MULTI-SPACECRAFT OBSERVATIONS OF PARTICLE EVENTS AND INTERPLANETARY SHOCKS DURING NOVEMBER/DECEMBER 1982

Transcription

1 MULTI-SPACECRAFT OBSERVATIONS OF PARTICLE EVENTS AND INTERPLANETARY SHOCKS DURING NOVEMBER/DECEMBER 1982 M.-B. KALLENRODE, G. WIBBERENZ, H. KUNOW, R. MÜLLER-MELLIN Institut für Reine und Angewandte Kernphysik, University of Kiel, Germany and V. STOLPOVSKII and N. KONTOR Nuclear Physics Institute, Moscow State University, Moscow, Russia (Received 25 October, 1992; in revised form 3 February, 1993) Abstraet. We present a sample of solar energetic particle events observed between November 18 and December 31, 1982 by the HELIOS 1, the VENERA 13, and IMP 8 spacecraft. During the entire time period all three spacecraft were magnetically connected to the western hemisphere of the Sun with varying radial and angular distances from the flares. Eleven proton events, all of them associated with interplanetary shocks, were observed by the three spacecraft. These events are visible in the low-energy (about 4 MeV) as weil as the high-energy (30 MeV) protons. In the largest events protons were observed up to energies of about 100 MeV. The shocks were rather fast and in some cases extended to more than 90 east of the flare site. Assuming a symmetrical configuration, this would correspond to a total angular extent of some interplanetary shocks of about 180. In addition, due to the use of three spacecraft at different locations we find some indication for the shape of the shock front: the shocks are fastest close to the flare normal and are slower at the eastern flank. For particle acceleration we find that close to the flare normal the shock is most effective in accelerating energetic particles. This efficiency decreases for observers connected to the eastern flank of the shock. In this case, the efficiency of shock acceleration for high-energy protons decreases fastet than for low-energy protons. Observation of the time-intensity profiles combined with variations of the anisotropy and of the steepness of the proton spectrum allows one in general to define two components of an event which we term 'solar' and 'interplanetary'. We attempt to describe the results in terms of a radially variable efficiency of shock acceleration. Under the assumption that the shock is responsible not only for the interplanetary, but also for the solar component, we find evidence for a very efficient particle acceleration while the shock is still close to the Sun, e.g., in the corona. In addition, we discuss this series of strong flares and interplanetary shocks as a possible source for the formation of a superevent. 1. Introduction Enhancements of energetic particles, particularly in the energy range below 20 MeV, associated with the passage of an interplanetary shock, are a well-known phenomenon, first reported by Bryant et al. (1962). These enhancements, called Energetic Storm Particle (ESP) events, were originally thought to be related mainly to the time of shock passage, thus the shock was believed to have only a local effect on the energetic particte population. Recent research has changed this picture (e.g., Cane, Reames, and von Rosenvinge, 1988, hereafter CRvR; Reames, 1990; Cane, Reames, and von Rosenvinge, 1991). In particular, CRvR showed in a study of 235 proton events observed from Earth-bound IMP spacecraft during a time period of 20 years that, at least for protons with energies less than about 20 MeV, (1) the most intense events come Solar Physics 147: , (~) 1993 KIuwer Academic Publishers. Printed in Belgium.

2 378 M.-B. KALLENRODE ET AL. from flares near central meridian, as do the strongest shocks, rather than from the magnetic connection longitude of Earth at about W60, (2) the greatest intensities are observed when the observer is magnetically connected to the strongest region of the shock and not when the weaker portion of the shock passes the observer, and (3) the acceleration of energetic particles at an interplanetary shock is not limited to energies of some MeV but can extend up to energies well above 100 MeV. Note that CRvR refer to the intensity of the overall intensity time profile, these are not necessarily the particles accelerated at or close to the Sun, but can also be particles accelerated in interplanetary space. The influence of the interplanetary shock is not only limited to the time of shock passage but the larger proton events often are dominated by particles accelerated at the interplanetary shock, either due to the acceleration of material from the solar wind (cf., Reames, 1990) or due to re-acceleration of particles (cf., Tan et al., 1990) accelerated in the preceding solar rare, e.g., in the flare-heated material or by a coronal shock. The analysis of solar particle abundances shows that the intensity profiles and abundances of such large proton events as a function of longitude best can be interpreted in terms oftwo particle sources: (1) particles accelerated directly in the rare with abundances that resemble the composition expected from the flare-heated material (Cane, Reames, and von Rosenvinge, 1991), as, e.g., Fe/O enhancements by a factor of ten compared to the coronal abundances, and (2) particles accelerated at coronal and interplanetary shocks with compositions comparable to the corona and the solar wind. While the flare-accelerated component (1) can be present in well-connected events only (cf., Reames, Cane, and von Rosenvinge, 1990), the shock-accelerated component (2) is present 'everywhere' up to angular distances of more than 100 away from the rare site and is dominant especially in events in which the observer is located close to the rare normal, where the shock hits directly, and in events with poor magnetic connection to the rare site. Thus the interplanetary shock is believed to play a crucial role in supplying particles to regions not magnetically connected to the flare site. As a consequence of this interpretation, an azimuthal transport process in the corona to account for the large angular separations between rare and observer (e.g., Reinhard and Wibberenz, 1974; van Hollebeke, Ma Sung, and McDonald, 1975) is no longer necessary (cf., discussion in CRvR). In addition to the composition, also the charge states of heavy elements allow some clues on the acceleration mechanisms. The flare-accelerated particles in the impulsive flares not only show a higher content of heavy elements but also higher charge states of these elements. In impulsive flares, Fe, for example, has a charge state of 20.5 (Lun et al., 1987), indicative of temperatures above 10 MK and therefore clearly originating from flare-heated material. In contrast, Fe in the large gradual events has a charge state of 14.1 (Luhn et al., 1985, 1987), similar to that in the solar wind. This is indicative of lower temperatures at the acceleration site (around 2 MK) and therefore gives evidence for the acceleration of these particles out of the corona or the solar wind by a shock wave, but not out of the flare-heated

3 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 379 material. The above picture of a shock being responsible for the acceleration and the 'azimuthal spread' of particles is attractive insofar as it gives an elegant and consistent explanation for many observations, in particular the differences between impulsive and gradual flares and the large azimuthal ranges over which particles can be observed in gradual flares. However, there are still some open questions (cf., Kallenrode, 1993), in particular: (1) does the interplanetary shock accelerate particles out of the solar wind or is a pre-accelerated particle component required, (2) how do the shock properties develop during its expansion and propagation, and (3) how does its efficiency for particle acceleration develop (a) due to the evolution of the shock and (b) due to changes in shock geometry: at small distances from the Sun the Archimedian spiral is nearly radial, therefore a shock tends to be quasi-parallel, while at distances beyond 1 AU the curvature of the magnetic field line increases, and the shock becomes more and more quasi-perpendicular. Our instrumentation does not allow us to distinguish between flare-accelerated and shock-accelerated particles (poor isotopic resolution, no charge states), however, there are some other indications for two components (not necessarily the same two components as described by Reames, Cane, and von Rosenvinge, 1991). The first evidence comes from the time scales and profiles (cf., Kallenrode and Wibberenz, 1992) observed in events accompanied by an interplanetary shock. In many events the intensity rises fast to a first maximum, well before the arrival of the shock. Kallenrode, Cliver, and Wibberenz (1992) found for 52 gradual events observed at different distances a mean value for the time to maximum of 6.2 hr for 10 MeV protons and 2.8 hr for 0.5 MeV electrons. Note that these maxima are not necessarily the absolute maxima of the entire time profile. Often the intensity stays constant or later rises up to a maximum around the time of shock passage. These short time scales, together with the interplanetary propagation, indicate particle acceleration on or at least close to the Sun, but not necessarily in the flare process. Therefore we will term this particle component the solar component. The second evidence for two particle components comes from the time development of the proton spectrum in the range 4-50 MeV (cf., Meyer, Wibberenz, and Kallenrode, 1992): the proton spectrum is rather hard early in the event (when the solar component is dominant) and shows a characteristic weakening as the shock approaches with a weakest spectrum around or shortly after the time of shock passage. At the time of the weakening of the spectmm, the intensity, especially of the low energy protons, is increased, either as a 'hump' on the decaying flank of the solar component (as in the classical picture by Bryant et al., 1962, cf. Figure 1 in CRvR), as a continuous increase of intensity or as a plateau in intensity (cf., e.g., Reames, 1990, especially Figures 1 and 2). This component with the weaker spectrum leading to this second maximum will be termed the interplanetary component and is related to continuous particle acceleration at the approaching interplanetary shock. One

