21 February 1994 event, there were bursty and intermittent whistler waves propagating away from the shock front in the frequency range of 1-4 Hz. On the other hand, in the lower frequency range of 0.1-1.0 Hz, the gradual and continuous enhancement of right-hand waves was seen both in 21 February 1994 and in 15 July 2000 events. The gradual increase is clearer in the latter event. This difference may reflect the difference of shock properties and physical processes of these IPSs. From the gradually increasing profile of wave power density and particle flux, it is suggested that the whistler waves are generated through the wave-wave interaction, namely, through the energy cascading from the Alfven waves to whistler waves. Generated whistlers in turn would contribute to the acceleration of low energy electrons. However, to confirm this scenario, the number of IPS events where whistler wave activity and the acceleration of low energy electrons were observed is still not enough because of their accidental occurrence. We are waiting for further interplanetary shock events. REFERENCES Coroniti, F. V., C. F. Kennel, and F. L. Scarf, Whistler mode turbulence in the disturbed solar wind, J. Geophys. Res., 87, 6029-6044, 1982. Fairfield D.H. Whistler waves observed upstream from collisionless shocks, J. Geophys. Res., 79, 1368-1378, 1974. Greenstadt, E. W, F. V. Coroniti, S. L. Moses, B. T. Tsurutani N. Omidi, K. B. Quest, and D. KraussVarban, Weak, wuasiparallel profiles of Earth's bow shock: a comparison between nummerical simulations and ISEE 3 observations on the far flank, Geophys. Res. Lett., 18, 2301-2304, 1991. Kennel, C. F., F. V. Coroniti, F. L. Scarf, B. T. Tsurutani, E. J. Smith, S. J. Bame and J. T Gosling, Plasma waves in the shock interaction regions at comet Giacobini-Zinner, Geophys. Res. Lett. 13, 921-924 1986. Kokubun, S, T.Yamamoto, M. H. Acuna, K.Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46, 7-21, 1994. Levinson, A., Electron injection in collisionless shocks, Astrophys. J., 401, 73-80, 1992. Melrose. D.B. "Kinetic Plasma Physics" in Saas-Free Advanced Course 24; Plasma Astrophysics, Kirk, J. G, D. B. Melrose, E. R. Priest, Springer-Verlag, 1994. Moses, S. L.. F. V. Coroniti, C. F. Kennel, and F. L. Scarf, Strog electron heat flux modes in Jupitar's foreshock, Geophys. Res. Lett. 9, 869-872, 1984. Mukai, T. S. Machida, Y. Saito, M. Hirahara, T. Terasawa, N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particle (LEP) experiment onboard the GEOTAIL satellite, J.Geomag. Geoelectr., 46, 669 - 692, 1994. Shimada, N, Diffusive Shock Acceleration Process of Electrons in the Solar Wind, in Their doctoral thesis, Dept. Earth and Planetary Physics Univ. of Tokyo, Tokyo, JAPAN, 1998. Shimada, N, T. Terasawa, M. Hoshino, T. Naito, H. Matsui, T. Koi, K. Maezawa, Diffusive shock acceleration of electrons at an interplanetary shock observed on 21 Feb 1994, Astrophy. Space Sci., 264, 481-488, 1999. Tsurutani, B. T, E. J. Smith and D. E Jones, Wave observed upstream of interplatnetary shocks., J. Geophys. Res., 88, 5645-5656, 1983. Tsurutani, B. T, E. J Smith, M. E. Burton, J. K. Arballo. C. Galvan, and Xiao- Yan Zhou, Oblique "1-Hz" whistler mode waves in an electron foreshock: The Cassini near Earth encounter, J. Geophys. Res., 106, 2001. -284-
WAVE-PARTICLE INTERACTION IN THE BASTILLE SHOCK OF YEAR 2000 K. Igarashi1, T. Terasawa1, T. Mukai2, Y. Saito2, K. Bamert3, R. Kallenbach4, and B. Klecker5 1
University of Tokyo, 7-3-1,Hongo,Bunkyo-ku, Tokyo, 133-0033, Japan Institute of Space and Astroatical Science, 3-1-1, Yoshinodai,Sagamihara,Kanagawa,229-8510, Japan 3 Physikalisches Institut, University of Bern,Siderstrasse 5, CH-3012 Bern, Switzerland 4 Internatinal Space Science Institute, Hallerstrasse 6, CH-3012 Bern, Switzerland 5 Max-Planck-Institut fur Extraterrestrische Physik, D-85740 Garching, Germany
2
ABSTRACT Diffusive shock acceleration process is regarded as one of the most important processes for astrophysical particle acceleration. We treat an ESP (energetic storm particle) event associated with an interplanetary shock (IPS) arrived at 1AU about 28 hours after a large flare/CME event happened in 14 July 2000 (known as the Bastille event). Using the magnetic field data from Geotail (X=24.9 Re, Y=7.4Re, Z=-1.8Re), and flux data of protons (~ several xlO2 keV) from SOHO (X=2.2Re, Y=-67.8Re, Z=12.3Re), we study the wave-particle interaction in the upstream region of this IPS.
INTRODUCTION From several hours before the arrivals of IPSs, fluxes of non-thermal particles (more than several tens ~ several hundreds of keV) often show smooth increases toward the maxima occurring near the shock fronts. This type of events, known as ESP since the time of IGY, is now considered as the in situ evidence of diffusive shock acceleration processes [Reames et al., 1996]. According to the standard Lee's theory [Lee, 1983] of ESP events, substantial energy exchange thorough the quasi-linear wave-particle interaction process is expected to occur in the upstream region of the IPS. While the quasi-linear process should depend on the shock parameters (e.g., the shock angle and Mach number), only limited information on such parameter dependence has been available so far (e.g., see [Kennel et al., 1986]). In this study, we show a case study for the relationship between particles and waves during one of the strongest ESP events. OBSERVATION For 10 years, Geotail spacecraft observed about 60 IPSs. The so-called Bastille event in 2000 we study here is one of the biggest events. On this report we use the data sets from the magnetic field experiment (MGF; Kokubun et al., 1994) aboard Geotail and the energetic particle experiment (HSTOF; Hovestadt et al., 1995 ) aboard SOHO. Combination of these Geotail and SOHO data sets is made because they were complementary during the Bastille event. (Geotail energetic particle data were not available during the event, and SOHO lacked the magnetic field measurement.) The shock arrival time at Geotail was 14:35:45 (UT) on 15 July 2000. Magnetic field, density, solar wind velocity and Alfven velocity of upstream region were 9.4 [nT], 5 [/cc], 633 [km/s] and 91.0 [km/s]. The shock speed in observer's frame was ~1100 [km/s], and shock angle was estimated to be ~46 [deg]. Figure 1 shows the magnetic field profile (magnetic intensity, 6 and ). The horizontal axis is 06:00 to 19:00 (UT) on 15 July 2000. It is seen in the panels of 6 and that the turbulent component enhanced near the shock front (from 12:00 to shock front). Figure 2 shows the proton differential flux observed by SOHO and magnetic field intensity observed by Geotail (whose time is shifted by -18 minutes taking into account -285-
Fi
Fig. 1. Magnetic field from 06:00 to 19:00 (UT) on 15 July 2000 observed by Geotail. Geotail (X=25Re, Y=6.8Re, Z=-1.6Re) observed the Bastille shock at 14:35:45 (UT).
g- 2. SOHO proton profiles. From upPer l i n e . energy ranges are 257keV, 333keV, 409keV, 521keV and 671keV. SOHO(X=202Re, Y=-67.8Re, Z=12.3Re) observed this IPS at 14:17:31. Magnetic field data are adjusted to SOHO's time. (Time lag is about 18 minutes.)
of the IPS propagation time between SOHO and Geotail). The vertical axis is differential flux (arbitrary unit), and the horizontal axis is from 06:00 to 19:00 (UT) on 15 July 2000. The energy range is 257keV, 333keV, 409keV, 521keV and 671keV. Each point represents the differential flux averaged for the interval ±15 minutes. In Figure 2, it is seen that the increases of the proton flux continued not only in the upstream region (before 14:17) but also in the downstream region (after 14:17). This behavior is explained in terms of the contributions from the diffusive shock acceleration as well as the second order turbulent acceleration in the downstream region [Bamert et al., 2002; Kallenbach et al., 2002]. In this report, we treat the former process focusing on the scattering process of energetic protons in the upstream region (13:00~14:00 at Geotail). According to the standard diffusive shock acceleration theory [e.g., Blandford and Ostriker, 1978], the flux / o f accelerated particles in the shock upstream region is given by,
/(«) = /-co + ( ^ - / c o ) e x p { v . £ ^ ^ d x ' } = /.00 + ( / o - / _ 0 0 ) e x p { - ^ }
(1)
where we assume that the steadiness of the acceleration process is realized. In (1) the x axis is taken parallel to the shock normal direction from the upstream to the downstream directions (with x = 0 at the shock front). /_oo is the background flux level at the far upstream limit x —> — oo, and /o the flux at the shock front. The upstream plasma comes into the shock front with the speed Vs (in the shock rest frame). Note that the spatial diffusion coefficient, D(x,E), generally depends on the position (a;) and the particle energy (E). We will use the spatially-averaged value D(E) in what follows. Since the flow velocity of the upstream plasma and the shock speed in the observer's frame are known, we can estimate D by fitting the functional form (1) to the observed flux time profile. The estimated values of D in the two lowest energy ranges (257 and 333keV) are listed in the second column of Table 1. We compare these D values with those from the wave observations in the next section. In other energy range, we do not estimate D because of the large statistical errors of the particle observation. WAVE PARTICLE INTERACTION We determine the wave number vector using minimum variance method [Sonnerup, 1967] for turbulent upstream magnetic field. The power spectrum are calculated from 3-sec averaged magnetic filed data. -286-
Figure 3 shows 1-hour averaged power spectrum (Transverse and Compressional component) of the solar wind magnetic field upstream of the Bastille shock. The vertical axis is power (P{f)) of the magnetic waves per frequency [nT2/Hz]. The horizontal axis is the observed wave frequency. The dashed lines are averaged for 10:00-10:59, and the solid lines are averaged for 13:00-13:59 just before the arrival of IPS. From these spectra, it is seen that the intensity of transverse magnetic field waves (0.01-0.1 Hz) just upstream of the IPS showed one order of magnetic enhancement over the value 3 hours away from the IPS.
