Sharp boundaries of small- and middle-scale solar wind structures

Transcription

1 JOURNAL OF GEOPHYSICAL RESEARCH, VOL. 110,, doi: /2005ja011307, 2005 Sharp boundaries of small- and middle-scale solar wind structures M. O. Riazantseva 1 and G. N. Zastenker Space Research Institute, Russian Academy of Sciences, Moscow, Russia J. D. Richardson Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, Massachusetts, USA P. E. Eiges Space Research Institute, Russian Academy of Sciences, Moscow, Russia Received 7 July 2005; revised 30 September 2005; accepted 5 October 2005; published 24 December [1] This work reviews the characteristics of sharp (less than 10 min) and large (> cm 2 s 1 ) solar wind ion flux changes which are not due to shocks. These changes are boundaries of small- and middle-scale solar wind plasma structures. We present examples and statistical results from simultaneous plasma and magnetic field measurements by Interball-1 and Wind from 1996 to The behavior of the solar wind bulk velocity, temperature, and interplanetary magnetic field (IMF) during these changes is described. We show that in many cases pressure balance is not maintained across these structures and discuss possible implications. In 50% of the events the sharp boundaries are pressure balanced structures; these events are probably tangential discontinuities. We use multipoint measurements of the sharp discontinuities to determine the orientation of the plasma fronts; many fronts are at large angles to the solar wind flow direction. We show that these events can be geoeffective and produce sharp changes in the magnetic field at geosynchronous orbit and on the ground. Citation: Riazantseva, M., G. N. Zastenker, J. D. Richardson, and P. Eiges (2005), Sharp boundaries of small- and middle-scale solar wind structures, J. Geophys. Res., 110,, doi: /2005ja Introduction [2] The solar wind is highly variable by nature. The physical properties of the solar wind have large-amplitude variations over timescales ranging from seconds to years [Hundhausen, 1972]. Concentrating on the shorter scales, solar wind variations range from fractions of a second (plasma waves and noise) to tens of hours (variations related to inhomogeneities of the solar corona) [Feldman et al., 1977; Schwenn and Marsch, 1991]. Gosling et al. [1977] and Shodhan et al. [1999] studied solar wind pressure changes on the scale of hours with relatively slow (tens of minutes) onsets. Particularly interesting features of the solar wind are the frequent, large variations in ion flux that occur on timescales of less than a second to several minutes but are not shocks [Dalin et al., 2002a; Dalin et al., 2002b; Riazantseva et al., 2003a; Riazantseva et al., 2003b]. These variations are characterized by sharp, large-amplitude changes of the solar wind dynamic pressure. Study of these large, sudden changes of the solar wind ion flux (and 1 Also at Skobeltsyn Institute of Nuclear Physics, Moscow State University, Moscow, Russia. Copyright 2005 by the American Geophysical Union /05/2005JA pressure) should lead to better understanding of the nature and properties of the solar wind plasma, of the stability of such structures or of the character of plasma instabilities which cause these steep fronts, and of the effect of these features on the magnetosphere. [3] This paper compiles and reviews results from proceedings and Russian language journals so that a summary of this work is readily available and also presents new results. We first define the criteria for identification of these features and show examples. Next, we present statistical results on the occurrence rates of these features and their properties. We then investigate whether these structures are in pressure balance. Next, we look at the orientation of these fronts and their geoeffectiveness, and finally we discuss and summarize our results. 2. Data Used in This Study [4] We use Interball-1 solar wind ion flux measurements from 1996 to Interball-1 was designed to investigate Sun-Earth connections, focusing on Earth s magnetosphere and the solar wind [Galeev et al., 1996]. It was launched on 3 August 1995, into an orbit with an apogee of about 200,000 km, a perigee of ,000 km, and an inclination of 62. The orbital period was about 4 days. During 8 months (from February to October) of each year the 1of11