4 380 M.-B. KALLENRODE ET AL. part of the interplanetary component may consist of particles stored in the turbulence around the shock front (cf., Reames, 1990; theoretical considerations of this problem in Lee, 1983). Thus we have two components, the designation of these components should give an indication of the acceleration site but not of the acceleration mechanism. The interplanetary component is accelerated continuously by the interplanetary shock as it travels outward and as long as it is capable of accelerating particles. The solar component can be accelerated in the flare process itself (in well-connected events sometimes evidenced by the high abundances ofheavy elements, cf., Reames, Cane, and von Rosenvinge, 1990) or by a shock wave travelling through the corona (cf., Section 4), in general connected with a coronal mass ejection (CME). Thus, in principle, one shock might be responsible for the acceleration of both components. Because these two components are not separated sharply and the interplanetary propagation tends to smear out and mix these components even stronger at the observer's site, we cannot separate clearly between them (as evidenced by the smooth changes in Fe/O in Figure 4 in Reames, Cane, and von Rosenvinge, 1991, or in the time development of the proton spectrum in Meyer, Wibberenz, and Kallenrode, 1992). We rather can say that for early times the solar component is dominant, while at later times the interplanetary component becomes the dominant one. Note that there are also many events which consist of an interplanetary component only (Cane, von Rosenvinge, and McGuire, 1990; Kallenrode and Wibberenz, 1992). In this study we will supply support for the above picture of the shock being responsible for transport and acceleration by means of a multi-spacecraft study. This kind of study allows one to analyse individual events observed at different radial and angular distances from the flare site and therefore provides information about the spatial extent of the shock and changes of shock speed with distance from the flare in a rather direct way. Five large proton events associated with interplanetary shocks (in the terminology of CRvR these events would be western flares) will be analyzed in detail with special consideration for the relative contributions of the solar and the interplanetary components. The role of interplanetary shocks for the particle acceleration is not only limited to isolated particle events. In addition, systems of shocks seem to play an important role in the formation of superevents (Müller-Mellin, Röhrs, and Wibberenz, 1986, hereafter MRW; Dröge, Müller-Mellin, and Cliver, 1992, hereafter DMC). These superevents differ distinctly from other intensity increases, e.g., solar energetic particle events, corotating events or quiet-time increases of electrons in the following characteristics: the intensity remains above background level for a period in the order of one month, the time scales of onset and decay are much longer than in flare-associated particle events, radial and azimuthal gradients are small, there is no pronounced anisotropy, the proton to helium ratio is constant, and ions and electrons are enhanced simultaneously. Obviously, because of their long time scales, superevents are not an isolated phenomenon, but they are superposed by flare-associated particle events. MRW identified four such superevents during nine

5 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 381 years of the HELIOS mission. By using 54-day running mean intensities instead of data with high time resolution, DMC identified additional 12 time periods with long-lasting high intensities, observed simultaneously by IMP, HELIOS, and PI- ONEER 10/11. Because the time period considered in this paper is defined as a superevent by DMC, we will discuss not only the five large events that seem to dominate the time period but also consider the entire activity in form of flares and interplanetary shocks. The paper is structured as follows: after this Introduction we will describe the instruments and the data selection in Section 2. In Section 3 we will describe the entire time period and will discuss five large events in detail. The relation between these observations and recent ideas about the importance of interplanetary shocks for the acceleration and propagation of energetic particles, as weil as for the formation of superevents, will be discussed in Section Instruments and Data/Sources / The particle data used in this paper were obtaine~by the University of Kiel particle instrument (E6) on board HELIOS 1, the Moscow State University Nuclear Physics Institute particle instruments KV-77 on board VENERA 13/14, and the Johns- Hopkins University Applied Physics Laboratory particle instrument on board the IMP 8 satellite. The University of Kiel particle instrument consists of a stack of five solid state detectors of increasing thickness and is shielded by an anticoincidence scintillator (Kunow et al., 1977). The instrument measures electrons between 0.3 and 3 MeV in three, protons between 4 and 51 MeV in four, and helium between 2 and 48 MeV nuc1-1 in five energy channels. Due to the revolution of the spacecraft (1 rotation per second) around an axis perpendicular to the ecliptic plane, for most of the energy channels, particles were counted into eight sectors, providing information about the angular distribution of the particles in the plane of the ecliptic. Crude information about higher energies can be obtained from the integral counting channels for protons (above 51 MeV) and helium (above 48 MeV nucl-1). The Moscow State University KV-77 instrument consists of three detector blocks viewing in directions 45 (MEDS), 225 (MEDT), and 90 (E1D2) of the Sun-spacecraft line (of., Logachev et al., 1988; Belyakov e~ al., 1984). The MEDS and MEDT blocks measure electrons above 30 kev by means of the difference method with two open-window, gas-discharge counters with active shielding of a plastic scintillator, and protons and helium above 1 MeV nuc1-1 with a single semi-conductor. The E1D2 block measures electrons between 70 kev and 1.5 MeV and protons between 25 and 230 MeV with a combination of scintillation and semiconductor detectors. Because of the rotation of the spacecraft around the Sun-spacecraft axis, the identical MEDS and MEDT instruments look during part of their measuring time along the nominal Archimedian field line. Thus a crude estimate of the angular

6 382 M.-B. KALLENRODE ET AL. distribution of the particles is possible by defining the anisotropy A as A= IMEDS -- IrMEDT /MEDS + ]'MEDT Note that this anisotropy has not exactly the same meaning as on HELIOS because it is not only mathematically derived in a different way, but in addition the revolution of the sensor is not in the plane of the ecliptic, as in case of the HELIOS sensor, but perpendicular to it. Within the intention of this study, however, these differences are of minor importance. For both, HELIOS 1 and VENERA, the data were directly available, the IMP 8 data were provided by the NSSDCA. The IMP 8 particle data were obtained by the Johns-Hopkins University Applied Physics Laboratory (JHU/APL) particle instrument. (Instrument description and details on data analysis techniques may be found in Sarris et al. (1976).) The telescope consists of three colinear sensors (two surface-barrier, totally depleted detectors followed by a lithium-drifted detector). All four spacecraft are in different orbits in the plane of the ecliptic: IMP 8 is an Earth-bound satellite in a near-ecliptic, roughly circular orbit at about 40 earth radii, HELIOS 1 is in a highly eccentric orbit around the Sun with radial distances ranging between 0.3 and 0.98 AU, and the VENERA spacecraft follow a trajectory basically oscillating between Earth and Venus with a separation of about 106 km and < 1 between the two spacecraft. Because of the small separation between VENERA 13 and 14, both spacecraft are considered as one observation point only, in all the following figures (with the exception of the 1982 December 6 event) we will show VENERA 13 data only OVERVIEW 3. Observations Figure 1 gives an overview of the proton intensities observed by the three spacecraft during the time period 1982 November 18 to December 31, as weil as of the flares that occurred during the time interval under study (for this method compare, e.g., Roelof et al., 1992, and references therein). The upper panel shows 4-13 and MeV protons measured by HELIOS, the second panel the intensities of and MeV protons obtained by VENERA 13, and the third panel gives the intensities of 1-2 and MeV protons measured by the JHU/APL instrument on IMP 8. The energy intervals are chosen as representative for lowenergy (--~4 MeV) and high-energy protons (--~30 MeV) to allow a check on the energy dependènce of shock acceleration. The lowest panel shows the longitudes of all flares with soff X-ray class between M1 and M4 (small open circles), M5 and M9 (large open circles), and >_X1 (filled cirlces) as weil as the parent flares for the eleven largest particle events (stars) observed during this time period. The ordering of the flares corresponds to the rotation of the active regions with respect to the central meridian.

7 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS g.... i.... t,,, 10 7 t! ~~'i, 106 "p I*('~ '" \ t,,.,o'- "" J 1 \ lo 3. "~. 10 ~~,~'~t'" 1,!l,..;il, v ~~I utv Protonl 4 -~ 2s-eouev ' ~ #~ \ ~o ~, \ e ~ ~3 i u;ù,r,~,,;w~.-t.' t~n~" "11-- ' I 1 -,8 M,,v pr ~Of~l 4 i 103- ~.: 10 2 " ~,.,, : -,, ',~~, «~--/..-'. -. :.%. 10 ".""."',...'.. "..,, ~ %%. 1. I "-.'". " "-~-..-.., ~". '%x 10, ~'''~.'~,'~ «':.," ~,~..- "- ""-Z. "- N.. 10: " "... "" " " 1~:... : '; ':' " I 1 3 II 7 I 10 -": AA AA AA A A A A AA 120 g,. 80 1~ ~ o 40 dr o ~ o i... q...,... i... I... ;d,os ' i"?? 9..? o 9.,o_,o,~ ù. l,,,~...,..j L,.,... "'%" o 0 o VENERA 0 0 IQ 11 tl ù.~,, I,,,, LX,, IUP I~ o o ù ~.,~. %'%...«- %E «',' w \ %. o. 0 o o~ o o ~~ 0 " 0 -L40 o. 0(:90 ~p ~o -80 o d9 0 ù,' ~,.i ' I.... I ' ' t.... I ' I 32, $ TIME Fig. 1. Proton intensities for low and high energies for the time pefiod 1982 November 18 to December 31 observed by HELIOS (upper panel), VENERA 13 (second panel), and IMP 8 (third panel) together with an overview of the solar activity (fourth panel, cf. text). The insert shows the spatial configuration of the IMP 8, HELIOS, and VENERA spacecraft during the time period 1982 November 18 to December 31. The dashed lines give Archimedian spirals calculated with a nominal solar wind speed of 400 km s-l The insert in the bottom panel in Figure 1 shows the trajectories of the spacecraft during this time. In addition to the trajectories, we show Archimedian spirals (dashed lines) calculated for a nominal solar wind velocity of 400 km s -1, illustrating the approximate range of magnetic footpoints for the time period under