Fig. 3. Power spectrum of transverse and compressional wave. The dashed lines are calculated for 10:00-10:59 data, and the solid lines are for 13:00-13:59. The former lines are for the interval ~4 hours before the shock arrival, while the latter lines are for the interval just before the shock arrival.
According to diffusive shock acceleration theory, the enhanced waves resonate with particles of upstream region. Through quasi-linear interaction, energy is exchanged one another. The condition of the cyclotron resonance is, k\\VR = OJ-QP.
(2)
where k\\ the wave number parallel to the averaged magnetic field, w the frequency of magnetic wave, ilp the proton cyclotron frequency, and VR the velocity of the resonating particles, respectively. The kinetic energy (ER) of resonating particle is given by
where mp is the mass of proton and Ec is the magnetic energy per thermal proton [Kennel!k,Petschek, 1966]. Between observed frequency wo(,s and frequency in the solar wind frame wsul, there is the Doppler relation, wojs — u>sw + k-V sw.
(5)
Using the dispersion relation ui = k-V A = ^V^cos^g ( VA is Alfven velocity and O^B ls the angle between wave number vector and magnetic field direction.), we can estimate the frequency of enhanced waves in solar wind frame. In this event d^B kept about 15-25 degree all the time (quasi-parallel propagation). When the k direction changes, the Doppler effect and the resonant condition change too. Considering these points, we -287-
estimate the frequency of the resonating waves with 257 and 333keV protons, and obtain 1.37-1.74Hz and 1.20-1.52Hz, respectively (Table 1, the fourth column). We analyze the waves in this range. Now we estimate the spatial diffusion coefficient, Dw from the magnetic power spectrum. The relation between diffusion coefficient and mean free path (A) is Dw = -\v,
(6)
v is the velocity of proton. The collision time (r) is given roughly 1 (SB-,-2
T
(7)
=n~M)
where fip the cyclotron frequency of proton (flp = 0.917[/s]). SB2 is the effective wave amplitude defined as P(f)f where / is the resonant frequency. For protons of 257 and 333keV, we estimate Dw accordingly, and tabulate the results in the third column of Table 1. Considering uncertainties in the estimation of k, wave intensities, etc., we regard that the agreement between Ds from two different methods (the second and third columns) is reasonable.
energy [keV] 257 333
D from flux 1019[cm2/s] 2.03 2.25
Dw from waves 10 19 [cm 2 /s]
2.58-3.10 2.70-3.93
resonated frequency 10-2[Hz] 1.37-1.74 1.20-1.52
velocity 103[km/s] 7.00 8.00
Table 1. Diffusion coefficients for protons (257 and 333keV) estimated from increase rates of proton fluxes (D), and from the power spectra of magnetic waves (Dw). Waves frequency resonating with protons.
SUMMARY In this study, we compared the particle and wave observations in the upstream region of the Bastille IPS, and found that the proton spatial distribution and the resonant wave intensity are consistently explained in terms of the quasi-linear theory, which is one of the basic ingredients of the standard diffusive shock acceleration theory. Further quantitative studies of the wave-particle interaction around IPSs are now under way. REFERENCES Bamert, K., et al., Solar Wind Ten, in press, 2003. Blandford, R. D., and Ostriker, J. P., Astrophys. J. 221, L29, 1978. Hovestadt, D., et al., Sol. Phys., 162, 441, 1995. Kallenbach, R., et al., Solar Wind Ten, in press, 2003. Kennel, C. F. and & Petschek, H. E., J. Geophys. Res., 71, 1, 1966. Kennel, C. F., et al., J. Geophys. Res., 91, 11917, 1986. Kokubun, S., et al., J. Geomag. Geoelectr, 46, 7,1994. Lee, M. A., J. Geophys. Res., 88, 6109, 1983. Reames, D. V., et al., Astrophys. J., 466, 473, 1996. Russel, C. T. and Hoppe, M.M., Space Sci. Rev., 33, 155, 1983. Shimada, N, Doctorial thesis, University of Tokyo, 1998. Sonnerup, B. U. 6.,and Cahill, Jr. L. J., J. Geophys. Res., 72, 171, 1967.
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PARTICLE-FIELD DYNAMICS IN THE SHOCK TRANSITION REGION N. Shimada1 and M. Hoshino2 'JSPS & Communications Res. Lab. Space Simulation G., 4-2-1 Nukui-Kitamachi, Koganei, Tokyo, Japan 2 University of Tokyo, Dept. Earth and Planetary science, 7-3-1 Hongo, Bunkyo-ku, Tokyo, Japan
ABSTRACT We carried out particle-in-cell simulation to investigate the mechanism for steady and unsteady shock front formation depending on Mach number (MA) and plasma / ? . When even under the same MA, depending on / ? , the electric and magnetic field structures in the shock transition region (STR) are much different. At lower P or/and higher MA condition, the shock front tends to be unsteady and electrons play more important role in the shock dissipation process through two-stream instability.
INTRODUCTION As widely noticed, the shock front does not always show steady propagation. Unsteadiness of the shocked profiles is known as shock cyclic behavior or shock reformation. At quasi-parallel shocks this unsteadiness is observed conspicuously and discussed from various points of view (e.g. Burgess, 1989; Lyu and Kan, 1990; Winske et al., 1990; Onsager et al. 1991; Scholer et al., 1993). For quasi-perpendicular shocks Quest (1986) and Krasnoselskikh et al. (2002) discuss on the unsteadiness of the shock front, but there are not many investigation about its mechanism. Except Krasnoselskikh (2002), since above previous papers mainly treated ion kinetics, we also include electron kinetics by carrying particle-in-cell simulation (Hoshino et al., 1992) and consider the mechanism for the steadiness and unsteadiness for the perpendicular shock viewed with the electron dynamics. Simulation conditions are as follows, mass ratio (M/m) of 20, and frequency ratio of the electron plasma to the cyclotron (CO VJ Q. ce) of 20, the shock front propagates to -x and the magnetic field is polarized to z. TIME VARIATION OF SHOCKED FIELDS Under the same MA, different /? (the ratio of the plasma pressure to magnetical pressure) makes different field reaction of the STR. Figure 1 presents propagation view of MA = 4.3 shock with/?= 0.5 (left two) and fi = 0.1 (right two). In each picture, the right is electric field (Ex) and the left is magnetic field (Bz). Vertical axis is time normalized by 2#7ft) pe and horizontal axis x is normalized by 2/TUo/£2ce. The shock front propagates from the lower right to the upper left in a steady manner for J3= 0.5 case, while for /3 = 0.1 case the shock front shows highly unsteady behavior sometimes with large-amplitude electric field structures. Depending on the phase of the ion reflection at the shock front, the two-stream instability (TSI) between the electron and the reflected ion is triggered, evolved and quenched. Corresponding to this cycle the Ex field shows enhancement and decay of these small but large-amplitude structures, which are seen as electron holes in the velocity space. The Bz field shows unsteady steepning process in a cyclic manner.
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PARTICLE-FIELD DYNAMICS IN THE SHOCK TRANSITION REGION (STR) Steady Shock Front Case In this subsection we consider the mechanism of the macro field structure under steady shock propagation. Left side of Figure 2 shows typical steady shock properties, from top, ion velocity in the x and y direction, electron velocity in the x and y direction, Ex, and Bz are plotted versus x. The velocity is normalized by injection speed u0, Ex and Bz are by motional electric field Ey0 (uoBz/c) and initial magnetic strength B o . The STR is around within the two vertical lines. Gradual depression followed by small hump is seen in Ex profiles within two lines. The decrease of Ex (upstream to (b2)) is resulted from the motion of Vy>0 reflected ion which must accompany the electron through E x B drift. The following part from (b2) to (bl) where ions with Vy<0 turn off the enhancement of the negative Ex to accompany the electron toward opposite direction. As illustrated in the right picture of Fig. 2 these points of the ion reflection (al), Ex recovery (bl), magnetically maximum (c), and bulk flow decelerated (d) are almost same. The upstream properties gradually merge into the downstream properties. No seed in the upstream for accumulation of the particle and enhancement of the magnetic field.
Fig. 1. Field variation of MA = 4.3 with /? e = /? i = 0.5 (left two) and /? e = P i = 0.1 (right two) shocks. In each pair, the right includes Bz and the left includes Ex. Vertical axis is time normalized by 17t I (H^ and horizontal axis x is by 27Zx\sj£}. cc.
Unsteady Shock Front Case In the unsteady shock front these points (al), (bl), (c), and (d) in the previous subsection are not the same anymore. Figure 3 presents the same picture with Fig. 2 but for M A = 11.0 and /?= 0.5 shock without the TSI evolution phase. Unlike the steady shock case narrow but large positive Ex field is generated due to the large momentum change of the reflected ion (as shown in the right of Fig. 3 as a shaded part in Ex profile). Corresponding to the reflected ion motion, Vy>0 part (left to (a)) makes Ex depression and Vy<0 part (right to (b)) makes Ex hump to accompany the electron into y direction. Bulk flow is already decelerated around (d) before the magnetically maximum point (c). Large Ex hump dams back the inflow ions as well as push forward backstreaming ion population. Accumulation of the particles helps steepning forward to the upstream. The wide reflection position (around (c) - (e)) also contributes to the accumulation of the particles there. Large velocity difference between the electron and the reflected ion triggers strong TSI and results strong dissipation around gray arrow seen in the ion trajectory picture (upper right figure). These facts lead new shock front formation around the position (e) ~ (b).
