Bulk properties of the slow and fast solar wind and interplanetary coronal mass ejections measured by Ulysses: Three polar orbits of observations

Transcription

1 Click Here for Full Article JOURNAL OF GEOPHYSICAL RESEARCH, VOL. 114,, doi: /2008ja013631, 2009 Bulk properties of the slow and fast solar wind and interplanetary coronal mass ejections measured by Ulysses: Three polar orbits of observations R. W. Ebert, 1,2 D. J. McComas, 2,1 H. A. Elliott, 2 R. J. Forsyth, 3 and J. T. Gosling 4,2 Received 28 July 2008; revised 7 October 2008; accepted 15 October 2008; published 30 January [1] We examined plasma and magnetic field observations from all three Ulysses polar orbits of the Sun to study the properties of the slow and fast solar wind and interplanetary coronal mass ejections (ICMEs). We derived equations to characterize the radial and latitudinal variations for these three types of heliospheric plasma and identify distinguishing features in their spatial variations. Most notably, the slow-wind proton temperature falls less rapidly with distance than does the fast wind, indicating a source of enhanced heating in the low-speed wind. After removing the radial variations from the measurements, only minor latitudinal gradients were identified. The fast wind has now been shown to be only weakly dependent on solar latitude for two successive solar minima. The spatial variations in the ICME properties do not differ significantly from the slow and fast solar wind, although the variability in their parameters is much larger. We also investigated solar cycle variations in the fast polar coronal hole (PCH) flows by comparing their properties measured over Ulysses 1st and 3rd orbits. While the latitudinal gradients were similar, slight differences were observed in the radial dependence for the proton density and magnetic field strength. Also, a slight reduction in the proton speed at 1 AU, along with more significant decreases in the proton temperature, density, dynamic pressure, and magnetic field strength, was observed for the 3rd orbit relative to that for the 1st. These results are consistent with recent observations of weaker PCH flows for the current solar minimum. Citation: Ebert, R. W., D. J. McComas, H. A. Elliott, R. J. Forsyth, and J. T. Gosling (2009), Bulk properties of the slow and fast solar wind and interplanetary coronal mass ejections measured by Ulysses: Three polar orbits of observations, J. Geophys. Res., 114,, doi: /2008ja Introduction [2] The solar wind continuously expands from the solar atmosphere out into interplanetary space. Parker [1958] proposed that the solar wind resulted from the supersonic expansion of the Sun s 10 6 K corona. Direct, in-situ, observations of the solar wind were first made in 1959 by the Soviet spacecraft Luna-1 while measurements from the solar plasma experiment on the Mariner 2 spacecraft revealed a continuous and highly variable stream of ionized particles from the Sun [Neugebauer and Snyder, 1962]. [3] Coronal mass ejections (CMEs) are a transient type of solar wind which originate from closed magnetic field regions in the solar corona and are ejected from the solar 1 Department of Physics and Astronomy, University of Texas at San Antonio, San Antonio, Texas, USA. 2 Space Science and Engineering Division, Southwest Research Institute, San Antonio, Texas, USA. 3 Space and Atmospheric Physics, Imperial College London, London, UK. 4 Laboratory for Atmospheric and Space Physics, University of Colorado, Boulder, Colorado, USA. Copyright 2009 by the American Geophysical Union /09/2008JA013631$09.00 atmosphere, transporting large quantities of plasma ( g) and magnetic flux (10 15 Wb) into interplanetary space [Gosling, 1996; Hundhausen, 1999]. They were first observed in the 1970s when events were detected in images from coronagraphs onboard OSO-7 [Tousey et al., 1974] and Skylab [Gosling et al., 1974]. [4] The interplanetary extensions of CMEs are commonly referred to as interplanetary coronal mass ejections (ICMEs). Nearly a third of all ICMEs are observed as magnetic flux rope structures [Gosling, 1990], which are called magnetic clouds when characterized by enhanced magnetic field strength, low plasma beta and smooth rotation in a component of the magnetic field [Klein and Burlaga, 1982; Gosling, 1996]. ICMEs also drive shocks in events where their bulk speed relative to the ambient wind ahead is greater than the fast magnetosonic speed of the upstream solar wind [Sheeley et al., 1985]. Gosling et al. [1994] identified a class of ICMEs observed at high latitudes, which are bounded by forward-reverse shock pairs driven not by the bulk speed of the ICME but by overexpansion caused by their high internal pressures. ICMEs are detected in the solar wind via signatures in their plasma, magnetic field and compositional properties. However, the detection of ICMEs in solar wind measurements remains a 1of17

2 difficult and rather subjective task because of the variability in their properties and signatures [Gosling, 1996]. A comprehensive review of ICME signatures in the solar wind can be found in the work of Gosling [1990], Neugebauer and Goldstein [1997], and more recently Zurbuchen and Richardson [2006]. [5] Over the past five decades, a number of spacecraft have been launched with instruments specifically designed to make in-situ measurements of the plasmas and magnetic fields that permeate interplanetary space. In this study, we focus on observations made by instruments onboard the Ulysses spacecraft. Ulysses is currently in its third 6.2-year orbit of the Sun on a trajectory that ranges in heliocentric distance from 1.3 AU to 5.4 AU and heliolatitudes between 80.2 S to 80.2 N. As the first, and only, spacecraft with a nearly polar orbit of the Sun, Ulysses has a unique vantage point from which to study the three-dimensional (3-D) heliospheric environment. Measurements from the first Ulysses orbit, taken through the declining phase of the solar cycle and around solar minimum, revealed a bimodal solar wind structure consisting of a fast, uniform wind at high latitudes and slower, more variable wind at lower latitudes [McComas et al., 1998a]. These observations established the fast wind as the dominant solar wind observed at high latitudes during solar minimum, a fast stream originating from regions of open magnetic field lines on the Sun called coronal holes [Kreiger et al., 1973]. The observed bimodal structure also supported the findings of Schwenn et al. [1978] that the fast streams from high latitudes were separated from the slower wind at low latitudes by relatively sharp boundaries. The solar wind flow detected during Ulysses second orbit was highly variable, with slow and fast wind observed at all heliolatitudes. This orbit coincided with a period of rising and enhanced solar activity, including solar maximum [McComas et al., 2003]. The origin of the slow solar wind is still under consideration with observations suggesting coronal streamers [Gosling et al., 1981], coronal hole boundaries [Neugebauer et al., 1998], and active regions [Neugebauer et al., 2002] as potential sources. [6] McComas et al. [2000] conducted an extensive statistical survey of the radial and latitudinal variations of the solar wind plasma and magnetic field properties for the midlatitude to high-latitude coronal hole flows measured over the first Ulysses orbit. The radial gradients were fit to power laws and showed that the proton temperature falls with distance as R 1.02, slower than would be expected for adiabatic expansion. Also, the radial and tangential components of the magnetic field dropped off more slowly (R 1.77 ) and rapidly (R 1.16 ), respectively, than theoretically predicted in a structureless solar wind. This slower drop off is likely due to the presence of large amplitude Alfvén waves in the high-latitude regions [Smith et al., 1995; Forsyth et al., 2002]. After correcting for the radial gradients, only small latitudinal variations were observed. For example, a latitudinal gradient in proton speed (V p )of1kms 1 deg. 1, in proton density (N p )of 0.02 cm 3 deg. 1, and in proton temperature (T p ) of 223 K deg. 1. Of these, McComas et al. [2000] only considered the velocity gradient to be statistically significant. The alpha/proton ratio (N a /N p ) was shown to be extremely constant in the fast coronal hole wind with a value of 4.4%, consistent with previous observations [see, for example, Bame et al., 1977]. Several other studies have examined the radial evolution in the properties of the solar wind [Gazis and Lazarus, 1982; Richardson et al., 1995], ICMEs [Liu et al., 2005] or both [Wang and Richardson, 2004; Liu et al., 2006]. These studies use data from a variety of spacecraft (Helios, Wind, ACE, Ulysses, Pioneer, and Voyager) located throughout the heliosphere. With the exception of Ulysses, these spacecraft were located predominantly near the ecliptic plane, where sampling of slow solar wind flows is more common. Not surprisingly, the radial evolution in some of the solar wind parameters found by these studies were not in agreement with those reported by McComas et al. [2000] for purely fast wind, owing to the much more complicated interactions between parcels of wind with different speeds generally observed at lower latitudes. A clear example is the proton temperature with estimates ranging from R 0.7 [Gazis and Lazarus, 1982] to R 0.62 [Wang and Richardson, 2004] out to 10 and 20 AU, respectively, and R 0.49±0.01 out to 40 AU [Richardson et al., 1995]. In each case, the decrease in T p deviates from adiabatic expansion and may be a result of heating caused by converting energy from the bulk flow into thermal energy by shocks or stream interactions [Gazis and Lazarus, 1982]. These differences provide further motivation to compare the radial and latitudinal variations in the properties of different types of heliospheric plasma. [7] In this paper we compare the radial and latitudinal variations of the bulk plasma and magnetic field properties for three distinct types of solar wind plasma: the slow solar wind, the fast solar wind and ICMEs. We use data from the Solar Wind Observations Over the Poles of the Sun (SWOOPS) [Bame et al., 1992] plasma instrument and the Vector Helium and Flux Gate magnetometer (MAG) experiment [Balogh et al., 1992] onboard the Ulysses spacecraft. Here we analyze the measurements from all three orbits. We also investigate solar cycle variations in the properties of the solar wind by comparing the properties and spatial variations of the fast polar coronal hole (PCH) flows measured over the 1st and 3rd orbits, which occurred at similar periods in the solar cycle. [8] Figure 1 shows an overview of the observations. Figure 1a displays the 1-hour averaged solar wind proton speed versus time from 18 February 1992 through 5 April This plot clearly shows the bimodal structure of the solar wind during the first orbit and the variable structure of solar wind for the second orbit. During the third orbit we see a return to the bimodal structure observed during the first orbit. We define all data above 650 km s 1 not associated with an ICME as fast wind and all data below 500 km s 1 not associated with an ICME as slow solar wind. The correlation between the frequency of ICMEs and solar activity is clearly evident. The two prolonged intervals (DOY 36, 1995 DOY 288, 1996 and DOY 112, 2006 DOY 184, 2007) where no ICMEs were identified occurred when Ulysses was at high latitudes and the solar activity was near a minimum. [9] In section 2 we provide a brief description of the instruments used and the physical properties explored in this report. The methodology used to cull and analyze the data is also described there. In section 3 we present the ICME events identified in the Ulysses data from 18 February 1992 through 5 April We describe the radial and latitudinal 2of17