8 384 M.-B. KALLENRODE ET AL. study. As a general rule, IMP 8 is located above central meridian and connected magnetically to about W60, the angle VENERA-Sun-Earth is between 31 and 49, which, considering the decreasing radial distance, corresponds to connection longitudes between W83 and W92. The much faster moving HELIOS spacecraft moves away from the other s/c during the time interval under consideration, with a spacecraft-sun-earth angle increasing from W60 to W127, corresponding to a connection longitude between W102 and W146. The entire time period is characterized by a relatively large number of flares as well as of particle events in interplanetary space. In Figure 1 numbered circles mark the 11 energetic particle events which will be considered in more detail in this study. The parent flares of these events all initiated interplanetary shocks. Triangles in Figure 1 mark the passage of interplanetary shocks at the observer; the numbers at the triangles refer to the numbers of the shock-initiating flares. Therefore, the same numbers in different panels indicate the same shock. Note that this record of interplanetary shocks is not complete, only for the HELIOS-spacecraft magnetic field and solar wind were data available for most of the time period; the shocks were identified by Schwenn (private communication). The VENERA magnetic field data were incomplete, but were available for most events discussed in detail data. The time of shock passage for IMP 8 was inferred from the sudden commencements published in the Solar Geophysical Data (SGD). The particle profiles during this time interval are characterized by five large (No. 2, 3, 6, 9, and 10) and six smaller proton events (No. 1, 4, 5, 7, 8, and 11) observed on all three spacecraft despite different radial and angular distances (cf. insert in the bottom panel in Figure 1). At least in the five large events, intensity increases are observed not only in the low-energy protons but also in about 30 MeV protons. The shocks observed at one of the spacecraft in general are also observed by the other spacecraft, either directly in plasma and magnetic field data of, if these data were not available, the intensity time profile gives evidence for the passage of a shock (cf. discussion of the individual events). The intensity, especially in the lower energy channels, stays at a relatively high level above background during most of the time period. From the lowest panel in Figure 1 we find that the ten large proton events associated with interplanetary shocks originated in the western hemisphere. This preference for flares in the western hemisphere is obviously due to our selection of large particle events from a time period during which the spacecraft were connected to the western hemisphere within about 40 of the western limb. The 45-day time period under study is characterized by a large number of solar flares and also by a fair number of interplanetary shocks. Many of the eastern events may contribute to the intensity profile observed on IMP 8 (cf., Cane, McGuire, and von Rosenvinge, 1986; CRvR), however, it is difficult to resolve the individual events in the particle data. The additional events on the western hemisphere can be partly identified in the HELIOS and VENERA intensities, however, these events do not stick out but lead to only small increases and sometimes mix with particles from preceding or

9 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 385 following events. Therefore statistical errors are too large for qualitative analysis. Note that these smaller events eventually may contribute to the relatively high background. In better time resolution for many of these flares smaller particle increases can be identified on at least one of the spacecraft. Table I summarizes the flare and spacecraft parameters for the eleven large events marked in Figure 1. Column (1) gives the date of the events, columns (2)- (4) the flare location, the flare classification in Ho~ and GOES soft X-rays, and the time of the Ho~ maximum (all from Solar Geophysical Data), column (5) gives the spacecraft (H1 =HELIOS, V=VENERA, IMP=IMP8), and columns (6)-(9) give the spatial parameters: the radial distance (6) in AU, the angle spacecraft-sun- Earth (7), the angle between the Sun-spacecraft line and a normal through the flare site (8), the coronal footpoint (9) calculated from the measured solar wind speed (11, ** corresponds to a nominal solar wind speed of 500 km s -1) according to the method of backmapping (Nolte and Roelof, 1973), and the coronal distance (10) as the angular distance between the flare site and the footpoint of the observers magnetic field line. A distinction between the two angles (8 and 10) is essential for an understanding of the different contributions of the solar and the interplanetary components (cf., discussion and figure in CRvR or Meyer e~ al. (1992)). Note that all flares in Table I are gradual. All flares in Table I were accompanied by interplanetary shocks; most of these shocks were clearly identified on two or three of the spacecraft (cf. Table II). Table II summarizes the shocks indicated in Figure 1. As mentioned above, VENERA data are incomplete, the timing of shocks on IMP 8 is inferred from the sudden commencements in SGD. Columns (1) and (2) give the date and the time of shock arrival at the observer. The classification is indicated in column (3). Columns (4) and (5) give the radial distance and the angular distance between the observer and the flare normal. Column (6) gives the locally measured shock speed (available for HELIOS only), column (7) gives the mean shock speed derived from the time delay between the flare and the shock arrival at the observer. Column (8) indicates the parent flare (numbers as in Table I) SELECTED EVENTS Five of the above events have been subjected to a closer inspection. They are presented in Figures 2-6, all figures are structured in the same way: in the upper right of the figure the spatial configuration of the three spacecraft is shown such that the flare is located on the central meridian. The Earth is at the same position as IMP 8, therefore one can easily rotate back to heliographic coordinates. Arrows from the Sun in the direction of the spacecraft in the spatial configuration plot give an indication for the average shock speed (column 7 in Table II) in that direction as inferred from the transit times. Around this spatial configuration we show the spacecraft data, organized in the following way: for IMP 8 (upper left) we show intensities of 1-2 and MeV protons. For VENERA (lower left) we show the intensities of , 25-60, , and MeV protons in the upper panel.

10 386 M.-B. KALLENRODE ET AL. ù~.~~ r.,o r~ r.~ r.,,o r~ r..o r.~ r.~ t.., (D <:~ oo ~~o~~o~~o~~o~~o~~o.~~o~~o O0 c~ ~ ~ t"-i 'q t".l < ~ O O O~ t"-- t",l ~ :.= ~ ~ oo o~ O~ ù_., d Z

11 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS ~+~+~ 0 ~ +1111~+11 c 1 r~ 0 er~ 0 ~ Cq 0 r,~ z o',, cq,., 0o z ~ z o O "O O d 2:

12 388 M.-B. KALLENRODE ET AL. TABLE II Shocks during November/December 1982 Doy Time s/c Rad. 4" n2/nl ~ Flare No :48 V E :17 IMP 1 E :48 V E :54 H W :22 IMP 1 E :53 H E ~01:00(?) V E :11 IMP :00 V W :29 IMP I :00 H E :48 V E :21 IMP 1 E :53 V :06 IMP 1 E :54 IMP :00 H W :20 IMP 1 El :08 H W :24 H El :58 V E :15 IMP :50 H W e~*: location of spacecraft with respect to flare. The middle panel shows the proton spectral index,, in the range 2 to 230 MeV, assuming a power law I(E) = IoE-'L Especially for such a broad energy range, a power law is only a very crude approximation; however, the spectral index, % is still a sufficient means to compare the time development of the intensities in different energy channels (cf., Meyer et al., 1992). The anisotropy of the MeV protons is given in the lower panel, calculated according to the method described above. Modulations in the anisotropy, sometimes also visible as modulations of the intensity, result from the rotation of the VENERA spacecraft: in the case of an anisotropic distribution, the anisotropy is highest when the instruments view along the magnetic field line and decreases when the instruments view perpendicular to it. A positive anisotropy indicates particles streaming from the Sun, a negative anisotropy indicates particles streaming towards the Sun. For HELIOS (lower right) we show the intensities of 4-13, 13-27, 27-37, and 37-5l MeV protons in