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Fig. 2. Simulation result (left, ion Vix and Viy, electron Vex and Vey, Ex, and Bz versus x) of the steady shock front of MA = 3.0 and /3- 0.5 and an illustration for the steady shock front mechanism (right). The STR is around within the two vertical lines in the left picture.
Fig. 3. The same format with Fig. 2 but for the unsteady shock front of MA = 11.0 and ft = 0.5. SUMMARY AND DISCUSSION The macro structure and mechanism of the steady weak shock and unsteady strong shock front are discussed. The width of the STR is about 3 in the unit 2^Tuo/£2Ce (or 1 in the unit uo/i^d), for all shocks investigated, where Q.c; is ion cyclotron frequency. In the unsteady shock front, in addition to the classical negative Ex region, strong narrow positive Ex region is observed corresponding to the abundant ion reflection and magnetic overshoot. Around the point where Ex has peak value (both of positive and negative), the ratio of the normalized Bz to Ex is about 1 independent of the plasma conditions. This means Uo (ion bulk motion speed) is comparable to cEx/Bz (= electron drift speed, VEB, where c is the light speed). This is consistent to the fact motion of the reflected ion (order Uo) sets up Ex field to forces the electron to obey in the y direction together in order to maintain quasi-neutrality by E X B drift. In Figure 4 the shock steadiness, macro-Ex profile, micro-Ex feature, and the ion to electron temperature ratio (Ti/Te) of the shocked plasma are overviewed under MA= 3-28 with J3= 0.5 (top) and p- 0.1 -291-
(bottom) conditions. With lower p, the Ex turbulence is triggered at lower MA because the TSI evolves easily. Above MA = 7.0 (for /?=0.5) we can see unsteadiness of the shock front as well as the enhancement of micro-scale turbulence (with coherent electron holes) in the Ex field, where at the same time Ti/Te begins to decrease because the nonlinear evolution of TSI enhances electron energization (e.g. Shimada and Hoshino, 2000; Hoshino and Shimada, 2002). As MA increases the shock front shows highly unsteady behavior but, it is interesting, in the much higher MA regime like > 20, unsteadiness becomes weak with showing strong turbulence at all times.
Fig. 4. Summary of the steadiness and shocked electric field profiles with the variation of Mach number (3 - 28) and j5 value (top for 0.5, bottom for 0.1). Shocked plasma temperature ratio Ti to Te is also presented. REFERENCES Burgess, D., Cyclic behavior at quasi-parallel collisionless shocks, Geophys. Res. Lett., 16, 345-348, 1989. Lyu, L. H. and J. R. Kan, Ion leakage, ion reflection, ion heating and shock-front reformation in a simulated supercritical quasi-parallel collisionless shock, Geophys. Res. Lett., 17, 1041-1044, 1990. Hoshino, M , J. Arons, Y. A. Gallant, A. B. Langdon, Relativistic magnetosonic shock waves in synchrotron sources - Shock structure and nonthermal acceleration of positrons, Astrophys. J., 390, 454-479, 1992. Hoshino, M , N. Shimada, Nonthermal electrons at high mach number shocks: electron shock surfing acceleration, Astrophys. J., 572, 880-887, 2002. Krasnoselskikh, V. V. , B. Lembege, P. Savoini, V. V. Lobzin, Nonstationarity of strong collisionless quasiperpendicular shocks: theory and full particle numerical simulations, Phys. Plasmas, 9, 1192-1209, 2002. Onsager, T. G., D. Winske, M. F. Thomsen, Ion injection simulations of quasi-parallel shock re-formation, J. Geophys. Res., 96, 21,183-21,194, 1991. Quest, K. B, Simulations of high mach number perpendicular shocks with resistive electrons, J. Geophys. Res., 91, 8805-8815, 1986. Scholer, M., M. Fujimoto, H. Kucharek, Two-dimensional simulations of supercritical quasi-parallel shocks: upstream waves, downstream waves, and shock re-formation, J. Geophys. Res., 98, 971-18,984, 1993. Shimada, N., M. Hoshino, Strong Electron Acceleration at High Mach Number Shock Waves: Simulation Study of Electron Dynamics, Astrophys. J. Lett., 543, L67-71, 2000. Winske, D. V. A. Thomas, N. Omidi, K. B. Quest, Re-forming supercritical quasi-parallel shocks. II - Mechanism for wave generation and front re-formation, J. Geophys. Res., 95, 18821-18832, 1990. E-mail address of N. Shimada: [email protected]
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ELECTROSTATIC QUASI-MONOCHROMATIC WAVES DOWNSTREAM OF THE BOW SHOCK: GEOTAIL OBSERVATIONS K. Shin1, H. Kojima1, H. Matsumoto1, and T. Mukai2 1
Kyoto University, Radio Science Center for Space and Atmosphere, Gokasho, Uji, Kyoto, 611-0011, Japan ^Institute of Space and Astronautical Science, 3-1-1 Yoshinodai, Sagamihara, Kanagawa, 229-8510, Japan
ABSTRACT The downstream region of the bow shock is a very turbulent region. Intense electrostatic waves in a wide frequency range can be observed. They are expected to be generated by electron beams which are accelerated in the transition region of the bow shock. In the present paper, we focus on the Electrostatic Quasi-Monochromatic (EQM) waves which are mostly observed in the downstream region. Similar waves have been reported by other spacecraft observations. However, their generation mechanism, energy source, and wave mode are still unclear. Geotail observations show good correlation of the EQM waves with the cold electron beam-like component which correspond to the electrons accelerated in the bow shock. In the present paper, we introduce the characteristics of the EQM waves and discuss their correlation with electron velocity distributions.
INTRODUCTION Intense plasma waves in a wide frequency range can be observed in the downstream region of the earth's bow shock. Many studies have been made on plasma wave generations in the bow shock and downstream regions. Fredricks et al. (1970) reported the first detailed observation results of the electric field spectra in the bow shock with OGO 5 satellite. Using the IMP 6 satellite data, Rodriguez et al. (1975) and Rodriguez (1979) classified the electrostatic waves observed in the near earth magnetosheath region into the Broadband Electrostatic Noise (BEN), narrowband electrostatic emission, and electron plasma wave. They are expected to be related to electron beams which are accelerated in a transition region such as the bow shock transition. The most common waves in the downstream region are the narrowband electrostatic waves whose frequencies are in the range between the electron plasma frequency fpc and the ion plasma frequency / p ;. Similar narrowband electrostatic emissions can be observed in the tail lobe region close to the magnetopause (e.g. Kojima et al., 2001), and foreshock regions (e.g. Anderson et al., 1981). However, no comprehensive model for the generation of these narrowband electrostatic emissions exists. Based on the observational studies in the last decades, the most plausible wave mode of this narrowband electrostatic emissions are believed to be the Doppler shifted ion acoustic waves (Rodriguez, 1979, Gallagher, 1985). However, the typical electron to ion temperature ratio in the downstream region is not satisfied with the theoretical generation condition (Tc/Ti > 1) of ion acoustic waves. In the present paper, we report the results of our analysis on the narrowband electrostatic emissions observed in the downstream region of the bow shock using Geotail plasma wave data. We address these narrowband electrostatic emissions as Electrostatic Quasi-Monochromatic (EQM) waves because of their characteristics of waveforms. We examine the electron velocity distributions and find good correlation between the EQM wave activities and the existence of cold beam-like components. Under the existence of two electron populations with different temperatures, electron acoustic wave mode can be destabilized. In this case, two electron populations mean the cold beam-like component and background electrons. In the present paper, we introduce the characteristics of the EQM waves and show one evidence that the EQM wave is the electron acoustic wave mode. OBSERVATIONS In this section, we show a typical example of the EQM waves. Figure l(a) shows the five minutes averaged electric field spectrum of the EQM waves observed in the downstream region of the bow shock during the period from 5:10 UT to 5:15 UT on January 15, 1995. The spectrum is generated by the Sweep Frequency Analyzer (SFA) data of the Plasma Wave Instruments (PWI) (Matsumoto et al., 1994) onboard the Geotail spacecraft. Intense narrowband spectra appear around 1 kHz which is in a range of the ion plasma frequency. Figure l(b) shows the Geotail orbit for the period of January 14, 1995 to January 16, 1995 in the Geocentric Solar Magnetospheric (GSM) coordinate
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Fig. 1. (a) Typical electric filed spectrum (averaged five minutes) of the EQM waves in the downstream region of the bow shock observed during the period from 5:10 UT to 5:15 UT on January 15, 1995 generated by the SFA data of the PWI. (b) Geotail orbit for the period of January 14, 1995 to January 16, 1995. system. A solid and dotted arrows show the directions of the ambient magnetic field and ion bulk flow, respectively. The magnetic field data were acquired from the Magnetic Field Instrument (MGF) (Kokubun et al., 1994). For this entire period, Geotail spacecraft was in the downstream region of the bow shock. Figure 2 shows the waveforms of the observed electric fields as well as their polarization on the antenna plane observed by the Wave Form Capture (WFC) receiver of the PWI. The top panel shows the snapshots of the parallel (upper panel) and perpendicular (lower panel) electric field components with respect to the ambient magnetic field observed during the period from 5:13:37.000 UT to 5:13:45.500 UT on January 15, 1995. This figure shows that the EQM waves are packet-like waves and their packet size are almost 0.5~l sec. We expand the observed waveforms in the left two panels of Figure 2(b) for the period of 40 msec, from 5:13:42.492 UT. The right panel displays the corresponding hodograph. From these panels, we find that the waveforms are quasi-monochromatic and the orientation of the electric field vector is almost parallel to the ambient magnetic field. Since the EQM waves are not accompanied with magnetic field components, they are purely electrostatic waves and their wave normal directions are parallel to the ambient magnetic field. Figure 3 shows the electron distribution observed by the Low Energy Particle (LEP) detectors (Mukai et al., 1994) onboard Geotail at 5:13:29 UT on January 15, 1995. The top panel shows contours of the phase space density of the electron velocity distribution function in theB — E x B velocity space plane, where E and B represent the ambient electric and magnetic fields, respectively. Here, the maximum contour level is 10~ 133 s 3 /m 6 and contour lines are drawn at an interval of 10° 2 s 3 /m 6 . The bottom panel shows the electron velocity distribution cut through f(v) along the ambient magnetic field. The positive velocity is referred to the ambient magnetic field direction. Dashed dotted lines show one count level of the instrument. In order to avoid confusing, we shaded the energy range contaminated by photoelectrons. bFigure 3 shows the existence of a beam-like component (displayed by arrows) around the velocity range of-2000km/s along the ambient magnetic field direction. We examined many EQM wave events and found that observations of cold electron beam-like components well correlate with those of the EQM waves. Since observed electron beam-like components are directed to the downstream away from the bow shock, they could be the result of the acceleration in the bow shock. DISCUSSION In the present paper, we show correlation of the EQM waves with the cold electron beam-like component in the downstream region of the bow shock. In last decades, the EQM waves in the downstream region of the bow shock have been considered as the Doppler shifted ion acoustic waves (e.g. Rodriguez, 1979, Gallagher, 1985). Gallagher (1985) proposed the model that the Doppler shifted ion acoustic waves are generated in the bow shock and ion foreshock regions and converted into the downstream region. However, in general, electrostatic waves are hard to propagate far from the generation region. Further, in the typical downstream region of the bow shock, the temperature ratio of electrons and ions is %IT\ < 1. Therefore, the ion acoustic wave is heavily damped. Thus, we conclude that the wave mode of the EQM wave is not the ion acoustic wave mode. The electron velocity distribution displayed in Figure 3 is similar to those observed by the VELA 4 satellite (Mont-
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Fig. 2. (a) Waveforms of the EQM waves of the parallel (top panel) and the perpendicular (bottom panel) components with respect to the ambient magnetic field observed during the period from 5:13:37.000 UT to 5:13:45.500 UT on January 15, 1995. (b) Expanded waveforms (left panels) and corresponding hodograph (right panel) for the period of 40 msec, from 5:13:42.492 •
Fig. 3. Electron velocity distribution tor the Jan1 5 j 1 9 9 5 e v e n t o b s e r v e d a t 5:1 3: 29 UT. Top and bottom
ls s h o w t h e c r o s s s e c t i o n in t h e
_ E x B p l a n e a n d t h e c u t through f(v) along the ambient magnetic field, respectively. The pos-
B
itive v d o c % is r e f e r r e d t 0 t h e a m b i e n t m a g n e t i c d i r e c t ion. The beam-like component is displayed by arrows.