2 propagation time from Wind to Interball-1; we assumed the flow was radial and used the solar wind speed measured by Wind. The coordinates of the spacecraft in the geocentric solar ecliptic (GSE) coordinate system were (X INT =12R e, Y INT = 6 R e,z INT = 13 R e ) for Interball-1 and (X Wind = 58 R e,y Wind =65R e,z Wind =3R e ) for Wind. An interesting feature of this event is that the solar wind ion flux change is accompanied by very little change in the bulk velocity, the thermal speed, the IMF magnitude, or the IMF direction. Hence in this case the flux changes are driven solely by plasma density variations. Figure 1. The shaded region highlights a sharp (6 min) decrease of the solar wind ion flux observed by Interball-1 and Wind on 26 June The panels show (a) the ion flux, (b) the radial velocity component (V X ) and the thermal speed (V th ) of the solar wind, (c) the IMF magnitude jbj, and (d) the polar and azimuthal angles of the IMF, f and q. satellite was in the solar wind near apogee. From 1995 to 2000, Interball-1 spent more than 15,000 hours in the solar wind. The solar wind data used in this study are at least 1 hour from the closest bow shock crossing. Data when Interball-1 was in the foreshock region were removed. The ion data were obtained by an integral Faraday cap [Zastenker et al., 2000] and have a time-resolution of 1 s or better. We also use interplanetary magnetic field (IMF) data from the Interball-1 flux-gate magnetometer FM-3I [Nozdrachev et al., 1998]. Plasma and magnetic field data from Wind, ACE, IMP-8, and Geotail were obtained from the NSSDC. 3. The Phenomenon and Its Properties [5] A large number of events with sudden, large-amplitude solar wind ion flux changes were selected from the Interball-1 data. We, somewhat arbitrarily, chose events in which the flux changed by more than cm 2 s 1 over a time period of less than 10 min. Figure 1 shows a typical event observed by both Interball-1 and Wind. The solar wind flux decreased by almost a factor of 2 in about 6 min. The Wind data were time-shifted by the solar wind 4. Occurrence Frequencies and Characteristics of Ion Flux Changes [6] About 10,000 hours of Interball-1 solar wind observations were used to determine the occurrence frequency of the sharp solar wind ion flux changes. Figure 2 shows the occurrence frequency; Figure 2a shows absolute changes and Figure 2b shows relative changes. The occurrence frequency in Figure 2a drops smoothly but quickly from about 80 events per day with amplitude changes of at least cm 2 s 1 to one event per 5 days with an amplitude change of at least cm 2 s 1. Figure 2b presents the occurrence frequency of the relative amplitude changes of the same events (the relative amplitude is defined as the absolute value of the change of the ion flux divided by the mean of the ion fluxes before and after the event). The occurrence frequency drops off quickly as the relative amplitude increases. We note that the distribution of these events is not uniform in time. These events are strongly clustered, but the clustering is not related to the speed or density of the solar wind [Riazantseva et al., 2003a]. [7] We selected the largest 207 events, those with ion flux changes larger than cm 2 s 1, for a more detailed study of their properties and characteristics. We first investigate the distributions of the amplitudes and the durations of these events. Figure 3 shows that the relative solar wind ion flux changes are most often in the range 1 4, although roughly 12% have values of more than 4. Similar results for the amplitude distribution of dynamic pressure changes of the solar wind were found using a smaller data set from Interball-1 and IMP-8 [Dalin et al., 2002a]. [8] Figure 4 shows a histogram of the durations of the flux changes for each event. These durations were obtained by analysis of the high-resolution (1 s) Interball-1 measurements. A majority, 63%, of these large flux changes take place in less than 1 min. About 21% have durations of 5 s or less. [9] Figure 5 shows two examples of events with very rapid flux changes measured with Interball s highest time resolution. Figure 5a shows a large pulse with a duration of nearly 6 min. observed by Interball-1, Wind-SWE, and Wind-3DP. The data from all three instruments are very similar. Figure 5b shows the first flux increase in Figure 5a at high (60 m s) time resolution; the ion flux increases by a factor of 2 in 1.2 s. The observations of the same front by the Wind-3DP instrument with 3 s resolution are also shown. These data agree very well in spite of the large separation between these spacecraft. The GSE coordinates of Wind were (X = 55.8, Y = 49.8, Z = 6.9) R e and of Interball-1 were (X = 14.8, Y = 2.3, Z = 14.2) R e. The 2of11