13 MULT1-SPACECRAFT OBSERVATIONS OF SHOCKS 389 the upper panel, the time evolution of the proton spectrum in this energy range in the middle panel, and the anisotropy of the 4-13 MeV protons in the lower panel. Note that the anisotropy in HELIOS is determined with respect to the local magnetic field direction, thus from the sign of the anisotropy alone we cannot infer whether particles are coming from the Sun or are streaming sunward. If no magnetic field data were available, the anisotropy was determined with respect to the direction of maximum particle flux (e.g., for the events in Figures 2 and 4). In these cases, the anisotropy is always positive. To clarify this situation we inferred the direction of preferential streaming from the angular distributions. A '+' sign on the anisotropy indicates particles streaming from the Sun, a '-' indicates particles streaming towards the Sun. Arrows 'S' mark shocks identified from the local magnetic field (on HELIOS and VENERA) or from the sudden commencements (on IMP 8), arrows. A question mark indicates probable shocks in which either the above sources were not reliable or not available The 1982 November 22 Event In discussing the individual events we will order them by the spatial configuration of the three spacecraft. Let us start with the 1982 November 22 event (doy 325, Figure 2) in which all three spacecraft are located nearly symmetrically around the Bare normal. The parent Bare was a gradual 1N/M7 Bare at S11 W36 with an Ha maximum at 18:17 UT. The angle between the Sun-spacecraft line and the normal is 25 for HELIOS, -2 for VENERA, and -36 for IMP 8. Here a negative angle indicates that the spacecraft is east of the Bare normal. While this configuration is symmetrical with respect to the Bare normal, the angular distance between the Bare and the footpoints of the observer's magnetic field lineare asymmetrical: all three spacecraft are connected to positions west of the Bare, with angular distances of 25 (IMP 8), 49 (VENERA), and 87 (HELIOS). Particles accelerated in a shock initiated by the preceding event on November 21 (Bare No. 1, Bare location indicated by a dashed line in the configuration plot) contribute to the low-energy proton intensity profiles on all three spacecraft. On IMP 8 and VENERA the shock 'S(1)', probably from event No. 1, arrived at the observer after the first particles of event No. 2 were observed. On HELIOS the data do not allow shock identification, however, considering the geometrical configuration (HELIOS is close to the Bare normal for No. 1) and the fact that an interplanetary shock is observed on the other two spaceprobes, we would expect a shock on HELIOS too, because of the smaller radial distance probably before the start of event No. 2. Note the different spatial configurations in events No. 1 and No. 2, the Bare normal in No. 1 is indicated by a dashed line in the configuration. In the November 21 (F1) event thefootpoints of the spacecraft's magnetic field lines are located nearly symmetrically around the Bare site while all spacecraft are east of the Bare normal (other examples for such a configuration will be presented later). This different spatial configuration is also reflected in the number of solar particles: on HELIOS, the spacecraft connected closest to this Bare, protons are observed

14 390 M.-B. KALLENRODE ET AL. ~I ~ JP,I 0.M r-, m~,~r.. ~-,, ù-, C~,-.«- (~1 ~ ~ I O I " ~'~ ù. L."Z ~,' Q L '" ~ ---~~',-----~--"---q--~---h- ~ ~... ~- 7-_%. --~.-~... ~ ~, +,..... " - -J I I i! "" ".....,,-~~ ]«. ~... ~ =,.': "....:.",._; ', ~ ',,o "~ il!... / : 1"- --»M'i... I '. '-~-- r "'-/ ~ o1' t' r I' i' i' t' 1' I' i' 1' O"q'"'i""t""l'"'l'"l'"';;"'l ~. := ~ o ", 0 o o.o ~~«~où~- zo~ ~~ -:.--'" \ :" ~ ~,!,, I, ~ 1,,, ù!... ~,.~ 0 ù. ~... - \ / -~, >>....,:....,..~ ù. - ~~~~ -. ; 7 ~" ~ ~~. ~_ -_,,~..-. ~ ~, :~ ~ -. =.~ i ll ::., : ù" -- -~ " 'T ~ ~ =.~ :~ ~~ ~~", '--" ~ ~,~... ~, : ~,., 1.-,/... '.,:" ~'. '~-: -~..,'..-.»"b, / t- H :], ~-L ] --"-..:r!: " "~~ [ "1... ; ~- "- # ~ ~,~, :~ 'I, ~ ~0 ù, ~ ~~,.I -~g 0 "~ ~ ~ ù'" ~'----."" ~ ~,' "F '. '', ~ :. 7 i-, ~-\ "-..~ i'. ~, ~ : ~ L 0,'m 12, ù"~o - ~--~..~ E " " : ~"1,«,:" - '~ Ii ]:""'".. ;::~f ~ /,J.IS N ":!INI Vl"il~lVO SINV o o o Io Io,LLISN3.LNI ~, 0 ~ I I I C~ I ~,~ ~.,,~ E'I -~ 'ê~ ~ - ~ 0~~ ~! ~~i > ~.l.õ

15 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 391 up to energies of about 30 MeV, on the other two spacecraft the proton intensity increases only in the lowest energy channel (a few MeV) but not in the higher energy channels. Because of the power-law spectrum the intensity of particles accelerated to higher energies was too low to produce any obvious rise above background. In the low-energy channels on IMP 8 and VENERA the intensity continues to increase, probably due to the approaching shock. In the November 22 event the situation is quite different: on IMP 8 and VENERA we observe a relatively large and fast intensity increase in the higher energy channels (not in the low-energy channels because of the high background from the preceding event and the approaching shock), these particles are the component accelerated near the Sun (either in the flare process or by a shock in the corona, cf. Section 1). On HELIOS this increase of the solar component is, at least in the lower energies, rauch slower because of the larger angular distance between the flare and the observer's magnetic footpoint, however, the higher energies give clear evidence for a solar component. Note that this solar component varies in intensity as weil as in time scales in the same way as the intensity in an impulsive flare without shock would vary with azimuthal distance (cf., Kallenrode, 1992). The approaching shock 'S(2)' of event No. 2 leads to different signatures on the three spacecraft: on IMP 8 and VENERA the intensity in higher protons decreases after the maximum of the solar component and starts to increase again as the shock approaches more closely. In addition, on VENERA it seems that a large number of protons are stored behind the shock front because the intensity increases after the passage of the shock. These particles could also be accelerated at the field line after the shock has passed behind the observer, however the anisotropy at the time of maximum of this hump suggests a net streaming of particles from the Sun rather than particles streaming from the shock front towards the Sun. On HELIOS the situation is quite different: after the slow-starting rise in intensity there are some features in the proton profiles where rather abrupt intensity changes were observed without velocity dispersion in protons and also in relativistic electrons. Four dashed lines at the HELIOS profiles mark these times. These changes probably are related to changes in the magnetic connection between shock and observer. The first jump may also be related to the arrival of the interplanetary shock initiated by event No. 1. The last intensity decrease is marked 'FD' and may be a Forbush decrease. The designation Forbush decrease originally is applied to decreases in gaiactic cosmic rays related to the arrival of an interplanetary shock and/or the driver gas of this shock, however, because of the common origin, we have adopted this term for the lower energies too. Until the end of doy 327, particles come preferentially from the solar direction. Around the second dashed line there is some evidence for a bi-directional streaming. After the shock-passage 'S' particles stream towards the Sun until after the Forbush decrease 'FD' there is again a net-streaming of particles from the Sun. Therefore the low-energy protons seem to be dominated by particles accelerated at the interplanetary shock, whereas the higher energy channels on HELIOS give evidence for a solar component early in the event.

16 392 M.-B. KALLENRODE ET AL. Let us add just one remark on the shock geometry in this event. The arrows in the configuration part of Figure 2 indicate the mean shock speed (that is the radial distance divided by the transit time) for the shock initiated by flare No. 2 (November 22) in these directions. Obviously, the shock extends at least -t-40 around the flare site and is able to accelerate particles in this range. Within this rather small longitude range (compared to the -t-90 or more in some other events) there seems to be no pronounced variation of the time profiles with longitude. In addition, the shock is fast close to the flare normal and becomes slower close to its flanks (cf. Table II) The 1982 December 19 Event We will now consider cases in which the flare moves successively to the west with respect to the observers. In addition, these events will become less complex because they are not influenced by preceding flares as strongly as the above event. The 1982 December 19 event (doy 353, Figure 3) originated in a 1B/M9 flare at N10 W75 with an Ho~ maximum at 16:32 UT. In this event the footpoint of the VENERA spacecraft is close to the flare site, while the other spacecraft are connected to positions 40 west and east, respectively. Apart from the absolute intensities, this should be reflected in very similar proton profiles of the solar component on all three spacecraft, probably with faster time scales on VENERA. For the solar component, especially in the higher proton energies, the HELIOS and VENERA spacecraft fit into this picture. For IMP 8 this picture seems to be true as far as the higher proton channel is concerned. However, this increase is rather meager and in the lower proton channel on IMP 8 no increase could be observed, both resulting from the relatively high background preceding the events (cf., Figure 1). Note that IMP 8 and HELIOS have comparable angular distances between the footpoint of the observer's magnetic field line and the flare site. The differences in the intensity profiles of the solar component may result, at least partly, from the larger radial distance of the IMP 8 spacecraft. Let us now turn to the influence of the interplanetary shock. On IMP 8 no sudden commencement, as evidence for an interplanetary shock, is reported, therefore the shock did not extend more than 75 to the east or was too weak to produce a sudden commencement there. The proton profiles also give no evidence for a shock, however we cannot exclude the possibility that farther in the shock had accelerated particles as it propagated outward. The low-energy protons on VENERA show clear evidence for shock-accelerated particles indicated by the long-lasting, nearly constant intensity. Because of a data gap in the magnetic field, no direct shock observations were available. There is also no indication for a Forbush decrease or for the arrival of the driver gas. The higher energy protons (above 25 MeV) on VENERA may be all solar particles with no or only a small interplanetary component contributing in the late phase of the event. Going closer to the flare normal, and therefore to the strongest portion of the shock, we find the interplanetary component becoming, at that time of the event, the dominant particle source in all

17 I MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 393 i ~ x~2m.maaß~ ~... T r,'.:" ~'~ ~ ~ ~!.. t ~ ~" ~- ~~- " '.~.i ' "~ - ~. ~ ~..... :!..'~/_:: ~ ~: t.~,i l ~ ~~~ i ~ ~ "'"'"""., ~L[ISN31NI Vl/t?tVO SINV I-'3 t i i ~ i I i r i 1, I,,l... oo i o_ O oo «, -..'. t : B i 1 õ, ;.!: --I _,,, ù.", e. to H fr3 fr3 0 fr3 /ISN31NI ù "t'""l'"'l",' oq ~). I I I V~RV9 SINV N... 1-"' t '~'' I'"" p",p"" p"' I-"' o o o Io Io ù,t_lisn31ni Fig. 3. Same as Figure 2 for the 1982 December 19 event.