3 Figure 1. (a) One-hour averaged Ulysses-SWOOPS solar wind proton speed measurements versus time covering the period from 18 February 1992 through 5 April Superimposed (red circles) over the speed profile are the mean speeds for each ICME event identified at Ulysses. The blue and purple dashed lines indicate the proton speeds of 500 and 650 km s 1, respectively, which we use to identify the fast and slow solar wind; we define all data above 650 km s 1 not associated with an ICME as fast wind and all data below 500 km s 1 not associated with an ICME as slow solar wind. (b and c) Time profiles for Ulysses heliolatitude and the solar sunspot number. variations in the properties of the slow and fast solar wind and ICMEs in section 4. A summary of the solar cycle variations in the fast PCH wind measured during Ulysses 1st and 3rd orbits is given in section 5. Finally, we discuss and summarize our findings in sections 6 and 7, respectively. 2. Data Analysis [10] We investigate the properties of the interplanetary plasma using 1-hour averaged measurements of V p, alpha particle speed (V a ), T p, N p, alpha particle density (N a ), and N a /N p, obtained from the SWOOPS ion instrument. Pitch angle distribution measurements of ev electrons from the SWOOPS electron instrument were used to help identify ICMEs. We also use hourly averages of the RTN magnetic field vector components (B R, B T, B N ) and the field magnitude (jbj) measured by the Ulysses MAG experiment [Balogh et al., 1992] to investigate the properties of the interplanetary magnetic field (IMF) out to 5.4 AU. [11] The extended (>17-year) timescale of the Ulysses mission provides an immense data set from which to study the heliospheric environment. This study covers the timeframe from 18 February 1992 through 5 April 2008 (29 February 2008 for MAG data), corresponding to hour averaged measurements of the interplanetary plasma. The start time selected for this study corresponds to Ulysses exit from the Jovian magnetosphere and the beginning of its polar orbit of the Sun. Using these measurements we created three distinct data sets: one consisting solely of low-speed (V p < 500 km s 1 ) solar wind measurements, one solely of high-speed (V p > 650 km s 1 ) solar wind measurements and a third set identified as corresponding to ICME events. These subsets were created such that no measurement appeared in multiple data sets. [12] A large portion of the ICME events were obtained directly from the existing Ulysses-SWOOPS ICME list ( which was compiled by Jack Gosling with assistance from Dan Reisenfeld and Robert Forsyth and covers the period from February 1992 through December As part of this study, we extended this list through 5 April A description of this extended ICME list and the techniques used to identify the ICMEs are provided in the next section. [13] The data sets for the slow and fast solar wind were created by first removing all measurements associated with the ICME events. We also removed all intervals within 24 hours on either side of each ICME interval. This was done to eliminate sheath regions that may surround the ICME. These subsets were constructed with the goal of creating a solar wind data set devoid of ICMEs and their local disturbances. We also eliminated any solar wind measurements with a proton speed ranging from 500 < V p < 3of17

4 650 km s 1 in an attempt to remove most of the data associated with stream interaction regions (SIRs). The measurements associated with SIRs in the fast wind data set, which were generally observed as well-formed co-rotating interaction regions (CIRs), were removed from our analysis by only selecting measurements taken above latitudes in which CIRs were observed. The SIRs in the slow-wind data could not be easily removed and thus were included in the analysis of the slow-wind measurements. Finally, the interval associated with the recent encounter with comet McNaught (3:30 DOY 35, 2007 to 8:30 DOY 39, 2007) [Neugebauer et al., 2007] was also removed. [14] Our resulting ICME, slow solar wind and fast solar wind data sets consist of 10996, and hour averaged measurements, respectively, corresponding to 7.9%, 35.6%, and 37.0% of the measurements in the Ulysses data set. The 500 < V p < 650 km/s solar wind data, sheath intervals and anomalous event data accounted for the remaining 19.5% of the Ulysses measurements. 3. Ulysses ICME List [15] The identification of ICMEs from in-situ solar wind measurements is not a straightforward process [Gosling, 1996; Zurbuchen and Richardson, 2006]. There is no single feature exhibited by all ICMEs and no standard signature that can be used to identify them all. In this study, we use a combination of signatures to find ICMEs in the SWOOPS and MAG data. These signatures include bi-directional electrons [Montgomery et al., 1974; Gosling et al., 1987], lower than expected proton temperature (T p /T ex < 0.5) [Gosling et al., 1973], enhanced alpha/proton ratio (N a /N p > 0.08) [Borrini et al., 1982], enhanced magnetic field strength, low plasma beta (b < 0.1), and smooth rotation in the magnetic field vector [Burlaga et al., 1981]. The expected proton temperature (T ex ) was calculated using the T p V relation from the Genesis mission [Neugebauer et al., 2003], which was derived from solar wind speed and temperature measurements from both the Genesis spacecraft and the SWEPAM instrument [McComas et al., 1998b] on the Advanced Composition Explorer (ACE) spacecraft. These equations are T ex ¼ þ 595:2V p 0:1623V 2 p 6697 =1:214 for V p < 450 km s 1 T ex ¼ þ 1217V p 0:5214V 2 p 6697 =1:214 for V p > 450 km s 1 [16] We found that these equations did a reasonably good job of describing the Ulysses proton temperature measurements when normalized to 1 AU (T p R). [17] All ICME events were identified by eye using the signatures outlined above. In general, we required at least two signatures within a given solar wind interval before we identified it as an ICME. It should be noted that owing to power limitations on the Ulysses spacecraft, the electron pitch angle distribution data used for identifying bi-directional electrons was unavailable from October 2004 through 25 March 2006 and available only intermittently after then. [18] The Ulysses ICME list presented here covers a period that coincides with the declining phase of solar cycle 22 and all of solar cycle 23. The updated list is provided in Appendix A. To this list we have added information on the spacecraft coordinates and the average proton speed and magnetic field strength along with their standard deviations for each ICME event. The spacecraft coordinates are given using the heliographic inertial (HGI) coordinate system. The heliolatitude corresponds to the latitude of the spacecraft at the beginning of the event while the heliolongitude is the value given for the day in which the event was identified. In total, 178 ICME events have been identified in the Ulysses plasma and magnetic field data. The 27 events reported after December 2002 are the ICME events newly identified in this study. We note that because of the nature of the Ulysses orbit, a nearly half (49%) of these 178 events were observed at R > 5.0 AU. The maximum number of events in a single year (39) occurred in 1999, coinciding with the ascending phase of solar cycle 23, whereas only a single event was identified in the years 1995, 2006 and The latter were periods when Ulysses was embedded in the fast PCH flows, near solar minimum. Given that events were identified by eye, some smaller or less clear events may have been excluded from our survey. [19] Figure 2 shows an ICME event identified in this study. The proton temperature and density have been normalized to 1 AU by multiplying the measurements by R and R 2, respectively. The dominant ICME signatures in this event are bi-directional electrons, low b p and smooth rotation in the polar angle (d) of the magnetic field. These signatures are consistent with a magnetic cloud type event. The start and stop boundaries of the event were defined by the period of smooth rotation in the magnetic field polar angle. [20] Figure 3 shows the distribution of mean proton speeds, hv p i, for the ICME events measured by Ulysses. We observed ICMEs ranging in mean speed from km s 1 with a large number of the events having a mean speed below 500 km s 1. The largest number of ICME events (31) was in the km s 1 bin. The median hv p i was 441 km s 1 while the mean was 471 km s 1. Only one event had hv p i below 300 km s 1 while only two events had hv p i greater than 800 km s Radial and Latitudinal Variations [21] In this section, we examine the radial and latitudinal variations for various properties of the interplanetary plasma and magnetic field from AU and ±80.2 in heliolatitude. These spatial gradients are characterized using a simple equation, which consists of a power law for the radial variations and a linear fit for the latitudinal variations. Specifically, we quantify the spatial variations in the properties of the three types of heliospheric plasma mentioned above by performing a Levenberg-Marquardt least squares fit ( to the following model function, Y ¼ R g ½a þ bqš: Here, Y is the solar wind parameter, R and q are the heliocentric distance and heliolatitude of the Ulysses spacecraft, respectively, and g, a and b are the free parameters in 4of17