3 Figure 3. A histogram of the relative solar wind ion flux changes for the 207 events with ion flux changes > cm 2 s 1 (df = jf 2 F 1 j), where F 1 and F 2 are the fluxes before and after each event. gyroradii. The agreement between the Interball-1 and the Wind-3DP data is not as good in this case; the Wind-3DP data also show a flux increase, but the amplitude of this increase is smaller and the profile is different. This difference may be due to local inhomogeneities on the ion flux front. 5. Statistical Analysis of the Plasma Parameters and of the IMF [11] The behavior of the other solar wind plasma parameters and of the IMF during these large and sharp changes of Figure 2. Occurrence frequency of sharp and large solar wind ion flux changes, showing (a) the absolute and (b) the relative value of the amplitude. The horizontal arrows indicate that all events with changes (absolute or relative) more than the value of the abscissa are counted. (i.e., Figure 2a shows that 80 events per day have df > cm 2 s 1 ) The total number of events detected by Interball-1 from is about 37,000. fronts were very sharp and the amplitudes of the ion flux increases were the same at each location. [10] Figures 5c and 5d show an event where the ion flux change occurred in less than 1 s. Figure 5c shows a pulse observed by three instruments with a duration of about 1 min. Figure 5d expands the very sharp first flux increase. A factor of 1.5 increase in the ion flux occurred in only 0.25 s. Since the solar wind speed during this event is about 600 km/s (as measured by Wind), the 0.25 s front passage corresponds to a spatial boundary width of about 150 km. The proton thermal velocity was km/s and the IMF magnitude was about 20 nt, so the boundary thickness is only 3 5 proton Figure 4. A histogram of the time durations of the 207 events with solar wind flux changes > cm 2 s 1. 3of11

4 Figure 5. Examples of very sharp solar wind ion flux increases. (a) A sharp solar wind ion flux increase on 5 June The plot shows 1 s INTERBALL-1 data (black thin line), 3 s Wind-3DP data (gray bold line), and 90 s Wind -SWE data (black bold line). (b) The 0.06 s INTERBALL-1 data (black thin line) and 3 s Wind-3DP data (gray bold line) for the first front of this pulse. The duration of the front is 1.2 s. (c) A very sharp increase of solar wind ion flux on 24 May The plot shows 1 s INTERBALL-1 data (black thin line), 3 s Wind-3DP data (gray bold line) and 90 s Wind-SWE data (black bold line). (d) The 0.06 s INTERBALL-1 data (black thin line) and 3 s Wind-3DP data (gray line) for the first front of this pulse. The duration of the front is 0.25 s. the ion flux could help us to understand these features. In some cases (such as in Figure 1), the observed flux changes occur because the plasma density changes; the plasma speed, temperature, and IMF remain constant. More commonly, speed changes are small but changes are observed in the IMF. Figures 6, 7, and 8 use histograms to quantify our results. The lighter histograms in all three figures bin the absolute values of the speed changes jdv x j across the 207 large flux change events. In almost 60% of these events, the bulk speed changes are less than 10 km/s. Thus in a majority of the events the flux change is driven only by the density change. The speed changes by more than 40 km/s for only about 5% of the events. [12] In Figure 6, the darker histograms inside each bin of dv show the distributions of the change in thermal speed dv th, with the width of each thermal speed bin 1 4 given in the figure legend. More than half the events with the smallest dv x also have small dv th. The value of dv th becomes larger in the larger dv x bins. [13] Figure 7 shows a plot similar to that of Figure 6, except that the darker, interior bins show the absolute value of the change of the magnetic field magnitude djbj. The two bins with the smallest dv x values have the majority of djbj values in the first two bins. The distributions of djbj are broader for the higher dv x bins. In Figure 8, the dark interior histogram shows the number of events where both the theta and phi angles of the IMF change by less than 30 across the event. About 50% of the low dv x events also have small changes in the field direction; the percentage of events