18 394 M.-B. KALLENRODE ET AL. energy channels on HELIOS. In the low-energy protons, the intensity continues to rise, therefore the maximum of the interplanetary component is higher than that of the solar component. In the higher energies the proton intensity decreases after the solar component and starts to rise again as the shock approaches. This raises the question whether the intensity increase around the time of shock passage might be, at least in the higher energy channels, indicative of particle storage close to the shock front rather than of a continuous acceleration and injection. The abrupt decrease in intensity is due to the arrival of the driver and therefore occurs simultaneously in all channels (also in electrons). After this decrease there is a net-streaming of particles from the Sun. This observation indicates that most of the particles accelerated at the shock are not able to escape towards the Sun, while in the low energy channel the continuously rising intensity seems to give evidence for a continuous acceleration and injection in the upstream direction of the shock. While in the low-energy channel on VENERA the intensity of solar and interplanetary particles is comparable, on HELIOS the intensity of the interplanetary particles exceeds that of the solar component by about one order of magnitude. In the higher energies there seems to be no increase due to interplanetary particles on VENERA. On HELIOS the relative amount of interplanetary protons is smaller in higher energies than in low energies, however there is a marked increase in intensity up to energies of about 50 MeV. The shock spike is visible even in the integral channel above 51 MeV The 1982 November 26 Event The 1982 November 26 event (doy 330, Figure 4) originated in a 1N/X4 flare at S12 W87 with Ho~ maximum at 02:36 UT. In contrast to the previous event, here all three spacecraft are located at more comparable radial distances between 0.73 and 1 AU and are located east of the flare normal. HELIOS and VENERA are connected magnetically to positions about 10 west and east of the flare, IMP 8 is connected to a position about 45 east of the flare. On all three spacecraft solar as weil as interplanetary particles are observed. On HELIOS, which is connected magnetically close to the flare site and is also located close to the flare normal, a fast intensity increase is observed, however, the intensity does not decrease, but it stays constant in high energies or increases in lower energies for about 12 hours. This seems to be a combination of solar particles and particles accelerated by the interplanetary shock on the magnetic field line between the Sun and the observer. A clear separation of these two particle populations by means of the intensity and anisotropy profiles only is not possible, but observations for well-connected events, such as this example, indicate the presence of a solar particle component early in the event (cf., Reames, Cane, and von Rosenvinge, 1990, especially their Figure 4). The passage of the shock leads to a pronounced shock spike in all proton energy channels with an intensity increase by up to two orders of magnitude. Before the shock spike, particles come from the solar direction; after the shockpassage particles stream preferentially towards the Sun. The shock therefore acts as

19 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 395 a moving source of particles rather than as a closed reservoir of particles which is simply swept across the observer. The slow intensity increase late on November 27 may be related to another injection. The magnetic connection for VENERA is very similar to that of HELIOS, however the distance to the flare normal is much larger. Therefore the rise of proton intensities is similar to that on HELIOS with slightly longer time scales due to the larger radial distance, however the later times that are dominated by shock-accelerated particles on HELIOS are different on VENERÄ: in the low proton energies the intensity continues to increase and forms a shock spike (we marked that time as the time of shock passage; magnetic field data for a direct identification of the interplanetary shock were not available), which is similar to the low-energy protons on HELIOS. In higher proton energies there seems to be a small interplanetary proton component becoming dominant only for a short time above the residuals of the solar component at that time and being responsible for the structures in the decaying flank. This interplanetary component peaks.well before the arrival of the interplanetary shock., On IMP 8 the intensity increase is much slower and smaller, as expected from the larger angular distance between flare and observer's magnetic footpoint. There is no shock observed at IMP. The protons of higher energy seem to be mainly solar particles and look comparable to the high-energy protons on VENERA. In the lower-energy proton channel the profile is rather long-lasting and very flat around the time of its maximum. This late and flat maximum can be understood in terms of continuous shock acceleration of particles on the observer's magnetic field line as the shock propagates outward; however, due to the curvature of the field line to the east, the shock does not reach the observer The 1982 December 7 Event In the 1982 December 7 event (doy 341, Figure 5) the configuration is very similar (only HELIOS has moved to a somewhat smaller radial distance). The parent flare was a 1B/X3 flare at S19 W89 with an Ha-maximum at 23:40 UT. The intensity profiles are very similar to the profiles in the previous event, the main difference lies in the interplanetary component: as in the November 26 event, the interplanetary component is dominant in the low energies on HELIOS but, in relation to the solar component, is smaller in the higher energies. However, the interplanetary component is still dominant up to energies of about 30 MeV. In these higher energies it even seems possible to distinguish between these two components, while in the lower energies the transition is more continuous with a strong influence of the interplanetary component relatively early in the event. In contrast to the previous event no shock spike is observed, but after the shock passage a decrease indicates the arrival of the driver gas. Particles come from the solar direction until, after the shock passage, there is anet streaming of particles towards the Sun. Therefore it seems that, as in the previous event, the shock is a moving source of particles and injects, even after the decrease, a relatively large

20 M.-B. KALLENRODI~ ET/kL. 396 ~ ~ _,.),It ù.:"-... -t... ~ ~ ',,~ }-... ~... ~., ".. -- le) I'-"... _~~ i ~~.~ ~ - ~,!... A / : :~ : " ~', i '.. ~ ~ ~-i.: ~!,- ù-~i... "i r ~ :." t - " - ~~~_. "~... i... ~".,..~.:..",:-.~ ~,-:, - v: ~"-~ "~, " i... :..õ... ~ i'_l' ~.,v,),,.,:.. -- ~0 ~ ""... ~ ~ -- ~' ~~" 'tl ~ ~ I0 ~ ~ ~ I ~ i. I X,~ ~.../.." " " ~ " / fr'} ~~ " i~~~~i!!!i! ~I i~ I'0.. ~-,_-r'' ~-'.L ««~. ~T : 1 ~, n o ~,, o ~ ' ~':: ~LLIS N':I N I Vl/~IRVO SINV.,.. 0 kl. ~ ; : Iwr) I0 I0 ~=. ~ AZISN3 NI Fig. 4, Same as Figure 2 for the 1982 November 26 event.

21 MULTI-SPACECRAFr OBSERVATIONS OF SHOCKS 397 number of particles. This is important insofar as, contrary to the earlier belief, the intensity profile after the shock passage is not the profile one would get without the influence of the interplanetary shock. (The ESP component was considered to be superposed only around the time of shock-passage and consists mainly of particles stored around the shock front.) Note that in the December 19 event, after the Forbush decrease, a net streaming of particles from the Sun was observed, indicating that in that particular case the shock has no or only a small number of particles injected toward the Sun. On VENERA the higher energy protons seem to be the solar component alone and are not, or only to a very limited amount, increased by a possible interplanetary component. In contrast, in the lowest proton channel the shock influence is clearly visible as a long-lasting, slow intensity increase. The picture on IMP 8 looks very similar to the profile on VENERA, the shock arrives at doy 344, 07:21 UT, outside the time frame used in the figure. In general, in this event the shock has, compared to the intensity of the solar component, a smaller effect on the intensity profiles than in the previous event. However, the qualitative changes, especially the dependence on the angle to the shock normal, are the same as in the previous event. The shock, and also the ability of the shock to accelerate particles, extend up to an angular distance of 90 from the flare site, however the shock has a much lower mean speed on IMP 8 and VENERA than on HELIOS (cf., arrows in the configuration plot in Figure 5 and Table II). Note that HELIOS is close to the Sun, therefore the curvature of the interplanetary magnetic field line is rather small and HELIOS is connected roughly to the same point on the shock surface during all the time until the shock hits the spacecraft. This is in contrast to observers at large radial distances where the magnetic connection point moves east on the shock surface as the shock moves outward The 1982 December 26 Event In the 1982 December 26 event (doy 360, Figure 6) the flare occurs further to the west with respect to the spacecraft. For this event no association with a flare on the visible disk was possible; a flare behind the western limb close to W126 seems to be very likely. HELIOS is magnetically connected close to that position and is also close to the flare normal. A striking feature in the HELIOS proton intensities is the abrupt intensity decrease at the time of shock arrival in the low-energy protons, combined with an intensity increase in higher proton energies and electrons and anet streaming of particles towards the Sun, which again indicates that the shock seems to be a moving source of particles. In addition, the time of shock passage coincides with a reversal pf the magnetic field direction for about 6 hours. After this depletion in low-energy protons, the protons come again from the solar direction. On VENERA only data for the lowest proton channel are available. This channel shows a slow rise, particles streaming continuously from the solar direction until the arrival of the interplanetary shock. A distinction between a solar and an interplanetary component seems to be difficult. On IMP there seems to be an