5 the function. For the fast and slow solar wind the model function was fit to the individual 1-hour measurements, while for the ICMEs, the mean value for each event was used. To verify this method, we compared our results for the properties of the fast PCH wind measured over Ulysses 1st polar orbit with those obtained by McComas et al. [2000] for the same period. While the analysis techniques are somewhat different, our results were nearly identical Fast and Slow Solar Wind [22] Figure 4 shows the 1-hour averaged proton temperature measurements (black circles) versus heliolatitude for (Figure 4a) fast and (Figure 4b) slow solar wind. In both cases, we scaled the individual measurements to 1 AU by multiplying them by the power law component of the function used to express their spatial variations (see upper right corner in Figures 4a and 4b, respectively). Also, as a means of summarizing the observations, we separate the data into 2 heliolatitude bins and calculate the mean and variability within each bin. The mean values for each bin are then plotted over the individual measurements, shown as purple circles in the fast wind plot and as blue circles in the plot of the slow wind. The technique used to express the variability is similar to that used by McComas et al. [2000] in that we specify the value at the 5, 25, 75 and 95% level within each bin. This is shown graphically using a vertical bar through each bin mean, with the 5 and 95% values designated by the bottom and top of the vertical line, and the 25% and 75% values by the lower and upper horizontal bars, respectively. For clarity, the individual 1-hour data points have been omitted from the remaining plots, although the equations for the spatial variations in the solar wind properties were derived using these data. [23] The equation describing the spatial variations in proton temperature measured in the fast solar wind shown in Figure 4a (upper right corner) is for data above ±36, the approximate latitude at which co-rotating interaction regions are no longer observed. We see that T p in the fast Figure 2. Example of an ICME identified in the Ulysses data for the time interval of 17:30 DOY 43 to 11:30 DOY 45, The spacecraft was located at a heliocentric distance of 4.6 AU and heliolatitude of 21 N. Plotted from top to bottom, the parameters shown are the pitch angle distribution of ev electrons, V p, N p (R/R o ) 2 where R o = 1 AU, proton beta (b p ), T p /T ex (R/R o ), N a /N p, the magnetic field strength (jbj), and the angles d and l describing the magnetic field direction with respect to the RTN coordinate system. d is the polar angle out of the R T plane which ranges from 90 to 90, while l is the azimuthal angle in the R T plane and ranges from 180 to 180 (0 being the antisunward direction). The horizontal blue lines in the T p /T ex, b p, and N a /N p panels represent the ICME signature threshold for these parameters used to identify ICMEs (0.5, 0.1, and 0.08 for T p /T ex, b p, and N a /N p, respectively) in this study. Values above (below) these thresholds for N a /N p (T ex /T p and b p ) may indicate the presence of an ICME. The vertical red lines in each figure represent the event boundaries. In general, these boundaries are defined as the start and stop time for the most dominant signature in the interval. Figure 3. A bar plot representing the distribution in mean ICME speeds for the ICME events detected at Ulysses over the course of the mission. The mean speeds were separated into 25 km s 1 bin starting at 250 km s 1. 5of17

6 Figure 4. One-hour averaged proton temperature measurements (black circles) versus heliolatitude for the (a) fast and (b) slow solar wind. The data were scaled to 1 AU by removing the radial variations in the observations. We summarized the measurements by separating the data into 2 latitude bins from 0 to 80.2, with the purple and blue circles representing the bin means for the fast and slow wind, respectively. The vertical lines through these circles are used to represent the variability in the data, with the top of the line representing the value below which lie 95% of the data, the upper horizontal line the 75% level, the lower horizontal line the 25% level, and the bottom of the line the 5% level (see bottom right corner for schematic). The red (orange) line through the data in Figure 4a (Figure 4b) indicates the fit to the individual measurements from 36 to 80 for the fast wind (0 to 80 for the slow wind), with the equation to this fit shown in the upper right portion of the plot. Not shown here are the uncertainties in the parameters derived for these equations. The fits to the equations, which include the uncertainties in the parameters, are listed in Table 1. wind falls with distance approximately as R 1, similar to the result found by McComas et al. [2000] over Ulysses first orbit and somewhat below the R 4/3 predicted for adiabatic expansion. The fast wind proton temperature has a latitudinal gradient of 191 K deg 1 (0.08% deg 1 ) and a value when normalized to 1 AU and 60 latitude of K. The proton temperature in the low-speed wind has a R 0.7 radial dependence and a slight decrease with latitude of 258 K deg 1 ( 0.3% deg. 1 ). The value for the slowwind T p when normalized to 1 AU and 0 latitude is K. Compared to the fast wind, the slow wind measured by Ulysses is much more variable, cooler and has a more gradual decline in temperature with distance. That the slow wind is cooler and more variable has been discussed in a number of studies [Bame et al., 1977; see review by Schwenn, 2006]. The radial dependence for the slow-wind proton temperature is consistent with Voyager observations out to 10 AU [Gazis and Lazarus, 1982]. [24] Figure 5 shows plots of V p, N p, the momentum flux (r i V i 2 ) and jbj versus heliolatitude for the fast solar wind. The momentum flux (or dynamic pressure) is given by m p (N p V p 2 +4N a V a 2 ) where m p is the proton mass. The radial gradients have been removed from the data using the same approach as described for the proton temperature measurements. As expected, the proton speed shows no substantial radial gradient (R ) and we continue to observe the 1 kms 1 (0.93 km s 1 ) latitudinal variation previously identified by McComas et al. [2000, 2003]. The radial dependence for the proton density is R 1.86, not quite the R 2 profile expected for a radially expanding, structureless solar wind. N p has a latitudinal variation of cm 3 ( 0.4% deg. 1 ) and a value when scaled to 1 AU (q = 60 ) of 2.12 cm 3. The radial dependence for the dynamic pressure is identical to that for the proton density. The latitudinal variation is npa ( 0.2% deg. 1 ) with a value at 1 AU of 2.36 npa. Finally, we consider the magnetic field strength, jbj. The radial dependence for jbj is R 1.40 with only a slight 0.3% deg. 1 variation in latitude and a magnitude of 4.05 nt at 1 AU. After removing the radial gradients we see only slight (<1% deg. 1 ) variations with latitude in the properties shown here, consistent with fast wind observations from the first orbit [McComas et al., 2000]. Also, above ±36 heliolatitude the variability in the parameters remains consistent and relatively small, generally less than a factor of two between the 25% and 75% levels, indicating that the plasma and magnetic field properties of the fast PCH flows measured by Ulysses continue to be steady and uniform. These plots also show that our model function does a good job of describing the trends in the data. [25] Figure 6 displays plots of N p, N a /N p, jb R j and jbj versus heliolatitude for the slow solar wind. Compared to the fast wind, the properties of the slow solar wind shown here are much more variable, with differences as large as a factor of 5 between the 25% and 75% levels. We calculate a radial dependence of R 1.93 for the proton density, within 5% of the R 2 variation expected for a radially expanding solar wind. The latitudinal gradient is cm 3 deg. 1 (0.6% deg. 1 ) with a value when scaled to 1 AU of 5.55 cm 3. This calculated 1 AU value is consistent with other observations of N p near 1 AU [see, for example, Hundhausen, 1968] and is a factor of 2.6 larger than the 1 AU (q =60 ) 6of17

7 Figure 5. Similar to Figure 4 but plotted in clockwise fashion are the scaled proton speed (V p ), proton density (N p ), magnetic field strength (jbj), and dynamic pressure (r i V 2 i ) observations versus heliolatitude for the fast solar wind. The radial gradients in the measurements have been removed (see text). The equations describing the radial and latitudinal variations in the observations were derived for measurements above ±36 in heliolatitude. proton density for the fast PCH wind. The normalized to 1 AU alpha particle to proton ratio is nearly a factor of 2 less than in the fast wind. jb R j and jbj have radial dependences of R 1.43 and R 1.22, respectively. That the radial dependence for jbj is closer to R 1 in the slow wind compared to the fast indicates that the tangential component of the magnetic field, jb T j, is more dominant in the low-speed wind at these distances. This is because of the tighter Parker spiral configuration of the IMF caused by the slower speed of the slow wind. Finally, in all four plots there is more scatter in the mean values and more variability in the data compared to the fast wind. In spite of this scatter and variability, the latitudinal gradients in the properties shown here are still relatively small (<1% deg. 1 ). [26] Table 1 summarizes the results of all the plasma and magnetic field properties (listed in column 1) considered for the low- and high-speed wind. In addition to the properties already described, we include the mass flux (r i V i ), proton beta (b p ), proton thermal pressure (n p kt p ), and magnetic pressure (jbj 2 /2m o ). The mass flux was calculated using the equation m p (N p V p +4N a V a ) while the proton beta, b p, was calculated by b p = m o N p kt p /jbj 2 where m o is the magnetic permeability in free space and k is the Boltzmann constant. Columns 2 4 display the fit parameters for the equations that characterize the properties of the fast solar wind. Column 5 displays the values of the fast wind properties normalized to 1 AU (q =60 ). Columns 6 8 show the fit parameters for the slow wind. Column 9 shows the values of the slow wind normalized to 1 AU (q =0 ). We also present uncertainties in the parameters used to characterize the spatial variations in the heliospheric plasma. These uncertainties are for the fit parameters derived from our model function and do not represent the variations in the properties themselves. On the basis of the values shown in Columns 2 4 and 6 8 we see that these uncertainties are generally very small. [27] A comparison of the proton and alpha particle speeds in Table 1 indicate that at 1 AU the alpha particle speed is slightly higher than the proton speed in both the slow and fast wind. The fast wind observations support the findings by Asbridge et al. [1976] and Marsch et al. [1981] that the alpha particles in high-speed streams are faster than the protons by up to the Alfvén speed (V A = jbj/(m i n i m o ) 0.5 ). At 1 AU and 60 latitude the Alfvén speed in the fast wind is 55 km s 1 while the difference between V p and V a is 24 km s 1. [28] The alpha particle density has nearly the same radial dependence as the proton density with only minor latitudinal variations. Compared to the large enhancement in the proton density at 1 AU for the slow wind relative to the fast, only a factor of 1.2 enhancement is observed for N a. The fact that N a is not as enhanced as N p in the slow wind is reflected by the 1 AU magnitude of the slow-wind alpha/proton ratio being roughly half that of the fast wind. [29] The mass flux is calculated from the proton and alpha particle properties. The fast wind radial dependence 7of17