with large field direction changes increases for events with larger dv x.to summarize Figures 6 8, events with small changes in the solar wind speed were more likely to have small changes in the proton thermal speed, IMF strength, and IMF direction than events with larger changes in solar wind speed. [14] In many events the large flux changes are caused by large density changes with little change in the other solar wind and IMF parameters. No significant dependence of the amplitude of the solar wind ion flux change on the magnetic field magnitude or on the bulk speed was observed. These observations confirm that these events are not related to interplanetary shock waves (which should be accompanied by increasing values of all plasma parameters and the IMF magnitude), nor to Alfven waves (in which variations of the bulk speed and IMF dominate), nor to the fast magnetosonic waves (in which the plasma density and IMF magnitude vary in phase and have similar relative magnitudes). 6. Pressure Balance Across the Boundaries of These Solar Wind Structures [15] The internal pressure gradient across these events is important for understanding the evolution of these small 4of11

5 Figure 6. A histogram of solar wind bulk speed changes dv x at these events (large gray boxes). The smaller, darker boxes show histograms of thermal speed changes dv th within each dv x bin. The total number of events is 207. solar wind structures. In many solar wind structures (called pressure-balanced structures), the plasma and magnetic field change in concert to maintain equilibrium of the total (P th + P m ) pressure, where P th is the thermal pressure and P m is the magnetic pressure. Figure 9a shows that the changes in the solar wind ion flux and in the magnetic field magnitude usually have opposite signs. In 95% of the cases, when the ion flux increases the magnetic field magnitude decreases and vice versa. Figure 9b shows that changes in the thermal and magnetic pressure also usually have opposite signs. Although these changes are in the right direction to maintain pressure balance, we must check the conservation of the total pressure balance quantitatively by comparing the sum of the thermal pressure plus magnetic pressure on each side of the flux change. We use the Wind density (n), ion temperature (T i ), electron temperature (T e ) (the use of the electron temperature is necessary since the contribution of electrons to the thermal pressure is significant in most events), and IMF magnitude (jbj) to calculate the total pressure change across each event. [16] Figure 10 shows a histogram of the relative total pressure changes dp/p 1 (where P 1 is the total pressure before the solar wind ion flux change event, P 2 is the total pressure after the event, and dp = jp 2 P 1 j). The figure shows that the total pressure change across the event is less than 10% in over 55% of the events but that in about 13% of the events the total pressure changes by more than by 30%. The darker interior histograms in Figure 10 divide the dp/p 1 bins according to the value of b upstream of each event. For dp/p 1 < 10%, the number of events with b > 1 is nearly equal to the number of events with b < 1, while for dp/p 1 > 10% the number of events with b > 1 is about 3 times greater than the number of events with b < 1. Thus when the total pressure is not in balance, the thermal pressure is more likely to be larger than the magnetic pressure. Riazantseva Figure 7. A histogram of solar wind bulk speed changes dv x at these events (large gray boxes). The smaller, darker boxes show histograms of IMF magnitude changes djbj within each dv x bin. The total number of events is 207. Figure 8. A histogram of solar wind bulk speed changes dv x at these events (large gray boxes). The darker boxes show events for which dq <30 and df <30. The total number of events is of11