22 398 M.-B. KALLENRODEETAL, ~., «.-... T... -,. ù ""-... ù. ù,! t~ t~ I,I ù...-..; , ù %,.,: :" ù.. ; ("4 I *... ~... :~... ' I' I' I' I' I' i' I" I' I' ' I ' \\~---~~~- % i% % % o,o \ ~..-~ ~~ I I ) ~..'/ LJ.ISN3.LNI VIBII,~V~) " ",-. i....., """ SINV ù'"",, f,, I,, I,,, I, I, m fr) ' t 3 : / ù J.f " i,: - r,,~ j : t~ / { i te), r~ C',I, te) l.d i : ù ": i ùt" -.~. t/~ ~ 0 t/~ I A.LISN3.LNI ' I ' t'n, J I ff'~off)'- d d i I StNV IQ I 0 L.LISN31NI Fig. 5. Same as Figure 2 for the 1982 December 7 event.

23 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS I, I, I, I, I, Il II Il I, Il,Ndji,,l,,I,ùI I ~ I, I / ;//>eee- t ~~~_~_, - ~ ~ _,, ( ' J~ ~ i" i i:,: " "~ " -. rs--. ~:. ~ ' [... i... i i : ' ù...,'"! "'"'""'- I : : t~ o~ ---~ -."...:...,... ~ ~ ~ ; '... ~, :~: :-- j..._ :..'..... i ò-- ' I' I' I' I t I' I"l' 1' : I j '"'P'"P"'I'"~i " I ' ]" I ' I ù %%%% O ~, tt~ rd{'~ Q 399 c-,i - cd O D Q cd,llisn3j.ni VI~I~ O SINV... i >... ~~ --, ~ I ~\,Ir, I,,,,h,,,hù~hù~l... hù,l,,,j,,,,i I~~t~ I="~ t-, t--,, k.ù t'-,~ ~, t-, t'-, ~ -~ ~ r~ ù ~.=~..,- - ~. L"~ = w i- ù~ p;ù im,, p;ù i=,:, p.ù p,, p,, I o L.-I_ISN31NI M~~'9 ~SINV I.LISN3JNI Fig. 6. Same as Figure 2 for the 1982 December 26 event.

24 400 M.-B. KALLENRODEET AL. intensity increase after the flare; however, the background is relatively high from a preceding event (December 25) with parent flare at E45 (cf., middle panel in Figure 6 in CRvR). The shock originating in this December 25 flare is indicated in Figure 6 and is associated with a new increase in intensity. There is no indication for an interplanetary shock originating in the December 26 event on IMP 8. This is not surprising, because a shock with an angular extent at least 130 east of the flare normal would be required to produce a sudden commencement at Earth. A combination of the VENERA and HELIOS shock observations confirmed the picture which had evolved above: the shock is fastest close to the flare normal and the shock speed decreases to the flanks of the shock SHOCK PROPERTIES 4. Discussion Properties of interplanetary shocks are of crucial importance to understand the generation and distribution of energetic particles in the interplanetary medium (see, e.g., Reames, 1993, for a recent summary). Fast shocks are required to accelerate particles efficiently (cf., e.g., Cane, von Rosenvinge, and McGuire, 1990); the shock supplies particles to field lines not connected to the flare region (cf., CRvR), therefore the azimuthal extent of the shock controls the distance up to which particles can be found; and variations of the shock speed with azimuth control the large-scale shock geometry, leading to average shock normals different from the radial direction (see, e.g., Richter et al., 1981, for an instructive example). From the individual events discussed above the following picture concerning shock geometry evolves: (1) Shocks can extend up to more than 90 east of the flare normal. If we assume a symmetrical configuration around the flare normal this would imply a total extent of the shock of 180, and (2) shocks are faster close to the flare normal than close to the flanks. Observation (1) is not only supported by the individual events discussed above, but also by the other shocks in Table II. This total extent is larger than 100 as suggested by Cane (1988). Note that the observed extent of 90 to the east is larger than the typical span of CMEs of 60 (cf., Sheeley et al., 1984; Howard et al., 1985; Kahler, 1987; Hildner, 1992). Cane (1988, and references therein) therefore suggested that there is not a one-to-one correspondence at each position between the driver (CME) and the more extended shockfront. Such a view is supported by our observations, directly showing the large extent of the interplanetary shocks. We have also noted that in several cases the arrival of the driver is only noticed on the spacecraft closest to the flare normal. The shock geometry is depicted in Figure 7, where we show the distribution of shock speeds in the different directions relative to the flare normal for the five events No. 2, 3, 6, 8, and 10. For comparison, the dashed curve shows a circular shape. The arrows are normalized in such a way that arrows of different events

25 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 401 (indicated by the different styles of the arrows, e.g., solid, dashed, etc.) have a comparable length close to the flare normal. In constructing Figure 7 we used the transit speeds of the shock. A shock, however, can accelerate or decelerate during its propagation. Therefore the shape depicted in Figure 7 might be biased by radial changes of the shock speed, especially when we consider that for the examples presented here the observer close to the flare normal in general was also at a smaller radial distance. However, we do not believe that this effect causes drastic changes. Let us assume that the shock is originally spherical and our shape comes from deceleration alone. In that case, up to the radial distance of the observer closest to the Sun, the shock has had the same transit speed and the same travel time also for the observer at the larger distance. Let us take as a numerical example event No. 3 which fits weil in the general shape depicted in Figure 7. The shock arrives at HELIOS at a distance of 0.72 AU after a travel time of about 15 hours, implying a transit speed of 1990 km s -1. VENERA 13 is at a larger distance of 0.85 AU, the shock arrives there about 22 hr after the flare or about 7 hr after it had arrived at the radial distance HELIOS is located at. Thus if the shock was spherical it would have propagated with a transit speed of 1990 km s-1 to a radial distance of 0.72 AU (HELIOS) and would have continued the remaining 0.13 AU with an average speed of 0.13 AU/7 hr=774 km s -1 to the radial distance VENERA is located at. This would imply a drastic drop in shock speed by a factor of nearly three. To our knowledge, no example for such a strong and abrupt change in shock speed has ever been reported. We would like to note that Wibberenz and Cane (forthcoming paper) find a similar shape of the shock, but in their examples the observer close to the flare normal is not always at the smaller radial distance but at a comparable or even larger distance. Note that the qualitative shape in Figure 7 is quite similar to the general shape of the shock in the meridional plane as derived from Doppler-scintillation measurements (Fengsi and Dryer, 1991) and to the shock shape and velocity distribution in the equatorial plane as inferred from the numerical calculations of Dryer et al. (1984). Observation (2) is in qualitative agreement with the inferred transit speeds for shocks (Cane, 1988); however, this analysis of multi-spacecraft events suggests a stronger decrease of shock speed to the flanks than suggested by Cane on the basis of a statistical study. This discrepancy may have several reasons: (a) the shocks in this paper are partly very fast with shock speeds up to 2000 km s -1 close to the flare normal, and in these fast shocks the speed might decrease faster as one goes to the flanks. (b) Detection of shocks and association of these shocks with a parent flare becomes more complicated as one goes to the slower flank of the shock, therefore observations from a single point in space may not easily identify the slow flanks of a shock properly. This may lead to a preference for relatively fast shocks at larger distances from the flare normal. (c) An existing variation of the shock speed with longitude might be reduced in a statistical study when averaging

26 / I / \ \ 402 M.-B. KALLENRODE ET AL. SUN ~--...,3~11 ~ ----,,, \ / / / Fig. 7. circular shape. FIaPe Normal Shock geometry as inferred from four events. For comparison, the dashed curve shows a over various shocks with different speeds in a certain longitude band. In principle, the non-spherical shock shape might have implications on our picture about the acceleration efficiency at different longitudes. Independent of how the interplanetary magnetic field is distorted by the shock, the angle Oßn between the magnetic field direction and the shock normal will be different for the eastern and the western flank of the shock, leading to a larger Oßn at the western flank of the shock front compared to the eastem flank. However, we do not believe that this difference is large enough to provide a transition between different acceleration mechanisms as suggested by Sarris, Decker, and Krimigis (1985). Note that there are other differences between the shock properties on the eastem and western flank which might be more relevant for particle acceleration in interplanetary space: (a) there is stronger compression of magnetic field lines towards the western flank (Cane, 1988; CRvR) leading to larger compression ratios. (b) Because of the spiral shape of the magnetic field lines observers on the western flank are connected much longer with the front of the expanding shock than on the eastem flank; this leads to a larger number of accelerated particles at a given point of observation.