fidd
gomery et al., 1970) and the ISEE 2 satellite (Feldman et al, 1982, 1983) in the transition region of the bow shock. They show that the directions of cold beam-like components are always toward the downstream region away from the bow shock. Thomsen et al. (1983) suggested the possibility that the ion acoustic wave mode and the electron acoustic wave mode are destabilized under the existence of the two electron populations (i.e. beam and flattop shape electron velocity distributions) in the bow shock. As we stated above, the plasma conditions in the downstream region of the bow shock are not suitable for the excitation of the ion acoustic mode. In contrast, the electron acoustic wave mode can be destabilized under the plasma conditions in the downstream region which have the hot background electrons and cold beam-like components shown in Figure 3. In order to examine the detailed correlation, we compare intensities of the EQM waves with the electron velocity distributions. Figure 4 shows the electric field spectra for the intense (solid line) and weak (dotted line) EQM wave events in the same format as Figure 1. The spectra for the intense and weak EQM wave events are averaged for the during the periods from 5:10 UT to 5:15 UT on January 15, 1995 and from 12:05 UT to 12:10 UT on January 7, 1995, respectively. Note that the frequency in the horizontal axis is normalized by the electron plasma frequency. The difference of the spectral intensities in the frequency range of f/jpC = 0.01 ~ 0.1 is clearly seen. Figure 5 is the same as bottom panel of Figure 3 but for corresponding two events in Figure 4. The electron velocity distributions of the intense and weak EQM wave events are observed at 5:13:29 UT on January 15, 1995 and 12:09:11 UT on January 7, 1995, respectively. In both electron velocity distributions, there exist the cold beam-like components (displayed by arrows). Note that the cold beam-like components appear in the opposite directions since the ambient magnetic field directions of the two events are opposite. These two cold beam-like components were in the downstream direction away from the bow shock. In Figure 5, the thermal velocities of the background electrons are different. We estimate the background electron thermal velocities using the same method as Feldman et al. (1983). The estimated background electron thermal velocities of the intense and weak EQM wave events are 3850 km/s and 3200 km/s, respectively. From the observed velocity distributions, it is hard to estimate the exact thermal velocity of the beam-like components. However, their thermal velocities seem to be in the same order. Since the linear growth of the electron acoustic wave mode depends on the temperature ratio of cold and hot electrons, we can expect that the difference of the background electrons lead to the difference of the EQM wave intensities shown in Figure 4.
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Fig. 4. Electric field spectra for the intense (solid line) and weak (dotted line) EQM wave events in the same format as Figure 1 observed during the periods from 5:10 UT to 5:15 UT on January 15, 1995 and from 12:05 UT to 12:10 UT on January 7, 1995, respectively. Frequency in the horizontal axis is normalized by the electron plasma frequency.
Fig. 5. Electron velocity distributions for corresponding two events in Figure 4 in the same format as bottom panel of Figure 2. The electron velocity distributions of the intense and weak EQM wave events are observed at 5:13:29 UT on January 15, 1995 and 12:09:11 UT on January 7, 1995, respectively.
In order to confirm the wave mode of the EQM waves theoretically, we conducted the linear dispersion analysis using the realistic parameters (not shown in the present paper). Our preliminary results show the electron acoustic wave mode is unstable under the existence of the two electron populations shown in the present paper. The detailed linear analysis results of the EQM wave and their parametric dependence will appear in a future publication. ACKNOWLEDGMENT We would like to thank the ISAS/NASA Geotail mission project team for their support. This research was supported by grant-in-aid 15204044. REFERENCES Anderson, R. R., G. K. Parks, T. E. Eastman, D. A. Gurnett, L. A. Frank, Plasma waves associated with energetic particles streaming into the solar wind from the earth's bow shock, J. Geophys. Res., 86, 4493, 1981. Feldman, W. C , S. J. Bame, S. P. Gary, J. T. Gosling, D. J. McComas, M. F. Thomsen, G. Paschmann, N. Sckopke, M. M. Hoppe, C. T. Russell, Electron heating within the earth's bow shock, Phys. Rev. Lett., 49, 199, 1982. Feldman, W. C , R. C. Anderson, S. J. Bame, S. P. Gary, J. T. Gosling, D. J. McComas, M. F. Thomsen, G. Paschmann, M. M. Hoppe, Electron Velocity Distribution Near the Earth's Bow Shock, J. Geophys. Res., 88, 96, 1983. Fredricks, S. A., G. M. Crook, C. F. Kennel, I. M. Green, F. L. Scarf, P. J. Coleman, C. T. Russell, OGO 5 observations of electrostatic turbulence in bow shock magnetic structures, J. Geophys. Res., 75, 3751, 1970. Gallagher, D. L., Short-Wavelength Electrostatic Waves in the Earth's Magnetosheath, J. Geophys. Res., 90, 1435, 1985. Kojima. H., M. Ashour-Abdalla, W. R. Paterson, H. Matsumoto, L. A. Frank, R. R. Anderson, R. L. Richard, S. Kokubun, T. Yamamoto, Generation of the narrowband electrostatic noise in the geomagnetic tail: Geotail observation, J. Geophys. Res., 106, 8483, 2001. Kokubun, S., T. Yamamoto, M. H. Acuna, K. Hayashi, K. Shiokawa, H. Kawano, The GEOTAIL Magnetic Field Experiment, J. Geomag. Geoelectr., 46, 7, 1994. Matsumoto, H., I. Nagano, R. R. Anderson, H. Kojima, K. Hashimoto, M. Tsutsui, T. Okada, I. Kimura, Y. Omura, M. Okada, Plasma Wave Observations with GEOTAIL Spacecraft, J. Geomag. Geoelectr., 46, 59, 1994. Montgomery, M. D., J. R. Asbridge, S. J. Bame, VELA 4 plasma observations near the earth's bowshock, J. Geophys. Res., 75, 1217, 1970. Mukai, T., S. Machida, Y. Saito, M. Hirahara T. Terasawa, N. Kaya, T. Obara, M. Ejiri and A. Nishida, The Low Energy Particle (LEP) Experiment onboard the GEOTAIL Satellite, J. Geomag. Geoelectr, 46, 669, 1994. Rodriguez, P., D. A. Gurnett, Electrostatic and Electromagnetic Turbulence Associated with the Earth's Bow shock, J. Geophys. Res., 80,19,1975. Rodriguez, P., Magnetosheath electrostatic turbulence, J. Geophys. Res., 84,917,1979. Thomsen, M. F., H. C. Barr, S. P. Gary, W. C. Feldman, T. E. Cole, Stability of Electron Distribution Within the Earth's Bow Shock, J. Geophys. Res., 88, 3035, 1983.