6 Y se = 20.8 R e,z se = 0.5 R e ) and Interball-1 (X se = 9.3 R e, Y se = 23.3 R e,z se = 9.8 R e ). Barkhatov et al. [2003] show that the total pressure inside the flux pulse is twice the total pressure on either side of the pulse. Figure 11c shows that the thermal pressure decreases by a factor of 3 inside the pulse and the magnetic pressure increases by a factor of 16. Barkhatov et al. [2003] also show that the observed evolution and that predicted by an MHD model are similar. Two other mechanisms which could explain the different pulse widths observed by Wind and Interball-1 are (1) the inclinations of the leading and trailing edges of the pulse are different and (2) the ion flux fronts are spatially inhomogeneous. The first mechanism was checked by comparing the front orientations using data from a third spacecraft; different inclinations were not observed. The second mechanism is unlikely to be important because Wind and Interball-1 are close in Y se and Z se. Another point in support of the pulse evolution hypothesis suggested above is the agreement of the observations with MHD models. [18] An interesting plasma physics problem is how these very sharp boundaries are maintained over the 200 R e separation between Wind and Interball-1 (and possibly much longer distances). To investigate this topic, we need to compare observations with very high time resolution from multiple spacecraft. This subject will be a topic of future study. 7. Inclination of the Sharp Solar Wind Fronts [19] The orientations of plasma and magnetic field structures within the solar wind were studied for middle-scale- Figure 9. (a) The solar wind ion flux change plotted versus the IMF magnitude change for each event. Squares (for increasing flux) and crosses (for decreasing flux) show events for which the sign of df is opposite to the sign of djbj. (b) The solar wind thermal pressure change (dp th ) plotted versus the magnetic pressure change (dp m ) for each event. Squares (for increasing thermal pressure) and crosses (for decreasing thermal pressure) show events for which the sign of dp th is opposite to the sign of dp m. The total number of events is 207. et al. [2005] give a detailed analysis of events which conserve (or do not conserve) pressure balance at solar wind ion flux change boundaries. Half of the events, the pressure balanced structures, may be tangential discontinuities. More detailed analysis to attempt to classify the observed discontinuities by type will be done in future work. [17] Since many events are not in pressure balance, these structures must evolve with time. Figure 11 shows an example of an evolving event; the solar wind ion flux and IMF pulse expands in width between Wind (X se = 226 R e, Figure 10. A histogram of solar wind total pressure changes dp/p 1 (P = P th +P m ) across the ion flux event boundaries. The darker interior histograms show the distribution of b (b 1 =P th1 /P m1 ) within each dp/p 1 bin (P 1 is the total pressure before the event). The dark bins separate the events into times when beta is more than, less than, or roughly equal to 1 (±20%). The total number of events is of11

7 Figure 11. Simultaneous Interball-1 and Wind observations of pulses of (a) the plasma flux and (b) the IMF magnitude. The Wind data are time-shifted with respect to the Interball-1 satellite data by 76 min (the solar wind propagation time between these two spacecraft). (c) The thermal, magnetic and total pressures. length plasma [Richardson and Paularena, 1998; Coplan et al., 2001; Dalin et al., 2002c] and magnetic structures [Collier et al., 1998; Nakagawa et al., 2002]. The inclination of sharp IMF fronts was investigated in a number of papers [see, e.g., Turner and Siscoe, 1971; Sergeev et al., 1986; Shukhtina et al., 1999]. However, the orientation of small-scale (duration of several minutes, spatial scales from several tens to several hundreds of thousands of kilometers) plasma structures such as the rapid ion flux variations discussed in this work have not previously been reported. The location of the first interaction of a sharp solar wind front with the magnetosphere (i.e., the place where a front is tangent to the bow shock and the magnetopause) may be important for space weather forecasting, as are the direction of the magnetic field and the change in the dynamic pressure. Accurate prediction of arrival times of MHD discontinuities and plasma fronts at Earth requires knowledge of the orientations of these features. [20] For middle-scale-length IMF and solar wind plasma structures, the orientation is usually determined from the lag which gives the best correlation between data from two spacecraft [see Richardson and Paularena, 1998; Coplan et al., 2001; Dalin et al., 2002c]. For MHD IMF discontinuities, the orientation of the front was determined from the magnetic field component rotation by the minimum variance method [Collier et al., 1998; Nakagawa et al., 2002; Turner and Siscoe, 1971; Sergeev et al., 1986; Shukhtina et al., 1999]. However, these methods are not valid for the