27 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS PARTICLE ACCELERATION Let us assume that for the large gradual events discussed here shocks starting close to the Sun are responsible for most if not all of the energetic particles. In spite of this common source we have attempted to define two separate components for these events, based on dynamical properties of the particle population (see Section 1). In these events the solar component is characterized by a smooth decay of the anisotropy and a weakening of the energy spectrum. This behaviour can be quantitatively modelled by a finite injection of particles close to the Sun combined with interplanetary diffusion. An additional component which we have termed interplanetary is in general characterized by a new increase of the anisotropy and a further weakening of the proton energy spectrum (for details see Meyer et al., 1992). In many cases this component reaches maximum intensity close to the arrival of the interplanetary shock. The relative importance of the two components, in terms of the absolute intensity of the accelerated particles, varies with particle energy as weil as the radial and azimuthal distance of the observer from the flare. This question - which part of an intensity profile is dominated by the solar and which by the interplanetary processes - is also discussed by CRvR, Reames (1990), and Meyer et al. (1992). For particle acceleration, especially of the interplanetary component, we can draw the following conclusions based on our observations: (3) The acceleration of protons at an interplanetary shock can extend to energies of at least 50 MeV. (4) Shock acceleration is most effective close to the flare normal, i.e., the fastest portion of the shock. In particular, it seems that the acceleration of high energy protons (above about 20 MeV) occurs preferentially close to the flare normal and ceases towards the eastern flank of the shock. (5) The shock may still be very efficient towards its flanks in accelerating low-energy protons (of the order of a few MeV). (6) All events with high intensities in the high energy range were observed to have good magnetic connection between the footpoint of the observer's magnetic field line and the flare location, and all have a solar component. (7) The anisotropy after the shock passage is directed away from the shock front in some events. This suggests an interpretation of the shock as a continuously acting moving source of particles (e.g., in No. 3 and No. 6). In other events the shock could be interpreted as a more or less closed reservoir of accelerated particles that is swept across the observer (No. 9), with low-energy particles escaping preferentially into the upstream region of the shock. (8) Direct observations indicate strong particle acceleration at the interplanetary shock inside 0.5 AU. (9) Shock spikes can occur up to energies of 50 MeV and more. The dependence of acceleration efficiency on the shock geometry, especially the high efficiency for acceleration close to the flare normal (observation (4)), is a well-known phenomenon and is discussed in CRvR and in Cane, Reames, and

28 404 M.-B. KALLENRODE ET AL. von Rosenvinge (1991). While the CRvR work is based on a statistical analysis of a large number of events observed by only one spacecraft, this work confirms the results in CRvR for individual events observed at three different points in space. The acceleration of protons up to energies of 100 MeV at the interplanetary shocks (observation (3)) is also weil known and discussed in detail in Reames (1990). According to observation (4) high energy protons are preferentially accelerated close to the fiare normal. At larger angular distances between spacecraft and flare normal there can still be a very efficient acceleration of low-energy protons. This also means that for low energies the interplanetary component of a large energetic particle event may be dominant compared to the solar component, in contrast to high energies, where the highest absolute intensities during an event generally occur early in the event. This different behaviour of low and high energies is also reflected in characteristic changes of the proton spectrum (cf., Meyer et al., 1992) with the approach of the shock front. It should be noted, of course, that this net effect may be a combination of acceleration and propagation processes, because protons of higher energy with a larger diffusion coefficient can move ahead of the shock front faster than low-energy protons, thus leading to a larger portion of low-energy protons in the vicinity of the approaching shock. How can we understand this distinction of a solar and an interplanetary component in terms of particle acceleration at the shock only? From profiles observed for the high energy protons there is a clear indication that in many events there is a first intensity maximum of the solar component (within a few hours after the flare), then the intensity decreases in a manner more or less similar to the intensity decreases in events not accompanied by an interplanetary shock (e.g., impulsive events), and close to the arrival of the shock the intensity rises again to a second ma:timum around the time of shock passage. We interpret this increase as being due to the proximity of the shock. For larger distances of the shock from the observer, the relative contribution of the number of particles at the observer decreases accordingly. This means that we need an additional strong source when the shock is very far away, that is still close to the Sun. We need therefore a very efficient phase of shock acceleration while it is still in the solar corona in order to explain the well-defined solar component with hard energy spectra and a large total number of particles. For 175 MeV protons, Kahler, Reames, and Sheeley (1990) plotted the intensities versus the height of the leading edge of the CME in the corona instead of time and showed, that these intensities peak at heights between 6 and 10 solar radii. Considering the elongation of the time profiles due to interplanetary scattering, these heights are already upper limits for the acceleration height of the particles. On the other hand, the shock without its component close to the Sun should lead to a slowly rising intensity profile (the shock as an approaching source of particles) with a maximum around the time of shock passage. Such kinds of particle profiles of a 'pure interplanetary component' are observed (even in energy ranges of some tens of MeV, cf., Cane, von Rosenvinge, and McGuire, 1990; Kallenrode and

29 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS 405 Wibberenz, 1992). Therefore the complete intensity profile can in simple terms be described as a superposition of the solar component (extended injection with interplanetary propagation) and the shock-accelerated profile with a maximum around the time of shock passage. This view is not only in agreement with the observed intensity profile (especially observation (6)), but also with the time evolution of the proton spectrum and the observed anisotropy, at least in the events in which the particles stream sunwards after the shock passage (observation (7)). For a few cases there might be an alternative interpretation for the interplanetary component. This relates to events in which no sunward streaming of particles is observed after the time of the shock passage (observation (7)). Here we might argue that the acceleration efficiency of the shock is smaller and therefore the contribution of the interplanetary component is mainly restricted to the time around the shock passage. In principle, the intensity increase might be due to particles stored in the turbulence around the shock front rather than accelerated locally at the shock. However, this picture can at best hold for a few cases and for energies where a small hump is superimposed on the decay of the solar component. In sum, continuous acceleration at the shock travelling outwards from the Sun seems to be a reasonable assumption. However, we still have not considered how the acceleration efficiency can vary with distance from the Sun. Possibilities for physical mechanisms include the plasma density (as a measure for the seed population available for injection), the shock geometry, or the shock strength. As far as the density is concerned, in the corona the shock would operate in a high density medium which is not simply scaled with r -2. As a consequence, we might also obtain a large number of accelerated particles close to the Sun. Farther out, when the solar wind has reached its nearly constant expansion velocity, the plasma density varies approximately with r -2, thus supplying a decreasing number of seed particles. All particles accelerated by the expanding shock will have to compete with the solar component, which at late times shows little radial variation, but decays in time. So, whether or not one sees a hump of the interplanetary component becoming dominant over the solar component depends on details of the interplanetary transport as well as on the shock efficiency and requires detailed quantitative modelling. Let us turn next to the shock geometry. The field line close to the Sun is nearly radial, therefore the shock is quasi-parallel. However, as the shock propagates outward, the field line curves to the east, leading to an oblique geometry, at larger distances, especially at the western flank, even to quasi-perpendicular shocks. Observation (8) clearly shows that the shock within 0.5 AU, when it is quasiparallel, is able to accelerate large numbers of energetic particles. This means that we do not have to invoke the condition of quasi-perpendicular shocks to have very efficient particle accelerators. This is in agreement with theory which suggests that, especially at higher energies, quasi-perpendicular shocks are less efficient particle accelerators (cf., Jones and Ellison, 1992, and references therein). But how does this situation change as one goes farther out? For the eastern