8 Figure 6. Similar to Figure 4 but plotted in clockwise order are observations scaled to 1 AU of the proton density (N p ), the alpha/proton ratio (N a /N p ), magnetic field strength (jbj), and radial component of the magnetic field (jb R j) versus heliolatiude for the slow solar wind. Measurements at all latitudes in which the slow solar wind was observed were used to derive the equations which characterize the spatial variations in these selected properties. for the mass flux is identical to that of the proton and alpha particle densities with a latitudinal variation of 0.2% deg. 1 and a magnitude at 1 AU of kg m 2 s 1. The slow-wind radial dependence differs slightly from the density gradient owing to the slight positive radial gradient in the proton and alpha speeds. The variation with latitude was 0.5% deg. 1 with a magnitude at 1 AU of kg m 2 s 1. [30] For the magnetic field properties, the radial dependence of jb R j and jb T j in both the fast and slow solar wind differ from the theoretical prediction of B R R 2 and B T R 1 for a structureless and radially expanding solar wind plasma. In the case of the fast wind, this deviation is attributed to the presence of Alfvén waves at high latitudes [Smith et al., 1995; McComas et al., 2000]. Also, jb R j falls more rapidly with distance in the fast wind compared to the slow wind which causes jbj to fall more rapidly as well. The latitudinal variations in the magnetic field properties are <1% deg. 1. Another interesting feature is that in the fast wind the 1 AU value for jb R j is larger than jb T j (2.37 nt compared to 1.33 nt) while in the low-speed wind we observe the opposite with jb T j = 3.13 nt and jb R j =2.46nT. Again, this is attributed to the tighter spiral structure of the IMF in the slow wind. [31] The proton thermal and magnetic pressures in the fast wind both fall as R 3. For the thermal pressure this is not surprising since the proton density and temperature fall as R 2 and R 1, respectively. The gradient in the magnetic pressure is due to the nearly R 1.5 variation in jbj. In the low-speed wind, the thermal pressure varies as R 2.4 owing to the fact that the proton temperature falls as R 0.7. The magnetic pressure falls as R 2.3 which is due to the different radial dependence for jbj (R 1.22 ). When combined, the total pressure exerted by these properties at 1 AU are 15.8 ppa and 19.8 ppa in the fast and slow wind, respectively, roughly 2 orders of magnitude less than the solar wind dynamic pressure at 1 AU ICMEs [32] We perform a similar analysis on the ICMEs measured at Ulysses in an attempt to define their spatial variations and 1 AU magnitudes in comparison to the plasma of the non-icme solar wind. We derive equations to characterize these spatial variations in a similar approach to that used for the solar wind observations except that we fit our model function to the mean value for each ICME event instead of the individual 1-hour measurements. 8of17

9 Table 1. Fitting Parameters for the Equations That Describe the Spatial Variations Along With Normalized-to-1-AU Values for Selected Properties of the Fast and Slow Solar Wind Measured at Ulysses Fast Solar Wind (q > j36j ) Slow Solar Wind g a b 1 AU Value (q =60 ) g a b 1 AU Value (q =0 ) Parameter Vp (km s 1 ) ± ± ± ± ± ± Va (km s 1 ) ± ± ± ± ± ± Tp 10 5 (K) 0.97 ± ± 0.03 (1.9 ± 0.4) ± ± 0.01 (2.6 ± 0.1) Np (cm 3 ) 1.86 ± ± ± ± ± ± N a (cm 3 ) 1.85 ± ± (3.4 ± 0.2) ± ± (2.0 ± 0.1) N a /N p ± ± (2.8 ± 0.3) ± ± (9.4 ± 0.3) rivi 2 (npa) 1.86 ± ± ± ± ± 0.02 (5.2 ± 0.4) rivi (kg m 2 s 1 ) 1.86 ± ± ± ± ± 0.04 (2.0 ± 0.1) jbrj (nt) 1.68 ± ± ± ± ± ± jbtj (nt) 1.09 ± ± ± ± ± ± jb N j (nt) 1.16 ± ± 0.01 (3 ± 4) ± ± 0.02 (4.0 ± 0.4) jbj (nt) 1.40 ± ± ± ± ± ± b p 0.03 ± ± ± ± ± 0.09 (0.7 ± 1) npktp (ppa) 2.94 ± ± ± ± ± 0.09 (7 ± 3) B 2 /2mo (ppa) 2.97 ± ± ± ± ± ± [33] Figure 7 shows plots of selected ICME properties versus heliolatitude for the events outlined in Appendix A. The data are summarized in the statistical representation used for the fast and slow solar wind observations in Figures 4, 5, and 6 except that instead of binning the data by heliolatitude and plotting the mean and variability for each bin, we plot the mean value and variability for each ICME event. A majority of the fast events were in the fast PCH wind and we omit these events from our fitting algorithm. [34] The plot of proton speed versus heliolatitude reveals that a large percentage of the ICME events were identified at latitudes below 40. This is not surprising as ICMEs are generally associated with the streamer belt and active regions, both of which are generally found below 40 for much of the solar cycle, especially during solar minimum [von Steiger and Richardson, 2006]. Above 45 there appear to be two populations of ICMEs, one with V p < 600 km s 1 and one with V p > 600 km s 1. The slow ICMEs at high latitudes were, in general, measured near solar maximum and the higher-speed ICMEs near solar minimum. The two highest-speed events were at lower latitudes, with the fastest ICME identified at 20. Both of these events appear to be embedded in stream interaction regions (see Figure 1), one observed during the first orbit and the other during the second. Also, most of the slower ICMEs show only small deviations from their mean proton speed within each event while the other properties all show large variations within the events. This is similar to the slow solar wind where V p is much less variable than the other solar wind properties [Bame et al., 1977]. [35] T p, N p and jbj all show radial variations comparable to those calculated for the slow and fast solar wind. Also, only small latitudinal variations are observed although there is much scatter in the mean values around these fits and in many cases large variation within the events themselves. This is reflected in the large uncertainties for the parameters (see Table 2) used to describe these spatial variations, especially in the latitudinal fits where the uncertainties in the slopes are of the same magnitude as the slopes themselves. [36] Table 2 presents results for all the plasma and magnetic field properties studied for the Ulysses ICMEs. The properties (Column 1) are the same as for the fast and slow solar wind. Since the ICMEs embedded in the fast PCH wind were removed from our analysis, we compare the bulk properties of these ICMEs with the results for the slow solar wind. [37] When normalized to 1 AU the values for the proton and alpha speeds, proton density and temperature are within 20% of slow solar wind values. The scaled to 1 AU alpha particle density and alpha/proton ratio are 0.32 cm 3 and 0.064, respectively, 3 times larger than observed in the slow wind. This is not surprising since an enhanced alpha/ proton ratio is one of the signatures used to identify ICMEs in the solar wind observations [Borrini et al., 1982]. The dynamic pressure and mass flux both have a radial dependence similar to the proton density with values when scaled to 1 AU of 2.71 npa and 5.54 kg m 2 s 1, respectively. Both of these values are larger than observed in the slow wind. [38] The values at 1 AU for the RTN field components and total field strength are significantly higher than for the low-speed wind. Again, this is not surprising since enhanced 9of17