sharp plasma fronts discussed here. Minimum variance analysis (MVA) methods are applicable for magnetic field discontinuities, but in many of our events the magnetic field changes are small and/or not as sharp as the flux change. The analogue to the MVA method [Sonnerup et al., 2004] for the ion flux vector can be used only if the velocity changes, and in most of the events the velocity change is small. Furthermore, this method requires high time resolution velocity data (for fronts of several minutes duration a time resolution of several seconds is needed) and the Wind data we use has a 90 s velocity resolution, so we cannot use these methods. [21] The arrival times of the same event at several spacecraft plus the solar wind velocity and spacecraft positions allow calculation of the front inclination, assuming a planar front [Riazantseva et al., 2003a; Sonnerup et al., 2004]. We show below that the planar front approximation is generally true. The large ion flux changes observed by Interball-1 were also observed by other spacecraft located in the solar wind. After the solar wind propagation time was removed, delays were observed between the event arrival times at different spacecraft. Figure 12 shows that these delays are most often less than 6 min, but some are min. We assume that the fronts do not propagate relative to the plasma. This assumption is valid for most events, since few of these events are rotational discontinuities (which propagate relative to the plasma). These delays are then used to calculate the normal to the plane of the front, i.e., the inclination of the front. [22] Figure 13 shows the calculation procedure schematically. The solar wind moves from point C to point A (from Wind to Interball-1) in time dt 1 = dx/v. For a front inclined at an angle a, the time delay is dt 2 = dyz tan(a)/v. We first consider cases with data from two spacecraft from which Figure 12. Histogram of the time delays between the time of arrival of ion flux change events at Wind and Interball-1. The total number of events is of11

8 Figure 13. A schematic diagram of front propagation from Wind to Interball-1, where C is the location of Wind, D is the location of INTERBALL-1, dx GSE is the distance between the spacecraft along the axis X GSE,dY GSE Z GSE is the distance between the spacecraft in the plane perpendicular to the X GSE axis, and V is the solar wind velocity. The line DB shows the front line when it passes Interball-1. a is the angle between front and the line AD. the one-dimensional angle of the front inclination can be determined. We only use the large flux change events for which we could determine the time delay dt 2 accurately and which had bulk velocity changes of less than 50 km/s. This selection excludes events associated with the passage of interplanetary shock waves. [23] Figure 14 gives the inclination angles of these events; Figure 14a shows the absolute values of the angles and Figure 14b shows the signed values of the angles. Since the accuracy of the delay dt 2 is approximately ±1 min, we also show the angle distributions which result from excluding fronts which have dt 2 < 2 min. The dark bins in Figure 14 show these results; exclusion of the dt 2 < 2 min events substantially reduces the number of fronts with small inclination angles (0 20 ) and has little effect on the number of fronts with large inclination angles (>20 ). The mean value of the inclination angle is about 30 because of the large fraction of events with small delays. However, this inclination exceeds 30 for half of the events. Figure 14b shows that the front inclination angles have a small asymmetry in their distribution. The number of positively oriented fronts (with an inclination in the same direction as the Parker spiral) is a slightly larger than that of the negatively oriented fronts with a maximum in the interval However, this difference is too small to make a statistically significant conclusion. Thus although the rotation of the Sun may effect the orientation of the solar wind plasma fronts, these fronts are not tilted nearly as strongly as the IMF. [24] We assumed above that the fronts are planar over distances of order several tens of R E. We check this assumption by comparing the inclination angles determined independently for the same event using two different pairs of spacecraft: Interball/Wind and IMP-8/Wind (or Geotail/ Wind). We selected the most reliable large events, with >2 min delays. We also