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GEOMAGNETIC ACTIVITY DEPENDENCE OF OCCURRENCE PROBABILITY AND SPATIAL DISTRIBUTION OF UPSTREAM EVENTS K. Keika1, M. Nose 2 , S. P. Christon3, and R. W. McEntire4 1
Department of Geophysics, Graduate School of Science, Kyoto University, Kitashirakawa-Oiwakechou, Sakyou-ku, Kyoto 606-8502, Japan 2 Data Analysis Center for Geomagnetism and Space Magnetism, Graduate School of Science, Kyoto University, Kitashirakawa-Oiwakechou, Sakyou-ku, Kyoto 606-8502, Japan 3 Focused Analysis and Research, Columbia, MD 21044, USA 4 The Johns Hopkins University Applied Physics Laboratory, 111000 Johns Hopkins Road, Laurel, MD 20723-6099, USA ABSTRACT We investigated upstream events observed by the ion composition system (ICS) sensor of the energetic particles and ion composition (EPIC) instrument on board the Geotail spacecraft. We examined how occurrence probability and spatial distribution of upstream events depend on the geomagnetic activity. The results showed that the upstream events were observed more frequently in the dawn side during intense geomagnetic activity in particular. We also analyzed carbon-nitrogen-oxygen ions during the upstream events. From the above results we discuss origin of the upstream energetic ions.
INTRODUCTION Many researchers have studied upstream energetic ion events for more than two decades since the first observation by Asbridge et al. (1968). There are two origins for the upstream energetic ions (> 50 keV) observed in the upstream of the Earth's bow shock and in the magnetosheath. They can be solar wind ions accelerated at the Earth's bow shock by either the Fermi acceleration mechanism at the quasi-parallel shock or the shock drift acceleration mechanism at the quasi-perpendicular shock (i.e., solar wind origin). They can be also a result of leakage from the magnetosphere where they are accelerated (i.e., magnetospheric origin). However, it is yet to be identified which origin is dominant. Sarris et al. (1978) and Scholer et al. (1981) suggested that energetic ions of magnetospheric origin are accompanied by bursts of energetic (> 75 keV) electrons and highly energetic (> 300 keV) protons, while that bow shock associated ions are not accompanied by such bursts and protons. Statistical analyses based on their criteria were made by Anagnostopoulos et al. (1999, 2000). They found that the occurrence probability of upstream energetic ion (50-220 keV) events accompanied by the relativistic (> 220 keV) electrons is more than 80 %. Highly energetic (> 300 keV) ions were observed over 90 % of the upstream ions. These results suggest that occurrence probability of upstream events of magnetospheric origin reaches as high as ~ 80 %. They also found a dawn-dusk asymmetry of spatial distribution in the upstream events that are accompanied by the relativistic electrons and the highly energetic ions. The upstream events on the afternoon side had lower intensities than those on the morning side. In this study we analyze occurrence probability and spatial distribution of upstream events from a viewpoint of their dependence on geomagnetic activity. There are a few studies that focused on dependence of upstream events on geomagnetic activity (Lin et al., 1974; Mitchell and Roelof, 1983; Anagnostopoulos et al., 1999). Though these studies used the Kp index, the time resolution of the Kp index (i.e., three hours) is not high enough to examine
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geomagnetic activity dependence of upstream events whose duration is from a few tens of minutes to nearly one hour. Therefore in this study we use 1 -min values of the SYM-H index, which is essentially the same as the Dst index, though it is calculated by using 1-minute values from different sets of stations (Iyemori, 1990). Using this 1-min index allows us to know geomagnetic activity during the periods better corresponding to occurrence time of upstream events.
DATA SET We used energetic ion flux data obtained by the ion composition system (ICS) sensor of the energetic particles and ion composition (EPIC) instrument (Williams et al., 1994) on board the Geotail spacecraft. The ICS sensor measures ion flux and ion mass. The energy band ranges from 58 keV to 3005 keV for protons which is divided into nine energy channels (P2—P10). The EPIC/ICS data are sectored into 16 azimuthal angular bins. There are two identical telescopes for measuring ions; both have a field of view of ~ 30° in polar direction, and the north (south) telescope is centered 23 degrees above (below) the spin plane which is nearly parallel to the ecliptic plane. In this study we used the averages of ion flux data obtained by the two telescopes, which are regarded as ion flux flowing nearly parallel to the ecliptic plane. The sampling rate of the flux data is 96 s. OBSERVATIONS
Fig. 1. Examples of upstream events observed in the upstream region of the Earth's bow shock.
Figure 1 shows proton flux data obtained by the P3 channel (77-107 keV) of the EPIC/ICS sensor at XGSE ~30 R E and the SYM-H index for the period 0300-0600 UT on August 18, 1999. The top four panels present proton flux in the dawnward, sunward, duskward, and tailward directions. The bottom panel displays a variation of the SYM-H index. The dawnward and sunward flux increases suddenly by more than two orders around 0410 UT, 0450 UT, and 0500 UT. In the first event the duskward flux increases around 0420 UT. These three events occurred during geomagnetically active period (SYM-H < - 3 0 nT).
STATISTICAL ANALYSIS We examined upstream events measured by the P3 channel of the EPIC/ICS instrument from 1999 to 2001 during which the Geotail satellite is in the near-Earth orbit of 9x30 RE- We analyzed how their occurrence probability and spatial distribution depend on geomagnetic activity. We identified an upstream event when flux was enhanced by more than two orders within ten minutes. The upstream region was defined by XQSE > 0 and 15 RE< V-^GSE + ^GSE —31 RE- The upstream region was divided into four radial bins and eight local time bins, resulting in each mesh having a radial distance extent of 4 RE and a local time width of 1.5 hours. The occurrence probability is calculated as the total duration of the events divided by the traveling time of the satellite in each mesh. Figure 2 shows spatial distribution of occurrence probability and its dependence on the SYM-H index. The occurrence probability is displayed in a gray scale. The center panel of Figure 2 presents the distribution including all
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Fig. 2. Occurrence probability and spatial distribution of upstream events. The center panel shows the probability for all SYM-H conditions. The left and right panel display the probability during SYM-H > 0 nT and SYM-H < - 3 0 nT, respectively.
Table 1. Detection rate of CNO ions during upstream events Detection rate of CNO ions YEAR 72.12% 1999 71.62% 2000 73.44 % 2001 72.17% TOTAL (1999-2001)
SYM-H conditions. The left and right panels present the distribution during SYM-H > 0 nT and SYM-H < - 3 0 nT, respectively. Occurrence probability for each mesh ranges from 5 % to 20 % in the all geomagnetic condition (SYMH ALL), 0 % to 15 % in the quiet geomagnetic condition (SYM-H > 0 nT), and 10 % to over 40 % in the disturbed geomagnetic condition (SYM-H <—30 nT). Average of occurrence probabilities over all meshes was 8.5 % in the all geomagnetic condition, 2.5 % in the quiet geomagnetic condition, and 12.6 % in the disturbed geomagnetic condition. Spatial distribution of occurrence probability in the all geomagnetic condition showed a slight dawn-dusk asymmetry with more events on the dawn side. The dawn-dusk asymmetry also appeared in the disturbed geomagnetic condition, although events occurred uniformly in the quiet condition. The asymmetry in the disturbed condition was much stronger than that in the all geomagnetic condition. These statistical results indicate that (1) upstream events can be seen more frequently as the SYM-H index becomes smaller, and (2) upstream events occur more frequently in the dawn side during geomagnetically disturbed periods in particular. DISCUSSION We have examined occurrence probability and spatial distribution of the upstream events, focusing on their dependence on the geomagnetic conditions. The result shows that the probability was higher in the dawn side in the geomagnetically disturbed condition. We can conclude that particles of magnetospheric origin are dominant in the upstream events during geomagnetically disturbed periods. Particles in the magnetotail are accelerated by a reconnection process when storms and/or substorms occur. These high-energetic particles can leak out of the magnetosphere and travel toward the dawnside upstream region. However, it might be possible for solar wind particles to experience acceleration more frequently at the dawnside bow shock in the geomagnetically disturbed condition. The first-order Fermi acceleration occurring at the dawnside bow shock can be enhanced by the shock and/or disturbances in the solar wind (e.g., high speed of the solar wind) which result in greater geomagnetic activity. Therefore, in order to confirm our conclusion mentioned above (that is, particles of magnetospheric origin are dominant in upstream events), we examined carbon-nitrogen-oxygen (CNO) ions in the energy range of 221 keV to
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275 keV measured by the M3 channel of the EPIC/ICS instrument. The M3 channel can measure only CNO ions. According to the previous studies, CNO ions detected by the EPIC/ICS instrument are O + and N + dominant (Jacquey et al., 1994; Lui et al., 1994). We analyzed existence of CNO ions during the upstream events. We consider ion counts more than three as a flux enhancement of CNO ions. Table 1 displays the detection rate of CNO ions during the upstream events. The rate reached more than 70 % for each year, suggesting that more than 70 % ions of upstream events were magnetospheric origin. This is almost consistent with the rate of magnetospheric origin calculated by Anagnostopoulos et al. (1999). We came to the reasonable conclusion as follows. More than 70 % of energetic (77 keV-107 keV) ions in the upstream of the Earth's bow shock and in the magnetosheath are of magnetospheric origin. Ions accelerated in the magnetosphere are leaking into the magnetosheath and traveling toward the dawnside upstream region along the interplanetary magnetic field (IMF) during the geomagnetically disturbed periods (i.e., storms and/or substorms periods). The ions of magnetospheric origin are drifting westward in the near-Earth magnetosphere and leaking out from the duskside magnetopause because of their large gyroradius. We will examine IMF orientation during the upstream events in future study to determine from which part of the magnetopause ions can leak out. A statistical analysis with larger data set will be also conducted. ACKNOWLEDGEMENTS We are grateful to D. J. Williams for making EPIC data available. We also thank all members of Geotail program team for their assistance. We acknowledge T. Iyemori for providing the SYM-H index. We are indebted to T. Terasawa, K. Takahashi, and S. Ohtani for their helpful comments. This work was partly supported by the Atmospheric Science Division of the National Science Foundation (grant ATM-0000225) to JHU/APL