10 Figure 7. Similar to Figure 4 but plotted in clockwise order are observations scaled to 1 AU of the proton speed (V p ), proton temperature (T p ), magnetic field strength (jbj), and proton density (N p ) versus heliolatiude for the ICMEs measured at Ulysses. The data are presented in the same statistical representation described in Figure 4 except that, instead of binning the data by heliolatitude, we plotted the mean and variability for each individual ICME event. ICMEs with a mean proton speed >650 km s 1 are shown as red circles, with the remaining events shown as black circles. The fast events (red circles) were not included in the derivation of the equations used to characterize the spatial variations in these events. magnetic field strength was also used as an ICME signature. The magnitude at 1 AU for the thermal pressure is 9.0 ppa, while the magnetic pressure is 33 ppa, both larger than observed in the slow wind. The ICME proton beta, which is the ratio of the proton thermal pressure to magnetic pressure, is 0.24 which is significantly reduced from the observed slow-wind value of This is also not surprising as low proton beta is another ICME signature used in this study. 5. Solar Cycle Variations [39] The extended duration of the Ulysses mission has also provided an opportunity to study the properties of the solar wind for two different solar cycles. Here, we discuss observations of the fast PCH wind measured during Ulysses first and third orbits, both of which occurred at and around solar minimum. During the previous two solar minima (solar cycles 21 and 22) the heliospheric current sheet tilt angle, which was calculated from a potential-field source surface model (R s s = 3.25 R s ) based on observations of the Sun s photospheric magnetic field from the Wilcox Solar Observatory (WSO), dropped to <5 (relative to the Sun s equator) [Hoeksema, 1995; html]. The tilt angle observed during the current minimum (solar cycle 23) has so far remained at or above 10 ( wso.stanford.edu/tilts.html), suggesting that this solar cycle minimum may be somewhat different. [40] Recent studies by McComas et al. [2008], Smith and Balogh [2008] and Issautier et al. [2008] have shown that some of the values of the plasma and magnetic field properties measured in the fast PCH flows during Ulysses current (3rd) orbit are reduced relative to those measured during the 1st orbit, which coincided with the previous solar minimum. McComas et al. [2008] observed a slight (3%) reduction in both the proton and alpha particle speeds, coupled with much larger decreases of 20% in the scaled proton density (N p R 2 ), 15% in the scaled proton temperature (T p R), and 20% in the scaled dynamic pressure (m p (N p V p 2 + 4N a V a 2 )(R/R o ) 2 ) and mass flux (m p (N p V p +4N a V a )(R/R o ) 2 ). 10 of 17

11 Table 2. Fitting Parameters for the Equations Which Describe the Spatial Variations Along With Normalizedto-1-AU Values for Selected Properties of the ICMEs Measured at Ulysses ICMEs Parameter g a b 1 AU Value (q =0 ) V p (km s 1 ) 0.01 ± ± ± V a (km s 1 ) 0.01 ± ± ± T p 10 5 (K) 0.86 ± ± 0.19 (3.6 ± 3.3) N p (cm 3 ) 1.78 ± ± ± N a (cm 3 ) 1.89 ± ± 0.05 (2 ± 1) N a /N p 0.24 ± ± 0.01 (3 ± 1) r i V 2 i (npa) 1.80 ± ± ± r i V i (kg m 2 s 1 ) 1.78 ± ± ± jb R j (nt) 1.60 ± ± ± jb T j (nt) 1.21 ± ± ± jb N j (nt) 1.07 ± ± ± jbj (nt) 1.29 ± ± ± b p 0.16 ± ± ± n p kt p (ppa) 2.38 ± ± ± B 2 /2m o (ppa) 2.32 ± ± ± Smith and Balogh [2008] observed a 1/3 reduction in the magnetic field strength while Issautier et al. [2008] showed that the electron density and temperature were down by 25% and 15%, respectively. Thus the fast solar wind from the PCHs is significantly weaker than that observed during the previous minimum. [41] On the basis of these findings, we have analyzed the fast PCH wind data from the first and third orbits separately, deriving corresponding expressions for their radial and latitudinal dependences. Table 3 summarizes the various properties (see Column 1). Columns 2 4 show the fit parameters of the expressions for the spatial variations and column 5 the magnitudes at 1 AU for the various properties measured above ±36 in heliolatitude during the 1st orbit. Columns 6 9 show similar results for the 3rd orbit. These expressions are for latitudes above ±40 and heliocentric distances within 4.5 AU. We chose a slightly higher latitude cut-off compared to the first orbit as the band of solar wind variability extends to higher latitudes for the current Ulysses orbit [McComas et al., 2006]. The reason for the 4.5 AU distance cut-off is that we have very little fast wind data past this distance. Unlike the first orbit, Ulysses will not return data from the complete northern portion of its current polar orbit. As of this writing, we have included plasma data up to 5 April 2008 where Ulysses was at a distance of 2.64 AU and 65.6 in latitude. [42] Comparing the values when normalized to 1 AU (q = 60 ) for orbits 1 and 3 we observe differences in the plasma properties similar to those described by McComas et al. [2008]. The proton and alpha particle speeds were slightly reduced (2% and 3%, respectively) while much larger reductions were observed in the values at 1 AU for the proton temperature (19%), proton density (33%), dynamic pressure (38%) and mass flux (47%). No changes in the value at 1 AU for the alpha/proton ratio between orbits 1 and 3 were observed. We confirm the nearly 1/3 drop in the magnetic field strength identified by Smith and Balogh [2008] along with significant decreases in the magnitude of the magnetic field RTN components. The proton thermal and magnetic pressures were reduced by a factor of 2 and 2.4, respectively, while the proton beta increased by a factor of 1.4 (1.3 compared to 0.91). [43] The latitudinal variations observed over the 3rd orbit are consistent with those observed during the 1st orbit indicating that the fast PCH wind shows only weak latitudinal gradients for two successive solar minima. We find that the direction and magnitude of these gradients are the same for both orbits in all properties except for the alpha particle density and alpha/proton ratio. This, coupled with the small uncertainties in the parameters describing these variations suggests that these small latitudinal variations are real. [44] We observe slight differences in the radial variations between orbit 1 and 3, most notably in properties that depend on the proton density and magnetic field. N p falls as R 1.7 compared to R 1.9 for the first orbit. This affected the dynamic pressure, mass flux and thermal pressure, all of which were shown to fall less sharply with distance as compared to the 1st orbit. Separate analysis of the results for the northern and southern components of the 3rd orbit show that this R 1.7 dependence appears to be a good representation of the data. Regarding the magnetic field, a less rapid decrease with distance in the RTN components and total field strength were observed for 3rd orbit. 6. Discussion [45] In this study we examined interplanetary plasma and magnetic field observations from Ulysses taken over three polar orbits of the Sun. These measurements were used to compute the radial and latitudinal variations and calculate the magnitudes at 1 AU for selected properties of the fast and slow solar wind and ICMEs. We fit the data to a simple equation where the radial gradients were approximated by a power law and the latitudinal variations by a linear function. After removing the radial gradients, only small latitudinal variations were observed in the three types of heliospheric plasma considered here. The latitudinal gradients in the slow and fast wind appear to be statistically significant owing to the relatively small uncertainties in the derived parameters used to express them. The very small variations in the fast wind (<1% deg. 1 ) are consistent with observations over the 1st orbit reported by McComas et al. [2000], 11 of 17

12 Table 3. Fitting Parameters for the Equations Which Describe the Spatial Variations Along With Normalized-to-1-AU Values for Selected Properties of the Fast Solar Wind Measured During Ulysses 1st and 3rd Polar Orbits 1st Orbit (q > j36 j) 3rd Orbit (q > j40 j, R < 4.5 AU) g a b 1 AU Value (q =60 ) g a b 1 AU Value (q =60 ) Parameter V p (km s 1 ) ± ± ± ± ± ± V a (km s 1 ) ± ± ± ± ± ± T p 10 5 (K) 1.02 ± ± 0.04 (2.2 ± 0.5) ± ± ± Np (cm 3 ) 1.92 ± ± ± ± ± ± Na (cm 3 ) 1.89 ± ± (3.3 ± 0.3) ± ± (4.7 ± 0.3) Na/Np ± ± (4.3 ± 0.4) ± ± (0.6 ± 5) rivi 2 (npa) 1.93 ± ± ± ± ± 0.04 (6.0 ± 0.1) r i V i (kg m 2 s 1 ) 1.92 ± ± ± ± ± 0.06 (2.0 ± 0.1) jb R j (nt) 1.78 ± ± ± ± ± ± jb T j (nt) 1.15 ± ± ± ± ± ± jbnj (nt) 1.19 ± ± 0.02 (4 ± 5) ± ± ± jbj (nt) 1.48 ± ± ± ± ± ± bp 0.08 ± ± ± ± ± ± n p kt p (ppa) 2.94 ± ± ± ± ± 0.15 (4.2 ± 0.2) B 2 /2m o (ppa) 0.97 ± ± ± ± ± ± of 17 indicating that the fast PCH wind has shown only minor variations with latitude for two successive solar minima. [46] We also studied solar cycle variations in the solar wind by examining the properties of the fast PCH flows measured during Ulysses 1st and 3rd orbits, both of which occurred during a period near solar minimum. We found for several plasma and magnetic field properties of the fast PCH wind that their magnitudes when normalized to 1 AU were reduced for the 3rd orbit relative to the 1st, the reduction being of the order of those described by McComas et al. [2008] and Smith and Balogh [2008]. The values for the proton and alpha particle speeds were only slightly reduced (2 3%) compared to the changes observed in the proton temperature (19%), density (33%), dynamic pressure (38%), mass flux (47%) and magnetic field strength (32%). McComas et al. [2008] reported that the observed reduction in the mass flux and dynamic pressure coupled with only slight changes in the speed for the PCH wind supported the model of Leer and Holzer [1980] for solar wind energization below the sonic point in the corona. Also, Schwadron and McComas [2008] have recently shown that the linear relationship between the solar wind power (rv 3 R 2 /2) and the Sun s open magnetic flux density (jhb r R 2 ij), with the Sun providing 600 kw/wb of power to the solar wind, could be a possible explanation for these observations. They suggest that the weaker dynamic pressure and mass flux observed in the PCH wind during the current minimum can be directly correlated to the weaker magnetic field strength which has the effect of reducing the amount of open magnetic flux from the Sun and, hence, the power available for energizing the solar wind. [47] For the spatial dependences, the latitudinal variations appear to be similar in both direction and magnitude for the 1st and 3rd orbits, showing only small variations for both orbits. The radial gradients for the PCH wind appear somewhat different. Our observations indicate an R 1.7 dependence for proton density during the 3rd orbit compared to the R 1.9 dependence during the 1st. The RTN components and total strength of the magnetic field fall less rapidly with distance in the 3rd orbit as well. [48] Although the properties of the fast PCH wind show distinct differences between Ulysses 1st and 3rd orbit, we provide values for the combined measurements from all three polar orbits to characterize the global properties of the fast wind. The fact that Ulysses measurements of the fast PCH wind only cover two periods of solar minima makes it difficult to say which of the observations are more representative. [49] This paper represents the first attempt to characterize the global behavior in the properties of the slow wind as measured by Ulysses. The values at 1 AU for the proton and alpha speeds, proton temperature, alpha/proton ratio and dynamic pressure were all significantly reduced compared to the fast wind while the proton density showed significant enhancement. The magnitude at 1 AU for the alpha density and radial component of the IMF (jb R j) were about the same in the slow and fast wind. In most cases, the radial and latitudinal gradients in the slow wind were similar to those for the fast PCH wind, the exceptions being the radial gradients for the proton temperature, magnetic field strength, and parameters dependent on these properties such as the proton thermal and magnetic pres-