chose only events where the plane passing through the Interball-1 and Wind spacecraft was nearly parallel to the plane X GSE Y GSE or to the plane X GSE Z GSE. We found 17 events with observations by three spacecraft which met these criteria (7 parallel to the X GSE Y GSE plane and 10 parallel to the X GSE Z GSE plane). In most cases, the angles determined from different pairs of spacecraft are fairly consistent [Riazantseva et al., 2003b]. This result confirms that the plasma flux change fronts are planar on scales of order 80 R e in the Y GSE Z GSE plane. [25] Since the fronts are generally planar, we can use observations of the same event by three or more spacecraft to determine the three-dimensional front orientation. We selected 38 events with well-determined delays observed by at least three spacecraft. Events with a large velocity change were again excluded. The front orientation is defined by two angles: the angle from the ecliptic plane X GSE Y GSE and the angle from the plane X GSE Z GSE. Figure 15 shows the front angles a XZ (between the normal to the front and the plane X GSE Z GSE, analogous to the azimuthal angle f) and a XY (between the normal and the plane X GSE Y GSE, analogous to the elevation angle q) for these 38 events. We use these angles as they are easy to interpret; the front perpendicular to the ecliptic plane corresponds to the angle a XY = 0. The estimated angles have a large scatter. The front inclinations vary from 0 to almost 90. However, the largest number of points is in the second quadrant (the angle a XZ is positive and the angle a XY is negative), consistent with a large role being played by the solar rotation. Approximately 30% of the inclinations fall within the square (shown by the dashed square in Figure 15). The inclination angles in the two planes, seen in Figure 15, do not seem independent. The fronts strongly inclined to the ecliptic plane are usually strongly inclined to the other plane as well. [26] We tested the accuracy of these estimates by comparing the inclinations of the sharp plasma and magnetic field fronts in cases when these parameters changed simultaneously. The plasma inclinations were determined by the time delay method; the magnetic field inclinations were determined by both the time delay method and the minimum variance method. These independent methods provide a check on the time delay method for the plasma fronts because the inclination of the plasma and magnetic fields front must be the same. We found that the plasma and IMF front orientations are the same to within an accuracy of about 10 [Riazantseva et al., 2003c]. Figure 14. Histogram of one-dimensional inclination angles of the front plane determined from observations of the front by two spacecraft for (a) the absolute value of the angle and (b) for the signed angle. 8of11

9 8. Geoeffectiveness of Sharp and Large Solar Wind Dynamic Pressure Changes [28] Sharp changes of the solar wind plasma and IMF can produce a geomagnetic response. Sibeck et al. [1996] show that large, sharp pulses of the solar wind dynamic pressure result in rapid changes of the Chapman-Ferraro currents on the magnetopause. These current changes lead to strong, short-term disturbances of the magnetic field at geosynchronous orbit in the dayside magnetosphere. Figure 17 shows an example of an event from Riazantseva et al. [2003a]. The dynamic pressure pulse is shown in Figure 17a and is observed by both Interball-1 and Wind. The bulk speed and the IMF do not change during this pulse (Figures 17b and 17c). Thus the geomagnetic disturbances caused by this event result only from the dynamic pressure variation, not from a change in the IMF. Figure 17 shows two different kinds of magnetospheric response to this pressure pulse: Figure 17d shows that magnetospheric field variations at geosynchronous orbit and at a ground-based dayside nearequatorial station are very similar to the pressure change profile; Figure 17e shows that large-amplitude (70 nt) geomagnetic pulsations in the PC 5 frequency range (T 400 s) are excited in the auroral region in apparent association with the solar wind pressure changes. Similar results are observed in several other cases, suggesting this response is typical [Parkhomov et al., 2005; Borodkova et al., 2005]. Figure 15. Two-dimensional orientation of solar wind plasma fronts with respect to the X GSE Y GSE (analagous to the elevation angle q) and X GSE Z GSE (analagous to the azimuthal