and the Sasagawa Scientific Research Grant from The Japan Science Society. REFERENCES Anagnostopoulos, G. C , G. Kaliabetsos, G. Argyropoulos, and E. T. Sarris, High energy ions and electrons upstream from the Earth's bow shock and their dependence on geomagnetic conditions: Statistical results between years 1982-1988, Geophys. Res. Lett., 26, 2151-2154, 1999. Anagnostopoulos, G. C , G. Argyropoulos, and G. Kaliabetsos, Spatial distribution of upstream magnetospheric > 50keV ions, Ann. Geophys., 18, 42-46, 2000. Asbridge, J. R., S. J. Bame, and I. B. Strong, Outward flow of protons from the earth's bow shock, J. Geophys. Res., 73, 5777-5782, 1968. Iyemori, T., Storm-time magnetospheric currents inferred from mid-latitude geomagnetic field variations, J. Geomagn. Geoelectr., 42, 1249-1265, 1990. Jacquey, C , D. J. Williams, R. W. McEntire, A. T. Y. Lui, V. Angclopoulos, S. P. Christon, S. Kokubun, T. Yamamoto, G. D. Reeves, and R. D. Belian, Tailward energetic ion streams observed at ~ 100 RE by GEOTAIL-EPIC associated with geomagnetic activity intensification, Geophys. Res. Lett., 21, 3015-3018, 1994. Lin, R. P., C. -I. Meng, and K. A. Anderson, 30 to 100 keV protons upstream from the earth's bow shock, J. Geophys. Res., 79, 489-498, 1974. Lui, A. T. Y, D. J. Williams, S. P. Christon, R. W. McEntire, V. Angelopoulos, C. Jacquey, T. Yamamoto, and S. Kokubun, A preliminary assessment of energetic ion species in flux ropes/plasmoids in the distant tail, Geophys. Res. Lett., 21, 3019-3022, 1994. Mitchell, D. G., and E. C. Roelof, Dependence of 50-keV upstream ion events at IMP-7 and 8 upon magnetic field bow shock geometry, J. Geophys. Res., 88, 5623-5634, 1983. Sarris, E. T., S. M. Krimigis, C. O. Bostrom, and T. P. Armstrong, Simultaneous multispacecraft observations of energetic protons bursts inside and outside the magnetosphere, J. Geophys. Res., 83, 4289, 1978. Scholer, M., D. Hovestadt, F. M. Ipavich, and G. Gloeckler, Upstream Energetic Ions and Electrons: Bow ShockAssociated or Magnetospheric Origin?, J. Geophys. Res., 86, 9040-9046, 1981. Williams, D. J., R. W. McEntire, C. Schlemm II, A. T. Y. Lui, G. Gloeckler, S. P. Christon, and F. Gliem, GEOTAIL Energetic Particles and Ion Composition Instrument, J. Geomagn. Geoelectr., 46, 39-57, 1994. E-mail address of K. Keika
[email protected]
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ENHANCEMENT OF SUNWARD DOUBLE-PROBE ELECTRIC FIELDS OBSERVED BY GEOTAIL DURING THE SOLAR FLARE Y.Takei1, T. Terasawa1, M. Nakamura1, T. Mukai2, H. Hayakawa2, A. Matsuoka2, H. Takasaki3, and K. Shibata3 1
Department of Earth and Planetary Science, University of Tokyo, Tokyo, 113-0033, Japan 2 Institute of Space and Astronautical Science, Sagamihara, 229-8510, Japan 3 Kwasan Observatory, Kyoto University, Kyoto, 607-8471, Japan ABSTRACT
We report significant enhancement of sunward electric field, i?x, detected by GEOTAIL during the November 24, 2000 flare (X2.3/2B). We found a good correlation of the time profile of Ex with that of hard X-ray (larger than 23keV) flux observed by YOHKOH. In addition, we have found 100 events that show similar enhanced double-probe E-* in association with hard X-ray flares for which results of a statistical study are also presented.
INTRODUCTION Solar flares are an explosive phenomenon that takes place in the atmosphere of the sun, which releases a big amount of energy in the form of high-temperature plasma and energetic particles (plus mass motions), resulting in detection of the fiare emission in various wavelengths e.g. from radio, visible, EUV, soft and hard X-rays, through gamma-rays. The solar flare may affect some types of spacecraft measurements. For example, Brace et al. [1988] reports that the photoelectron currents detected by the Langmuir Probe onboard the Pioneer Venus Orbiter are enhanced and that the current enhancement is caused by the solar EUV emission. Also the change in energy spectrum of the solar flux causes the change in energy spectrum of photoelectrons emitted from the surface of the probes and the spacecraft [Grard, 1973], and results in the change of the spacecraft potential [Schmidt and Pedersen, 1987]. In this paper we report an event of significant enhancement of sunward electric field (£*) offset, detected by EFD-P onboard GEOTAIL which seems to occur in association with the November 24, 2000, flare (X2.3/2B, flare location N22 W07). We also present a statistical study of similar events. The EFD-P utilizes double-probe technique [Tsuruda et al., 1994], which measures the DC electric field vector (E*, Ey) in the satellite spin equator (within ~ 2° of the ecliptic plane) by two sphere probes which deployed 50 meters from the spacecraft sidewall. It is noted that this technique is sensitive to the plasma environment around the spacecraft. For example, Takei et al. [2004] reports that sunward electric field, Ex, measured by EFD-P has a positive offset because of the asymmetrical distribution of the photoelectron cloud around the spacecraft. The spacecraft potential, Vs/ci with respect to the surrounding plasma is also monitored, and usually ~ several volt in the solar wind environment. In the tenuous plasma Vg/c depends on the current balance between photoelectrons emitted from the sunlit surface of the spacecraft and ambient electrons coming into the spacecraft [Garrett, 1981]. The average energy of the photoelectrons is several eV [Grard, 1973], and the low-energy photoelectrons confines to the vicinity of the spacecraft. Recently multi-component photoelectrons yielded from the surface of the spacecraft in space are reported in several literatures [Pedersen, -301-
Fig. 1. Time profiles of GEOTAIL electric field measurement, GOES soft X-ray flux, and YOHKOH hard X-ray flux of the November 24, 2000 solar flare (X2.3/2B). Panel (a) shows time profile of sunward component of the double-probe electric field, -Ex (a solid line), —(V x B ) x electric field (a dotted line), and spacecraft potential Vg/c (a dashed line) obtained by GEOTAIL. The scales of the electric field (mV/m) and the potential (V) are given on the left and right sides of the panel, respectively. Panels (b) and (c) show time profiles of solar soft X-ray fluxes (1-8 A and 0.5-4 A) monitored by GOES and those of solar hard X-ray flux obtained by YOHKOH/HXT.
1995; Escoubet et al., 1997; Nakagawa et ah, 2000; Ishisakaet al., 2001]. In multi-component photoelectrons, the low-energy component is dominant in the vicinity (~ several meters) of the probes and the spacecraft body [Nakagawa et al.. 2000], while the high-energy component plays important roles in modifying largescale sheath structure. Since the probe length is much longer than the Debye length (~ 10m) in the solar wind environment, EFD-P could be sensitive to the high energy part of photoelectrons. EVENT STUDY Time profiles of GEOTAIL electric field measurement, GOES soft X-ray flux and YOHKOH/HXT hard X-ray flux from 1500UT to 1530UT during the November 24, 2000, solar flare (X2.3/2B) are shown in Figure 1. GEOTAIL was in the solar wind at (X,Y,Z)GSK = (13.7, -26.1, 3.2)RR during the event. Panel (a) of Figure 1 shows time profiles of the sunward component of the double-probe electric field Ex (a solid curve), —(V x B) x electric field (a dotted curve) calculated from ion bulk flow [Mukai et al., 1994] and magnetic field [Kokubun et al., 1994] measurements, and spacecraft potential Vg/c ( a dashed curve) obtained by EFDSP [Tsuruda et al., 1994]. The sunward double-probe electric fields, Z?x, gradually increased from 1500UT, reached a maximum at 1508:08UT impulsively, and then gradually decreased after 1514UT. Time scales of the impulsive and gradual enhancement were several minutes and several tens of minutes, respectively. There was no corresponding change in the —(V x B) x profile. Thus we conclude that these variations of £ x are not due to changes in the solar wind electric field, but due to the change of the Z?x offset. Note that no remarkable variations were found in the spacecraft potential (a dashed curve) and duskward electric field, Ey (not shown). Searching the cause of this event, we have found simultaneous enhancement of solar X-ray flux. Figure 1 (b) and (c) show time profiles of solar soft X-ray flux (1-8 A and 0.5-4 A) monitored by GOES and those of solar hard X-ray flux (L-band: 14-23, Ml-band: 23-33, M2-band: 33-53, and L-band: 53-93 keV) monitored -302-
Fig. 2. Position of the GEOTAIL satellite with respect to 100 E^ enhancement events associated with the solar flares. Most of events are observed in solar wind, though some of large E* enhancement events are observed in the magnetosphere.
by YOHKOH/HXT [Kosugi et al., 1991], respectively. It is remarkable that the time profiles of hard Xray flux in Ml, M2, and H-bands during the flare resemble that of the impulsive 2?x enhancement. The maximum time of hard X-ray flux above 23 keV is 1508:02UT. The time difference between the maximum times of hard X-ray flux and Ex is as small as 6 seconds. This result suggests that there is close relationship between the impulsive enhancement of E* and the solar hard X-ray flux (>23 keV). STATISTICAL STUDY In this section we compare enhancement of iJx with YOHKOH hard X-ray observations statistically. Starting from the hard X-ray flare listing summarized by the YOHKOH/HXT team (http://www.solar.isas.ac.jp/ HXT/), we have found 100 events for the period from September, 1993 to December, 2000, in which JE?X enhancement is clearly detected in association with hard X-ray flux increase. The 100 events correspond to 5.9% of 1708 hard X-ray flares in the listing (X class: 34, M class: 460, C class: 1214). This percentage increases to 41% if we restrict our discussion to X class flares. Figure 2 shows the GEOTAIL location at the time of 100 events. In this figure it is seen that most of the events are observed in solar wind, though some of large Ex enhancement events are observed in the magnetosphere. The reason a larger number of events are observed in solar wind than in the magnetosphere may be interpreted as follows. In the magnetosphere, the fluctuation level of the convection electric field, - ( V x B) x , is higher so that the solar-flare-associated variations of E^, if any, are easily masked. On the other hand, the fluctuation level of — (V x B) x in the solar wind is relatively low. Figure 3 shows scatter plots for the increment of E* (A£?x) versus the YOHKOH/HXT observations. There is close relationship between AEy and Ml, M2, and H band peak fluxes (b, c, d), while the relationship between AEX and L band peak flux (a) is less clear. No remarkable variations are found in the spacecraft potential like the event study. DISCUSSIONS In this paper we have studied the enhancement of sunward electric field, £7X, detected by EFD-P onboard GEOTAIL. which takes place in association with hard (and soft) X-ray emission from the solar flare. The event study of the November 24, 2000 flare and the statistical study clearly reveal that the sunward electric field, Bx detected by the double-probe technique, has its origin in the solar flare. The connection needs be
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Fig. 3. Scatter plots of the increment of £ x ( A £ x ) versus the YOHKOH/HXT observations for the 100 events (see text). The A £ x during the solar flare is defined as the value after subtraction of the background. R is the correlation coefficient.