13 sures. The radial dependence of T p in the slow wind was R 0.7 compared to the R 1.0 profile measured for the fast wind. Both of these deviate from the R 4/3 variation predicted for adiabatic expansion. That the proton temperature gradient is steeper in the high-speed wind than in the slow wind has been discussed previously [Mihalov, 1983] and seems to imply a source of heating which is enhanced in or unique to the low-speed wind. Using Voyager measurements from 1 to 10 AU Gazis and Lazarus [1982] showed that the proton temperature at low latitudes did not decrease adiabatically, instead varying as R 0.7±0.2. These authors proposed that this was caused by the conversion of bulk flow kinetic energy into thermal energy either from shocks or from the interaction of high- and low-speed streams. In our observations, we noted the presence of SIRs in both the slow and fast wind data sets. These interaction regions were easily removed from our analysis of the fast wind data set owing to the fact that they were observed only up to specific latitudes. For the slow wind, these SIRs were much more difficult to remove and hence, were included in our analysis. [50] Heating due to SIRs has been studied in detail. Gosling et al. [1972] found at 1 AU that roughly half of the temperature rise observed at the leading edge of highspeed streams could be attributed to compressional heating. Burlaga and Ogilvie [1973] reported that heating due to the adiabatic compression of streams in the interplanetary medium accounted for an increase of 15% in the average temperature of the solar wind over a timescale of several solar rotations. Jian et al. [2005] surveyed the properties of SIRs in the Wind data set from In cases where the interaction was subsonic they found only weak heating of the solar wind with the heating confined to a region near the stream interface. For a supersonic interaction, they reported that heating from shocks resulted in more significant heating of the solar wind with the heating region moving away from the stream interface. Whang et al. [1990] investigated solar wind heating from shocks using observations from 1 to 30 AU in conjunction with an MHD shock simulation. They found that shock heating alone could be responsible for heating the solar wind from 1 to 15 AU. Gazis and Lazarus [1983] also showed that shock frequency peaked at 5 AU while the shock strength peaked between 4 and 5 AU. Thus it is likely that compressional heating from shocks formed at the interface of SIRs is responsible for the different T p radial dependences observed in the slow and fast solar wind between 1.3 and 5.4 AU. Other potential sources include turbulent heating [Smith et al., 2001], heating from kinetic Alfvén waves [Leamon et al., 1999] or, perhaps, through magnetic reconnection [Leamon et al., 2000], although this last scenario appears somewhat doubtful, at least in the fast wind [Gosling, 2007]. [51] Finally, we extended the Ulysses-SWOOPS ICME list through the 5 April 2008, essentially to the end of the mission. We found that the majority of ICMEs had speeds below 500 km s 1. We also performed a statistical analysis on the average properties of these ICME events, deriving expressions characterizing their spatial variations and magnitudes at 1 AU. The radial dependence in the ICME properties did not differ drastically from those measured for the slow and fast solar wind. The latitudinal gradients were also comparatively small although the large uncertainties in the parameters used to describe these variations make it difficult determine whether these variations are statistically significant or not. 7. Summary [52] In this study we examined observations of the interplanetary plasma and magnetic field from Ulysses taken over three polar orbits of the Sun. These measurements were used to derive empirical equations to characterize the radial and latitudinal dependences and calculate the magnitudes at 1 AU for selected properties of the fast and slow solar wind and ICMEs, covering the entire (>16 year) out of ecliptic segment of the Ulysses mission. The key observations from this study are as follows. [53] 1. The Ulysses-SWOOPS ICME list was extended essentially through the end of the Ulysses mission, with 27 new events identified from January 2003 through 5 April [54] 2. A majority of the ICMEs identified at Ulysses had a mean proton speed below 500 km/s. [55] 3. The equations describing the spatial variations revealed that the proton temperature in the slow-wind falls less sharply with distance than in the fast wind. This is likely due to compressional heating from shocks formed at the leading edge of SIRs in the highly variable slow solar wind. [56] 4. The total IMF strength was also shown to fall less rapidly with distance in the slow wind than the fast, a result of the tighter Parker spiral structure in the low-speed wind. [57] 5. The fast and slow wind showed only minor latitudinal gradients (<1% deg. 1 ) and the slow solar wind was cooler, more dense and had a lower alpha/proton ratio than the fast wind, consistent with previous observations. [58] 6. The tangential component of the IMF (jb T j) was much stronger in the slow wind relative to the fast wind while the total IMF strength at 1 AU was enhanced as well. [59] 7. The magnitude at 1 AU of the alpha particle density and the radial component of the IMF (jb R j)are about the same in the slow and fast wind. [60] 8. The spatial variations in the ICME properties were similar to those observed for the solar wind, although with much larger uncertainties in the parameters used to describe them. [61] 9. The solar cycle variations in properties of the solar wind were explored by comparing the spatial variations and magnitudes at 1 AU for the fast PCH flows observed during Ulysses 1st and 3rd orbits. A reduction in the magnitude at 1 AU was observed for the proton and alpha speeds (2 3%), proton temperature (19%) and density (33%), dynamic pressure (38%), mass flux (47%) and magnetic field strength (32%) in the PCH wind measured in the 3rd orbit compared to the 1st. These results are in agreement with reports of weaker PCH flows for the current solar minimum. [62] 10. The proton density and IMF field strength in the fast PCH wind appeared to fall less sharply with distance in the 3rd orbit than in the 1st. Appendix A: Ulysses ICME List ( ) [63] Here we present the Ulysses ICME list for the period from 18 February 1992 through 5 April The events from 1992 through 2002 were obtained directly from the 13 of 17

14 Table A1. Ulysses ICME Event List Year DOY/Start Time DOY/Stop Time Heliolatitude (Deg.) Heliolongitude (Deg.) Ulysses-Sun Distance (AU) hv p i (km s 1 ) hbi (nt) /17:30 73/06: S ± ± /04:00 77/19: S ± ± /05:00 81/21: S ± ± /12:30 85/18: S ± ± /07:30 88/09:00 8.3S ± ± /09:20 107/00: S ± ± /17:30 133/19: S ± ± /13:00 201/12: S ± ± /09:20 209/06: S ± ± /22:00 316/22: S ± ± /15:00 323/00: S ± ± /17:00 10/10: S ± ± /07:00 113/22: S ± ± /21:20 164/01: S ± ± /01:15 205/20: S ± ± /20:00 242/11: S ± ± /17:00 41/14: S ± ± /10:45 59/16: S ± ± /05:45 112/08: S ± ± /06:00 36/20: S ± ± /08:00 295/16: N ± ± /03:00 350/20: N ± ± /23:00 11/06: N ± ± /19:00 43/10: N ± ± /21:30 62/13: N ± ± /02:00 78/02: N ± ± /12:30 148/12: N ± ± /15:30 219/22: N ± ± /03:00 231/05: N ± ± /00:00 236/01: N ± ± /22:00 246/21: N ± ± /07:00 278/05: N ± ± /12:00 287/02: N ± ± /21:00 303/10: N ± ± /06:00 320/14: N ± ± /03:30 355/00: N ± ± /13:30 361/14: S ± ± /06:30 9/05: S ± ± /13:30 26/00: S ± ± /10:00 33/15: S ± ± /19:00 38/10: S ± ± /18:30 45/18: S ± ± /18:00 54/10: S ± ± /01:00 61/00: S ± ± /18:00 71/12: S ± ± /13:30 87/18: S ± ± /07:30 95/19: S ± ± /10:00 98/03: S ± ± /07:00 101/12: S ± ± /03:00 111/09: S ± ± /02:00 134/00: S ± ± /09:00 154/11: S ± ± /15:00 165/18: S ± ± /03:00 173/23: S ± ± /21:00 180/11: S ± ± /05:30 187/22: S ± ± /04:30 198/10: S ± ± /18:00 210/08: S ± ± /18:00 214/09: S ± ± /08:30 225/19: S ± ± /02:00 227/12: S ± ± /18:00 231/22: S ± ± /22:00 235/17: S ± ± /16:00 247/14: S ± ± /09:30 252/20: S ± ± /01:00 264/00: S ± ± /04:00 278/22: S ± ± /07:00 288/18: S ± ± /01:00 292/01: S ± ± /08:00 297/13: S ± ± /15:00 303/14: S ± ± /04:00 315/10: S ± ± of 17