angle f) planes determined from observations of the front arrival times at three spacecraft. The typical error bar is Conclusions [29] 1. Small-scale solar wind structures often have very sharp boundaries; we have studied the characteristics of [27] Figure 16 shows a comparison of data from five spacecraft: Figures 16a and 16b present plasma fronts (solar wind ion flux changes) and Figures 16c and 16d present magnetic field magnitude fronts, all observed by five spacecraft, Interball-1 (X GSE =16R e,y GSE = 0.1 R e,z GSE = 12.3 R e ), Wind (X GSE = R e, Y GSE = 35.9 R e, Z GSE =26R e ), ACE (X GSE = R e,y GSE = 6.4 R e, Z GSE = 22.8 R e ), Geotail (X GSE = 17.8 R e, Y GSE = 0.8 R e,z GSE = 2.2 R e ), and IMP-8 (X GSE =11.4R e, Y GSE = 28.4 R e,z GSE = 25.3 R e ). All data are time-shifted to the Interball-1 position using the solar wind propagation time. During this event, the solar wind ion flux increases at the first front and decreases at the second front. The magnitude of the IMF, on the contrary, decreases at the first front and increases at the second. Note that the magnetic field Y-axes in Figures 16c and 16d decrease upward. Figures 16e and 16f compare the normal angles to the plasma and IMF fronts. The difference between these normals is small, about 4 for the first front and 6 for the second, less than the roughly 10 average error for this method [Riazantseva et al., 2003c]. So the time delay method does provide a reasonable measure of the sharp plasma front inclinations. Figure 16. A pulse with sharp fronts in both (a and b) the plasma and (c and d) IMF observed by five spacecraft. Vertical dashed lines mark the fronts. (e and f) A comparison of the normals to the plasma and IMF front planes. 9of11

10 by more than 30%. When pressure is not in balance, the events should evolve and evidence of this evolution is observed. [34] 6. The planes of the sharp and large plasma fronts are often significantly inclined to the Sun-Earth line. The distribution of the inclinations is broad: roughly equal numbers have angles between the plane of the front and the Sun-Earth line of greater than and less than 60. [35] 7. Large and sharp changes of the solar wind ion flux (or dynamic pressure) can be geoeffective. These changes cause disturbances of the magnetospheric and ground-based geomagnetic field (such as compression and excitation of oscillations) even if the bulk velocity and IMF are constant. [36] 8. The sharp ion flux fronts could be remnants of sharp changes in ion flux at Sun and which are maintained as the solar wind moves to Earth. Riazantseva [2003a] suggested that these events are a manifestation of the fine structure of solar coronal streamer and that they should be observed mainly in the heliospheric current sheet. Observations only partly agreed with this hypothesis. The sharp ion flux changes could also be created as the solar wind propagates from the Sun to Earth, but a plasma instability which could produce the sharp solar wind ion flux fronts has not been identified. This important topic needs more investigation. Figure 17. The geomagnetic response to a change in the dynamic pressure of the solar wind plasma on 11 April The panels show (a) the solar wind pressure, (b) V x, (c) the IMF magnitude and components, (d) the magnitude jbj and component X of the magnetic field at the GOES-8 geosynchronous satellite and the low-latitude station Tamanrasset, and (e) the magnetic component H at the high-latitude station Narsarsuaq. events in which the ion flux changes by more than 20% in less than 10 min. [30] 2. Most of these ion flux change events have small velocity changes (they are mainly density variations) and in some cases IMF changes are also small. [31] 3. The average occurrence frequency of these events varies as a function of their amplitude from approximately 80 events per day for ion flux changes > cm 2 s 1 to one event every 5 days for changes > cm 2 s 1. [32] 4. Large solar wind ion flux changes are often very rapid; the change occurs in less than 1 min in 63% of the cases and in less than 10 s in 33% of the cases. Thus the boundaries of these solar wind structures can be as thin as tens of gyroradii. [33] 5. In 95% of the events, the changes of the ion flux and the IMF magnitude and the changes of the thermal pressure and the magnetic pressure are opposite in sign so that the pressure changes at least partially compensate. 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