carefully examined, however. As the double-probe technique is sensitive to the photoelectron sheath surrounding the spacecraft, the £ x enhancement event in discussion is most probably associated with photoelectrons that are yielded by the solar flare emission. Hard X-rays reported in this study are caused by the high-energy electron precipitation into the atmosphere of the sun. and it is known that three kinds of electromagnetic waves (microwaves, EUV and hard X-rays) show a similar impulsive time profile in solar flares [Priest, 1981]. Since the microwaves do not affect the photoelectron yield, and we focus on EUV and hard X-rays emissions to explain the mechanism of E-x enhancement. It has been known that flare EUV emission consists of two components: one is a gradual component, resembling soft X-rays, and the other is an impulsive component, resembling hard X-rays [Kane et al., 1979]. We have made a preliminary analysis of the TRACE EUV/UV data and found that during the November 24, 2000, flare (Figure 1) the 1600 A photon flux at 1508 UT increased by a factor of ~2.5 from the pre-flare background level. It is noted that the ratio of the photon flux of the hard X-ray to that of the EUV is 10~10 or less even during solar flares [Golub and PasachofF, 1997]. The photoelectrons yielded by hard X-rays (several ten eV) would be at most 104 (hard X-ray energy / work function of material (several eV) ~ 104), but should be much less than this value in reality. The photoelectron yield for EUV photons
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(> lOeV) is about 0.1 [Grard, 1973], on the other hand. Therefore it is likely that the EUV component of the solar flare impulsive emission has a causal relation with the observed changes of the E^ offset. Hard X-ray flux during the solar flare is an important proxy in this study since it is difficult to obtain whole EUV line emissions from available satellite measurements. We mentioned in the introduction section that photoelectrons consist of multi components. The highenergy component of photoelectrons can produce the asymmetry of the high-energy photoelectron cloud beyond the Debye length scale, which would result in the local enhancement of the sunward electric field. Finally let us compare our observations with a similar work using a detector onboard the Pioneer Venus Orbiter (PVO) by Brace et al. [1988]. In the PVO case, it was reported that photoelectron current enhancement was measured in association with solar flares by the Langmuir probe. This current enhancement detection is possible because the Langmuir probe has a negative bias so that emitted photoelectrons can escape away. On the other hand, as GEOTAIL is charged positively, low-energy photoelectrons are easily trapped, and enhancement of low-energy photoelectrons (< several eV) do not necessarily result in the enhancement of the total photoelectron current and the change of the spacecraft potential. REFERENCES Brace, L. H., W. R. Hoegy, and R. F. Theis, Solar EUV measurements at Venus based on photoelectron emission from the Pioneer Venus Langmuir probe, J. Geophys. Res., 93, 7282-7296, 1988. Escoubet, C. P., A. Pedersen. R. Schmidt, and P. A. Lindqvist, Density in the magnetosphere inferred from ISEE 1 spacecraft potential, J. Geophys. Res., 102, 17595-17609, 1997. Grard R. J. L., Properties of satellite photoelectron sheath derived from photoemission laboratory measurements, J. Geophys. Res., 78, 2885-2906, 1973. Garrett, H. B, The charging of spacecraft surfaces, Rev. of Geophysics and Space Physics, 19, 577-616, 1981. Golub. L., and J. M. Pasachoff, The solar corona, Cambridge University Press, 162, 1997. Ishisaka, K., T. Okada, K. Tsuruda, H. Hayakawa, T. Mukai, and H. Matsumoto, Relationship between the Geotail spacecraft potential and the magnetospheric electron number density including the distant tail regions, J. Geophys. Res., 106, 6309-6319, 2001. Kane, S. R., K. J. Frost, and R. F. Donnelly, Relationship between hard X-ray and EUV sources in solar flares, Astrophys. J., 234, 669-682, 1979. Kokubun, S., T. Yamamoto, M. H. Acuria, K. Hayashi, K. Shiokawa, and H. Kawano, The GEOTAIL magnetic field experiment, J. Geomag. Geoelectr., 46, 7-21, 1994. Kosugi, T., K. Makishima, T. Murakami, T. Sakao, T. Dotani, M. Inda, K. Kai, S. Masuda, H. Nakajima, Y. Ogawa, M. Sawa, and K. Shibasaki, The hard X-ray telescope (HXT) for the solar-A mission, Solar Physics, 136, 17-36, 1991. Mukai, T., S. Machida, Y. Saito, M. Hirahara, T. Terasawa. N. Kaya, T. Obara, M. Ejiri, and A. Nishida, The low energy particle (LEP) experiment onboard the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 669-692, 1994. Nakagawa, T., T. Ishii, T. Tsuruda, K. Hayakawa, and T. Mukai, Net current density of photoelectrons emitted from the surface of the GEOTAIL spacecraft, Earth Planets Space, 52, 283-292, 2000. Pedersen, A., Solar wind and magnetosphere plasma diagnostics by spacecraft electrostatic potential measurements, Ann. Geophysicae, 13, 118-129, 1995. Priest, E. R., Solar flare magnetohydrodynamics, Gordon and Breach Science Publishers, 1981. Schmidt, R. and A. Pedersen, Long-term behavior of photo-electron mission from the electric field double probe sensors on GEOS-2, Planet. Space Sci., 35, 61-70, 1987. Takei, Y., T. Mukai, Y. Saito, H. Hayakawa, and K. Tsuruda, GEOTAIL study of comparison between the double-probe electric fields and the convection electric fields in the distant tail, in this issue. Tsuruda, K., H. Hayakawa, M. Nakanura, T. Okada, A. Matsuoka, F. S. Mozer, and R. Schmidt, Electric field measurements on the GEOTAIL satellite, J. Geomag. Geoelectr., 46, 693-711, 1994. E-mail address of Y. Takei: [email protected]
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QUEST FOR WAVES EXCITED BY INTERSTELLAR HELIUM PICKUP IONS M. Oka1 and T. Terasawa1 1
The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo, 113-0033, JAPAN ABSTRACT
'Torus-like' velocity distribution of the helium pickup ions of interstellar origin has directly been observed by the GEOTAIL spacecraft. These distributions possess a free energy to excite an electromagnetic wave owing to the anisotropic distribution. Although there have been several reports on the quest of waves induced by hydrogen pickup ions based on the observation beyond ~4AU, there has been no report on the waves of helium pickup ions. In this brief report, power spectrum analysis for helium pickup ion events obtained by GEOTAIL is carried out to search for such a wave excitement. We report, however, that the expected wave could not be found owing to the low density of the helium pickup ions at 1AU. The upper limit of the waves is determined to be 24.3% above the background level.
INTRODUCTION It is widely known that the Heliosphere is filled with the pickup ions (PUIs) of interstellar origin. The density of the PUIs is monitored throughout the year by both in situ and remote-sensing measurements in order to deduce the parameters of the local interstellar cloud. On the other hand, there is a consensus that the PUIs are the sources of the anomalous cosmic rays, and various theories have been proposed for their acceleration mechanisms. For the study of such various phenomena related to PUIs, either in observation or theory, it is important to know their velocity distribution function and their transport properties. The velocity distribution of the PUIs was first obtained by the earth orbiting satellite AMPTE/IRM utilizing the mass-spectrometer (Mobius et al.,1985). Since then, the PUI observation has been successively continued by advanced spacecraft such as ACE and ULYSSES. However, the biggest constraint on the observations was that the distribution functions were obtained only in one or two dimension, and there was always an assumption that the pitch-angle scattering rate is large and the distribution shape of the PUIs is of a complete sphere-shell. Recently, on the other hand, measurements onboard GEOTAIL have shown that the 'torus-like' structure of the velocity distribution exists, and such a pattern may be observed for about 20% of the observation time of an earth-orbiting satellite (Oka et al., 2002; hereafter Paper I). In this brief report, we go one step further by analyzing the magnetic field fluctuations accompanied by the events reported in Paper I. As the 'torus-like' distribution is far from isotropic, its relative streaming with the solar wind drives instabilities, and electromagnetic waves are expected to be excited (e.g. Wu and Davidson, 1972). Although the waves accompanying the cometary PUIs had been found and studied intensively (e.g. Glassmeier et al., 1989), there had been only a few decisive evidence for wave enhancements due to the interstellar hydrogen PUIs (Smith, 1993; Goldstein et al., 1993; Murphy et al., 1995; Intriligator et al., 1996). For helium part, there has been no report on the quest of wave excitement in spite of the prediction of a quasi-linear theory (Lee and Ip, 1987). In the next sections, we present the results of our spectrum analysis and report that any kind of wave excitement by the helium PUIs was not confirmed. -306-
Fig. 1. Annual variation of the He + PSD data utilized in the statistics. The DOY (=day of year) 330 corresponds to November 25 or 26.