15 Table A1. (continued) Year DOY/Start Time DOY/Stop Time Heliolatitude (Deg.) Heliolongitude (Deg.) Ulysses-Sun Distance (AU) hv p i (km s 1 ) hbi (nt) 319/20:00 322/20: S ± ± /02:00 336/10: S ± ± /07:00 346/10: S ± ± /15:00 352/18: S ± ± /01:00 14/23: S ± ± /17:00 25/11: S ± ± /19:00 39/05: S ± ± /05:30 50/12: S ± ± /18:00 64/22: S ± ± /14:00 78/05: S ± ± /01:00 98/16: S ± ± /21:00 102/16: S ± ± /07:30 124/18: S ± ± /03:00 127/01: S ± ± /16:00 141/03: S ± ± /12:00 169/22: S ± ± /05:00 176/05: S ± ± /01:30 184/06: S ± ± /14:00 187/17: S ± ± /20:00 217/10: S ± ± /23:30 232/12: S ± ± /04:00 247/06: S ± ± /07:30 280/22: S ± ± /12:00 289/22: S ± ± /16:00 303/10: S ± ± /16:30 348/12: S ± ± /10:30 21/02: S ± ± /22:00 44/21: S ± ± /09:00 78/09: S ± ± /03:00 88/14: S ± ± /12:00 92/07: S ± ± /00:00 102/12: S ± ± /12:00 111/22: S ± ± /10:30 124/04: S ± ± /22:30 135/04: S ± ± /18:00 148/20: S ± ± /01:00 180/00: S ± ± /00:00 198/02: S ± ± /11:30 225/09: S ± ± /23:00 257/19: S ± ± /00:00 294/22: S ± ± /01:30 314/12: S ± ± /09:30 345/06: S ± ± /23:00 24/20: S ± ± /15:00 33/12: S ± ± /01:00 78/23: S ± ± /18:00 93/08: S ± ± /21:30 102/02: S ± ± /10:00 114/05: S ± ± /06:00 132/10: S ± ± /00:00 142/02: N ± ± /07:30 160/04: N ± ± /07:00 164/23: N ± ± /13:30 179/12: N ± ± /10:00 188/13: N ± ± /11:00 206/02: N ± ± /18:00 225/16: N ± ± /13:00 237/01: N ± ± /12:30 272/14: N ± ± /02:30 304/22: N ± ± /14:40 314/15: N ± ± /12:50 319/16: N ± ± /09:30 333/00: N ± ± /08:00 003/15: N ± ± /00:00 019/12: N ± ± /15:00 032/23: N ± ± /18:00 043/23: N ± ± /09:00 074/07: N ± ± /18:00 098/12: N ± ± /06:00 127/06: N ± ± /20:00 147/11: N ± ± /10:00 157/04: N ± ± of 17

16 Table A1. (continued) Year DOY/Start Time DOY/Stop Time Heliolatitude (Deg.) Heliolongitude (Deg.) Ulysses-Sun Distance (AU) hv p i (km s 1 ) hbi (nt) 164/18:00 171/10: N ± ± /12:00 199/17: N ± ± /00:00 276/00: N ± ± /00:00 292/00: N ± ± /16:00 301/16: N ± ± /18:00 327/06: N ± ± /02:00 23/08: N ± ± /01:00 35/19: N ± ± /17:00 45/11: N ± ± /18:00 58/02: N ± ± /08:00 100/16: N ± ± /12:30 133/05: N ± ± /15:00 161/01: N ± ± /08:00 169/06: N ± ± /06:00 232/20: N ± ± /01:00 280/10: N ± ± /04:00 324/23: N ± ± /00:00 354/07: N ± ± /03:00 28/16: N ± ± /22:00 58/10: S ± ± /01:00 72/16: S ± ± /22:00 209/10: S ± ± /23:00 239/02: S ± ± /06:00 263/14: S ± ± /06:00 a 38/23: S ± ± /08:00 a 48/10: S ± ± /10:00 a 72/22: S ± ± /04:00 a 114/20: S ± ± /02:30 a 140/02: S ± ± /16:30 a 145/17: S ± ± /18:00 a 360/18: S ± ± /14:00 113/22: S ± ± /13:20 186/15: S ± ± 2.48 a Electron data not available for event; DOY = day of year, V p = proton speed, B = interplanetary magnetic field. Ulysses ICME list ( Events from the start of 2003 onward were identified by the authors of this manuscript. The method of identifying these events was described in section 3. For each event we provide the day of year (DOY) and event start/stop time, spacecraft heliolatitude, heliolongitude, and radial distance using the heliographic inertial (HGI) coordinate system and the average proton velocity and magnetic field strength of the plasma within the event timeframe (Table A1). [64] Acknowledgments. This work was funded by NASA under the Ulysses program. We thank those responsible for providing data on the tilt angle of the heliospheric current sheet from the Wilcox Solar Observatory. R.W. Ebert would like to thank Richard Menchaca for his assistance with the figures and the National Science and Engineering Research Council of Canada for their financial contribution to his work. [65] Amitava Bhattacharjee thanks Richard Marsden and Paulett Liewer for their assistance in evaluating this paper. References Asbridge,J.R.,S.J.Bame,W.C.Feldman,andM.D.Montgomery (1976), Helium and hydrogen velocity differences in the solar wind, J. Geophys. Res., 81(16), Balogh, A., T. J. Beek, R. J. Forsyth, P. C. Hedgecock, R. J. Marquedant, E. J. Smith, D. J. Southwood, and B. T. Tsurutani (1992), The magnetic field investigation on the Ulysses mission Instrumentation and preliminary scientific results, Astron. Astrophys., Suppl. Ser., 92(2), 221. Bame, S. J., J. R. Asbridge, W. C. Feldman, and J. T. Gosling (1977), Evidence of a structure-free state at high solar wind speeds, J. Geophys. Res., 82(10), Bame, S. J., D. J. McComas, B. L. Barraclough, J. L. Phillips, K. J. Sofaly, J. C. Chavez, B. E. Goldstein, and R. K. Sakurai (1992), Ulysses solar wind plasma experiment, Astron. Astrophys., Suppl. Ser., 92(2), 237. Borrini, G., J. T. Gosling, S. J. Bame, and W. C. Feldman (1982), Helium abundance enhancements in the solar wind, J. Geophys. Res., 87(A9), Burlaga, L. F., and K. W. Ogilvie (1973), Solar wind temperature and speed, J. Geophys. Res., 78(13), Burlaga, L., E. Sittler, F. Mariani, and R. Schwenn (1981), Magnetic loop behind an interplanetary shock: Voyager, Helios, and IMP 8 observations, J. Geophys. Res., 86(A8), Forsyth, R. J., A. Balogh, and E. J. Smith (2002), The underlying direction of the heliospheric magnetic field through the Ulysses first orbit, J. Geophys. Res., 107(A11), 1405, doi: /2001ja Gazis, P. R., and A. J. Lazarus (1982), Voyager observations of solar wind proton temperature: 1 10 AU, J. Geophys. Res., 9(4), Gazis, P. R., and A. J. Lazarus (1983), The radial evolution of the solar wind, 1 10 AU, in Solar Wind Five, edited by M. Neugebauer, NASA Conf. Publ., CP 2280, Gosling, J. T. (1990), Coronal mass ejections and magnetic flux ropes in interplanetary space, in Physics of Magnetic Flux Ropes, Geophys. Monogr. Ser., vol. 58, edited by C. T. Russell, E. R. Priest, and L. C. Lee, pp , AGU, Washington, D. C. Gosling, J. T. (1996), Corotating and transient solar wind flows in three dimensions, Annu. Rev. Astron. Astrophys., 34, 35. Gosling, J. T. (2007), Observations of magnetic reconnection in the turbulent high-speed solar wind, Astrophys. J., 671, L73 L76. Gosling, J. T., A. J. Hundhausen, V. Pizzo, and J. R. Asbridge (1972), Compressions and rarefactions in the solar wind: Vela 3, J. Geophys. Res., 77(28), Gosling, J. T., V. Pizzo, and S. J. Bame (1973), Anomalously low proton temperatures in the solar wind following interplanetary shock waves Evidence for magnetic bottles?, J. Geophys. Res., 78(13), Gosling, J. T., E. Hildner, R. M. MacQueen, R. H. Munro, A. I. Poland, and C. L. Ross (1974), Mass ejections from the Sun: A view from Skylab, J. Geophys. Res., 79(31), Gosling, J. T., G. Borrini, J. R. Asbridge, S. J. Bame, W. C. Feldman, and R. T. Hansen (1981), Coronal streamers in the solar wind at 1 AU, J. Geophys. Res., 86(A7), of 17