30 406 M.-B. KALLENRODE ET AL. flank the increasing spiral angle is partly compensated by the large-scale shock front which deviates from a hemispherical shape, and therefore the geometry still is quasi-parallel. Extrapolating our geometry of the eastern flank of the shock to a symmetric shock shape, we would get a smaller Oßn at the eastern flank, compared to the western flank, which would imply a rather perpendicular geometry (cf., Figure 9 in Sarris, Decker, and Krimigis, 1985). Therefore we would expect a strong east-west asymmetry in the total number of particles related with the arrival of the shock if the shock efficiency were largely controlled by the angle 0B,» Observational results on this point are unequivocal: CRvR suggest in their sketch of representative profiles (Figure 15 in their paper) that the intensity of the shock-accelerated component is small at the eastern flank and large at the western flank without a pronounced angular dependence on the western flank, while Meyer, Wibberenz, and Kallenrode (1992) report indications for a broad scatter of the sizes of the interplanetary component with no evidence for a symmetric distribution of event sizes. However, going to the flanks of the shock, the acceleration efficiency always decreases for observers both east and west of the flare normal. We conclude tentatively that the property of the shock, quasi-perpendicular versus quasi-parallel, does not play a crucial role for the acceleration of particles in the several to tens of MeV range. In addition to the geometry, shock-inherent parameters, such as the shock speed or compression ratio, might determine the acceleration efficiency. Cane, von Rosenvinge, and McGuire (1990) ordered the different kinds of profiles in events accompanied by a CME and an interplanetary shock as a function of CME speed. If the shock produced both a solar and an interplanetary component, the CME speed was higher than in events which consisted of an interplanetary component only. In addition, the size of the interplanetary component is larger in events with high in - situ shock speeds than it is in events with lower shock speeds. On the other hand, Kallenrode and Wibberenz (1992) argued that simple shock parameters as local shock speed, compression ratio, and Mach-number alone are not sufficient to describe the influence of an interplanetary shock on the energetic particles they only report a general tendency that events that are large in these parameters also tend to be accompanied by efficient particle acceleration. Anyhow, it seems to be hopeless to find a single shock parameter which is capable of ordering the efficiency of particle acceleration. Apart from the variation of the shock speed with distance from the Sun, we are faced with the situation that the total number of particles is determined by the connection of the observer to the shock and the resulting local geometry at the shock front while the shock propagates out from the Sun to large distances. To understand the dependence of the acceleration efficiency on local (Oßn) and global (angle to the flare normal) geometry, as weil as on shock properties, it seems important to extend the analysis to distances beyond 1 AU where shocks in general are quasi-perpendicular. Observations of high-energy protons accelerated at interplanetary shocks at larger distances have been reported (e.g., Krimigis,

31 / MULTI-SPACECRAFI" OBSERVATIONS OF SHOCKS ; Wibberenz et al., 1992); however, it is not clear whether all shocks beyond 1 AU are able to accelerate protons to energies of some tens of MeV, or whether this acceleration is limited to extremely strong shocks or systems of shocks (cf., Wibberenz et al., 1992, and references therein) RELATION TO SUPEREVENTS By the definition of DMC the whole time period discussed in this paper is a superevent (No. 14 in their Figure 1) in the sense that (a) the 54-day running mean intensity stays above background for about one month on spacecraft at different radial and angular positions throughout the solar system, and (b) the time period is accompanied by a decrease in 35-day running averages in neutron monitor data. In protons this super-event was observed by IMP 8 at 1 AU as weil as by PIONEER 10/11 at a radial distance of about 29/13 AU. An electron increase is observed at IMP 8 but not on PIONEER 10. Superevents are assumed to form out of a series of large solar flares, especially in geometrical configurations where the flare initiated shocks were able to form some closed shell around the Sun, confining and accelerating particles. The original definition ofa superevent, however, is slightly different (see MRW): in high time resolution data, e.g. hourly averages, there should be a long-lasting (more than one month) intensity increase by at least one order of magnitude above background that cannot be related to individual solar events or the approach of individual interplanetary shocks. Some larger solar events stick out of that profile, but the important fact is the long-lasting, smooth intensity profile which simultaneously increases for several channels, including MeV electrons and protons with energies of tens of MeV. We will refer to this original idea as superevents of type I, superevents in the sense introduced by DMC will be referred to as type II. We think that the time period under study represents a superevent of type II rather than one of type I. Although the intensity is above background on HELIOS, IMP 8, and VENERA for an extended time period, the profiles seem to be strongly dominated by the large individual events summarized in Table I and the time periods between these large events are characterized by a large number of smaller flares (as indicated in Figure 1). Therefore an averaging procedure should produce a relatively high intensity for a long time. The data seem not to fit the original idea of some smooth profile unrelated to specific solar activity, and not even related to non-local acceleration at an interplanetary shock as it propagates from the Sun to the observer, especially in the high-energy protons. Depending on the definition of superevents one prefers, our observations seem to indicate that an active Sun producing a fair number of strong energetic particle events, and also some smaller particle events, should easily be able to initiate a superevent of type II, however it seems that for the formation of superevents of type I, some additional feature is required. The formation of a closed shell at large distances seems to be an attractive idea for the acceleration and storage of these particles (cf. DMC). Considering the fact that nine strong interplanetary shocks

32 408 M.-B. KALLENRODE ET AL. were initiated in the western hemisphere during this time period, and that these shocks have different propagation speeds, the merging of shocks seems to be very likely, leading to a highly efficient acceleration (cf., Levy, Duggal, and Pomerantz, 1976; Wibberenz et al., 1992). For example, the first three shocks in Table II seem to be likely candidates for such a merging: the shocks are initiated within five days on the western hemisphere with increasing shock speeds: the first shock has a shock speed of roughly 900 km s-l; 1½ day later a shock with a shock speed of 1200 km s-1 starts about 40 east of the first one, and again 3 ½ day later an extremely fast shock with a speed between 1500 and 2000 km s -I starts from the same position as the first one. All three shocks have large angular extensions (presumably of at least 180 ) so it is very likely that the faster shocks will catch up with the slower ones: shock No. 2 would catch up with shock No. 1 about 6 days after the start of shock No. 1 at a radial distance of about 3 AU; shock No. 3 would catch up with this merged shock after an additional 4 days at about 5 AU. Assuming that the active regions that initiated these shocks during their rotation have also initiated shocks at other longitudes, the formation of a closed or nearly closed shell seems to be possible. In addition, these merged shocks might turn out to be much more efficient particle accelerators than single shocks would be beyond 1 AU. However, considering the intensity time profiles in Figure 1, even if such a closed shell has formed at larger distances, it seems to be unlikely that a large number of particles are accelerated at this shell and travel back into the direction of the Sun, or that the solar particles accelerated in later flares are trapped very efficiently in this shell, because the intensity, especially in the higher energy channels, decreases between some of the events to background or nearly background levels. While this holds for the 'global' event we cannot exclude local trapping of particles between individual shocks. An injection of shock-accelerated particles into the direction of the Sun after the shock has passed behind the observer might lead to a slower decrease of the intensity profile. Note that this lack of significant acceleration seen at or within 1 AU does not exclude the possibility that the shell is responsible for the simultaneously observed modulation of the galactic cosmic rays or for the energetic particles seen at larger distances from the Sün. In sum, this time period seems to be an example of an active period rather than of a superevent of type I. But it seems that a similar analysis conducted with data obtained at larger distances (e.g., observations as presented in McDonald and Selesnik, 1991, for a different time period) should help us to understand the evolution of the individual and merged shocks. In addition, a comparison of intensity time profiles with the interplanetary shocks should give first hints on whether the particles observed at these large distances are accelerated locally close to the observer or are mainly accelerated in the inner heliosphere.

33 MULTI-SPACECRAFT OBSERVATIONS OF SHOCKS Summary One result of the present study concerns the shock geometry, especially the sometimes large angular extent of the shock of more than 90 to the east of the flare normal, and the decrease in shock speed as one goes to the flanks of the shock. The second result is related to the acceleration efficiency of the interplanetary shock for protons: shock acceleration is most efficient close to the flare normal and decreases to the eastern flank. This decrease is reflected in smaller intensities and lower maximum energies, however, in the low-energy protons shock-accelerated particles are observed even at angular distances of about 90 to the flare normal. A solar component can be identified on all spacecraft and is, compared to the component accelerated at the interplanetary shock, larger at the eastem flank than close to the flare normal. This corresponds to the usual well-connected events. The shock close to the Sun is a very efficient particle accelerator; to explain the solar component in terms of shock acceleration the acceleration has to become less efficient as the shock propagates outward. This is, at least qualitatively, in agreement with our observations. A valuable tool for understanding the time profiles in detail and also the evolution of the interplanetary shock seems to be an attempt to model shock acceleration in interplanetary space by solving the transport equation not with an injection function fixed in space and extended in time, but with an injection that moves outwards, or in other words considering the interplanetary shock as a moving source of energetic particles. Some kind of 'fitting' of individual events by such a method might turn out to determine the injection function at the shock depending on radius. This would help to understand the evolution of the acceleration efficiency during the outward motion of the shock, especially the question whether the shock has to be more efficient as long as it is quasi-parallel (within about 0.5 AU), or whether it becomes more efficient as it becomes more quasi-perpendicular. Acknowledgements We are grateful to all members of the University of Kiel HELIOS team and to all members of the Moscow State University VENERA team. We thank J. Meyer and H. V. Cane for helpful discussions, the National Space Science Data Center A (Goddard Space Flight Center) for providing the IMP 8 particle data, and R. Schwenn and G. Musmann for providing solar wind and magnetic field data. We wish to express our thanks to the referee for many helpful comments which helped to clarify the outline of the paper. This work was supported by the Deutsche Forschungsgemeinschaft under contract Wi-259/8-1. References Belyakov et al.: 1984, Cos. Res. 22, 738. Bryant, D. A., Cline, T. L., Desai, U. D., and McDonald, E B.: 1962, J. Geophys. Res. 67, 4983.

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