Fig. 2. The B spectra for the 'spherical' event obtained on July 1, 1996. The magnetic field is decomposed into right (R) and left (L) hand polarized component and compressional (C) component.
DATASET We have utilized the PUI events reported in Paper I. The GEOTAIL data base of the time period between February 1994 and September 2001 was used to select the events. These events are determined by analyzing the measurement of the plasma (LEP/EAI, Mukai et al., 1994; CPI, Frank et al., 1994) and magnetic field (MGF, Kokubun et al., 1994) experiments. In Paper I, the events were categorized by the PUI velocity distributions into four categories: 'torus-like', 'spherical', 'unidentified', and 'unclassified'. 13 and 10 events were categorized as 'spherical' and 'torus-like', respectively. In this report, we deal with the 'spherical' and 'torus-like' events. Figure 1 shows the phase space density (PSD) of each 'torus-like' and 'spherical' events aligned by the date of year of the observation. The annual variation of the He+ PSD shows an enhancement in the early December. Since the earth is known to cross the gravitational focusing cone formed by the neutral particles from the interstellar cloud every winter, the enhancement in the figure supports the appropriateness of our methodology of detecting the PUIs of the interstellar origin with an electrostatic analyzer. THE MAGNETIC FIELD SPECTRUM Lee and Ip (1987) first derived the quasi-linear spectrum of waves that result from the pitch-angle scattering of ring-like distributed interstellar PUIs to a sphere-shell distribution. Besides, Johnstone et al. (1991) utilize an alternative approach which assumes that the bispherical distribution constructed by parallel and anti-parallel propagating Alfven waves represents a good approximation to the asymptotic PUI distribution. In this report, we follow the analysis of Huddleston and Johnstone (1992) who presented an application of the theory to the observation of the waves excited by the cometary PUIs. We have utilized the magnetic field (MGF) measurement of 3 sec-sampling and obtained power spectra for each 'spherical' and 'torus-like' distribution events. The frequencies, fsc, measured in the spacecraft frame, are transformed into the wavenumber, k, in the solar wind frame, according to Huddleston et al. (1991) and Huddleston and Johnstone (1992). Note that the spectra of waves propagating upstream and downstream directions are not separated. The peak of the power spectrum contributed from the interstellar PUIs may be predicted by the wavenumber ki, which corresponds to the approximate minimum wavenumber generated during the pickup process, i.e., from the resonance condition ki ~ il/u^. The peak, if exists, should appear at around or higher than the wavenumber k^. For the identification of the waves generated by the PUIs, it is important to determine the 'background' power spectrum of the solar wind. However, it cannot be determined consistently for the cases of interstellar PUI observation because the solar wind magnetic field may be affected by the interstellar PUIs at any time period. Here, we assume that its spectrum varies as a power law with an index 7 = 2, as is usually the case for the nominal solar wind condition. Figure 2 shows an example of the wave spectrum obtained during a 'spherical' event on July 1, 1996. This event has already been reported in the context of the identification of He+ PUIs being shock drift accelerated -307-
Table 1. Observed Parameters
Parameter Vsw, km/s Bo, nT dc, 0-90° Nsw, /cc VA, km/s
Jul. 1, 1996 309.3 2.3 42.3 8.1 17.8
Dec. 4, 2000 363.0 7.4 64.0 3.8 83.4
Fig. 3. The B spectrum for the 'torus-like' event obtained on December 4, 2000.
(See Figure 3 of Oka et al., 2002b for the velocity distribution). We first divided the magnetic field data into three components: right (R) and left (L) hand polarized components and compressional (C) component. Then, for each component, we took 256-sample-point spectra (corresponding 13.1 minutes time period) and performed a half sliding average to produce the spectra that span 3.5-hour time period, resulting with the coefficient of variance to be 18.0%. The arrow in the figure indicates the value of ki. We do not find any statistically significant feature in the spectra. They only feature a power-law with a spectral index 7 = 2, identical to the spectra of the nominal solar wind condition. A little bump near ki is within a statistical error, and it can be concluded that the signal for interstellar PUI was not detected. Figure 3 shows an example of the wave spectrum obtained during a 'torus-like' event on December 4, 2000. This event was first reported and thoroughly described in Paper I. We took 128-sample-point spectra (corresponding 6.5 minutes time period) and performed a half sliding average to produce the spectra that span 1-hour time period, resulting with the coefficient of variance to be 24.3%. Again, the spectra are divided into R, L, and C component. The spectra do not feature a single power law. The power spectral density does not decrease constantly in the wavenumber range of 4 x 10~7 to 2 x 10~6 [/m]. In the high wavenumber range above 2 x 10~6 [/m], there is a noisy power enhancement. The analysis of the high resolution (16Hz) MGF data suggests that the noise may be a part of ULF waves typically observed in the bow shock upstream region. However, we do not discuss on this as the peak of PUI wave, if exists, should not show itself in this high frequency range. We rather call attention to the change in the power spectral density at around 9 x 10 [/m]. The bump has a marginal statistical significance, and we decided to further study the spectrum. In order to determine whether the change in the spectrum is due to the effect of PUIs of interstellar origin, we compared the observation with the theoretical spectrum according to the formulation applied in Huddleston and Johnstone (1992). The parameters for the formulation are taken from the observation (Table 1). For the density of the helium PUIs, we used Njje+ — 2.2 x 10~4 [/cm3] as was estimated in Paper I for this event. For the background spectrum, the index of power law is set to be 7 = 2 as noted above. Then, the theoretical calculation indicated that the enhancement above the background level should be only 1.9% which is well below the observed statistical error 24.3%. Therefore, it is unlikely that the peak in the spectrum was the signal of the wave excited by the PUIs of interstellar origin. CONCLUDING REMARKS We have shown that the signature of the waves excited by helium PUIs was not detected in both examples of 'spherical' and 'torus-like' events. The waves were not found in all the other events as well. This result indicates that the pitch-angle scattering occurs owing predominantly to the magnetic field turbulence intrinsic to the background solar wind. Note, however, that the statistical errors for all the events we used were as large as 24.3%. Therefore, we still have not discarded the possibility of detecting the waves of the helium PUIs if we could suppress the statistical error. In order to investigate in which condition we could observe the wave excitement, we carried out a parameter survey of the theoretical power spectrum given by Johnstone et al. (1991). We assume that the magnetic field Bo — 5.0nT, the density of helium -308-
PUI NHe+ = 0.0002/cc, the solar wind speed Vsw = 400km/s, and the index and the normalization of the background spectrum are given by 7 — 2.0 and < <5B2 >= O.OlBo2) respectively. Under this condition, the cone angle 0c must be larger than ~85° and the solar wind density Nsul must be below ~0.4/cc, for the signal to have the significant level of 62% (~3 sigmas in our analysis) above the background power spectrum. Therefore, the PUI generated spectrum may be obtained in the low density solar wind events. It has been recognized that the density rarefaction (
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AUTHOR INDEX Anderson, R.R.
205
Balogh, A. Bamert, K. Baumjohann, W. Bogdanova, J. Bougeret, J.-L.
177 285 177 177 205
Christon, S.P.
100
Escoubet, C.P.
9
Goto,Y. Green, J.L. Hashimoto, C. Hashimoto, K. Hayakawa,H. Hori,T. Hoshino, M. Huang, X. Ieda,A. Igarashi, K. Imada, S. Ishisaka,K. Kaiser, MX. Kallenbach, R. Kamide, Y. Kasaba,Y. Kasahara, Y. Kawano, H. Keika, K. Kistler, L. Klecker,B. Kojima, H. Kokubun, S. Kondoh, K. Kurth, W.
v 172 113 161
Machida, S. 186 Malova, H.V. 100 Matsumoto, H. 75, 205, 220, 224, 247, 293 Matsuoka, A. 64,301 McEntire, R.W. 48, 198, 297 Mizuta, T. 261 Mukai, T. 1, 19, 34, 43, 75, 79, 172, 177, 186,247,281,285,293,301 Murata, K.T. 220
48, 198,297
Delcourt, D.C.
Fehringer, M. Fujimoto, M.
Lanzerotti, LJ. Le, G. Lui, A.T.Y. Lyons, L.R.
9 19, 38, 123, 130, 143, 177
Nagai, T. Nagano, I. Nakai,H. Nakamura, M. Nakamura, R. Nakamura, T.K.M. Nakata, K. Nishino, M.N. Noda, H. Nose, M. Nowada, M.
228 235 130 205,220 75,79,301 190 v, 28,34, 108,261,289 235 186 285 34 75
Ohtani, S.-I. Oka, M. Okada, T. Omura, Y.
205 285 194 205,247 228 172 297 177 177,285 205,293 19,190 135 220
186 306 75 v, 247
Parks, G.K. Paschmann, G. Popov, V.Yu. Puhl-Quinn, P.
172 177 100 177
Quinn, J.
177
Reinisch, B.W. Reme, H. Rostoker, G. Russell, C.T. Saito, Y. Sakurai, T.
-311-
87, 177, 186, 190 224 194 301 177, 190 38 281 28 177 48, 198,297 43,71
235 177 172,205 172 79,172,186,281,285 43,71
Sauvaud, J. A. Sharma,A.S. Shibata,K. Shimada,N. Shimizu, T. Shin, K. Shinkai,Y. Shinohara,I. Shue, J.-H. Song, P. Steinberg, J.-L.
177 100 301 281,289 135,139 293 71 123,281 186 235 205
Takano, H. Takasaki,H. Takei, Y. TanDokoro, R. Terasawa, T.
224 301 79, 301 130, 143 28, 267, 281, 285, 301, 306
Tonegawa, Y. Torbert, R. Torkar, K. Troshichev, O. Tsuruda, K. Uchiyama, H. Ugai, M.
-312-
71 177 177 54 79 228 135,139
Vaith, H.
177
Yagitani, S. Yokoyama, T. Yoon, P.H.
224 147 251
Zelenyi, L.M.
100