17 Gosling, J. T., D. N. Baker, S. J. Bame, W. C. Feldman, R. D. Zwickl, and E. J. Smith (1987), Bidirectional solar wind electron heat flux events, J. Geophys. Res., 92(A8), Gosling, J. T., S. J. Bame, D. J. McComas, J. L. Phillips, E. E. Scime, V. J. Pizzo, B. E. Goldstein, and A. Balogh (1994), A forward-reverse shock pair in the solar wind driven by over-expansion of a coronal mass ejection: Ulysses observations, Geophys. Res. Lett., 21(3), Hoeksema, J. T. (1995), The large scale structure of the heliospheric current sheet during the Ulysses epoch, Space Sci. Rev., 72, Hundhausen, A. J. (1968), Direct observations of solar wind particles, Space Sci. Rev., 8, 690. Hundhausen, A. J. (1999), The Many Faces of the Sun: A Summary of the Results From NASA s Solar Maximum Mission, edited by K. T. Strong et al., p. 143, Springer, New York. Issautier, K., G. Le-Chat, N. Meyer-Vernet, M. Moncuquet, S. Hoang, R. J. MacDowall, and D. J. McComas (2008), Electron properties of highspeed solar wind from polar coronal holes obtained by Ulysses thermal noise spectroscopy: Not so dense, not so hot, Geophys. Res. Lett., 35, L19101, doi: /2008gl Jian, L., C. T. Russell, J. T. Gosling, and J. G. Luhmann (2005), Measurements of heating at stream-stream interactions, in Connecting Sun and Heliosphere, Proc. Solar Wind 11/SOHO 16, edited by B. Fleck, T. H. Zurbuchen, and H. Lacoste, p. 731, SP-592, ESA, Noordwijk, Netherlands. Klein, L. W., and L. F. Burlaga (1982), Interplanetary magnetic clouds at 1 AU, J. Geophys. Res., 87(A2), Kreiger, A. S., A. F. Timothy, and E. C. Roelof (1973), A coronal hole and its identification as the source of a high velocity solar wind stream, Sol. Phys., 29, 505. Leamon, R. J., C. W. Smith, N. F. Ness, and H. K. Wong (1999), Dissipation range dynamics: Kinetic Alfvén waves and the importance of b e, J. Geophys. Res., 104(A10), 22,331 22,344. Leamon, R. J., W. H. Matthaeus, C. W. Smith, G. P. Zank, D. J. Mullan, and S. Oughton (2000), MHD driven kinetic dissipation in the solar wind and corona, Astrophys. J., 537, Leer, E., and T. E. Holzer (1980), Energy addition in the solar wind, J. Geophys. Res., 85(A9), Liu, Y., J. D. Richardson, and J. W. Belcher (2005), A statistical study of the properties of interplanetary coronal mass ejections from 0.3 to 5.4 AU, Planet. Space Sci., 53, Liu, Y., J. D. Richardson, J. W. Belcher, J. C. Kasper, and H. A. Elliot (2006), Thermodynamic structure of collision-dominate expanding plasma: Heating of interplanetary coronal mass ejections, J. Geophys. Res., 111, A01102, doi: /2005ja Marsch, E., K.-H. Mühlhäuser, H. Rosenbauer, R. Schwenn, and K. U. Denskat (1981), Pronounced proton core temperature anisotropy, ion differential speed, and simultaneous Alfvén wave activity in slow solar wind at 0.3 AU, J. Geophys. Res., 86(A11), McComas, D. J., et al. (1998a), Ulysses return to the slow solar wind, Geophys. Res. Lett., 25(1), 1 4. McComas, D. J., S. J. Bame, P. Barker, W. C. Feldman, J. L. Phillips, P. Riley, and J. W. Griffee (1998b), Solar wind electron proton alpha monitor (SWEPAM) for the advanced composition explorer, Space Sci. Rev., 86, 563. McComas, D. J., B. L. Barraclough, H. O. Funsten, J. T. Gosling, E. Santiago-Muñoz, R. M. Skoug, B. E. Goldstein, M. Neugebauer, P. Riley, and A. Balogh (2000), Solar wind observations over Ulysses first full polar orbit, J. Geophys. Res., 105(A5), 10,419 10,433. McComas, D. J., H. A. Elliott, N. A. Schwadron, J. T. Gosling, R. M. Skoug, and B. E. Goldstein (2003), The three-dimensional solar wind around solar maximum, Geophys. Res. Lett., 30(10), 1517, doi: / 2003GL McComas, D. J., H. A. Elliott, J. T. Gosling, and R. M. Skoug (2006), Ulysses observations of very different heliospheric structure during the declining phase of solar activity cycle 23, Geophys. Res. Lett., 33, L09102, doi: /2006gl McComas, D. J., R. W. Ebert, H. A. Elliott, B. E. Goldstein, J. T. Gosling, N. A. Schwadron, and R. M. Skoug (2008), Weaker solar wind from the polar coronal holes and the whole Sun, Geophys. Res. Lett., 35, L18103, doi: /2008gl Mihalov, J. (1983), Solar wind proton temperature gradient from Pioneer 10 and 11 and closer data, in Solar Wind Five, edited by M. Neugebauer, NASA Conf. Pub. CP-2280, p Montgomery, M. D., J. R. Asbridge, S. J. Bame, and W. C. Feldman (1974), Solar wind electron temperature depressions following some interplanetary shock waves: Evidence for magnetic merging?, J. Geophys. Res., 79(22), Neugebauer, M. and R. Goldstein (1997), Particle and field signatures of coronal mass ejections in the solar wind, in Coronal Mass Ejections, Geophys. Monogr. Ser., vol. 99, edited by N. Crooker, J. A. Joselyn, and J. Feynman, p. 245, AGU, Washington, D. C. Neugebauer, M., and C. W. Snyder (1962), Solar plasma experiment, Science, 138, Neugebauer, M., P. C. Liewer, E. J. Smith, R. M. Skoug, and T. H. Zurbuchen (2002), Sources of the solar wind at solar activity maximum, J. Geophys. Res., 107(A12), 1488, doi: /2001ja Neugebauer, M., J. T. Steinberg, R. L. Tokar, B. L. Barraclough, E. E. Dors, R. C. Wiens, D. E. Gingerich, D. Luckey, and D. B. Whiteaker (2003), Genesis on-board determination of the solar wind flow regime, Space Sci. Rev., 105, Neugebauer, M., et al. (1998), Spatial structure of the solar wind and comparisons with solar data and models, J. Geophys. Res., 103(A7), 14,587 14,599. Neugebauer, M., et al. (2007), Encounter of the Ulysses spacecraft with the ion tail of comet McNaught, Astrophys. J., 667, Parker, E. N. (1958), Dynamics of the interplanetary gas and magnetic fields, Astrophys. J., 128, 664. Richardson, J. D., K. I. Paularena, A. J. Lazarus, and J. W. Belcher (1995), Radial evolution of the solar wind from IMP 8 to Voyager 2, Geophys. Res. Lett., 22(4), Schwadron, N. A., and D. J. McComas (2008), The solar wind power from magnetic flux, Astrophys. J., 686(1), L33 L36, doi: / Schwenn, R. (2006), Solar wind sources and their variations over the solar cycle, Space Sci. Rev., 124, Schwenn, R., M. D. Montgomery, H. Rosenbauer, H. Moggenrider, K. H. Muehlhaeuser, S. J. Bame, W. C. Feldman, and R. T. Hansen (1978), Direct observations of the latitudinal extent of a high speed stream in the solar wind, J. Geophys. Res., 83(A3), Sheeley, N. R., Jr., R. A. Howard, M. J. Koomen, D. J. Michels, R. Schwenn, K. H. Muehlhaeuser, and H. Rosenbauer (1985), Coronal mass ejections and interplanetary shocks, J. Geophys. Res., 90(A1), Smith, E. J., and A. Balogh (2008), Decrease in heliospheric magnetic flux in this solar minimum: Recent Ulysses magnetic field observation, Geophys. Res. Lett., 35, L22103, doi: /2008gl Smith, E. J., A. Balogh, M. Neugebauer, and D. McComas (1995), Ulysses observations of the Alfvén waves in the southern and northern solar hemispheres, Geophys. Res. Lett., 22(23), Smith, C. W., W. H. Matthaeus, G. P. Zank, and N. F. Ness (2001), Heating of the low-latitude solar wind by dissipation of turbulent magnetic fluctuations, J. Geophys. Res., 106(A5), Tousey, R., R. A. Howard, and M. J. Koomen (1974), The frequency and nature of coronal transient events observed by OSO-7, Bull. Am. Astron. Soc., 6, 295. von Steiger, R., and J. D. Richardson (2006), ICMEs in the outer heliosphere and at high latitudes: An introduction, Space Sci. Rev., 123(1 3), , doi: /s z. Wang, C., and J. D. Richardson (2004), Interplanetary coronal mass ejections observed by Voyager 2 between 1 and 30 AU, J. Geophys. Res., 109, A06104, doi: /2004ja Whang, Y. C., S. Liu, and L. F. Burlaga (1990), Shock heating of the solar wind plasma, J. Geophys. Res., 95(A11), 18,769 18,769. Zurbuchen, T. H., and I. G. Richardson (2006), In-situ solar wind and magnetic field signatures of interplanetary coronal mass ejections, Space Sci. Rev., 123(1 3), 31. R. W. Ebert, H. A. Elliott, and D. J. McComas, Space Science and Engineering Division, Southwest Research Institute, P.O. Drawer 28510, San Antonio, TX , USA. (rebert@swri.edu) R. J. Forsyth, Space and Atmospheric Physics Group, Blackett Laboratory, Imperial College London, Prince Consort Road, London SW7 2BW, UK. J. T. Gosling, Laboratory for Atmospheric and Space Physics, University of Colorado, 1234 Innovation Drive, Boulder, CO 80303, USA. 17 of 17