Plasma convection in the nightside magnetosphere of Saturn determined from energetic ion anisotropies

Planetary and Space Science 91 (2014) 1–13

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Plasma convection in the nightside magnetosphere of Saturn determined from energetic ion anisotropies M. Kane a,n, D.G. Mitchell b, J.F. Carbary b, S.M. Krimigis b a b

Harford Research Institute, 1411 Saratoga Dr, Bel Air, MD 21014, USA Johns Hopkins University, Applied Physics Laboratory, Laurel, MD, USA

art ic l e i nf o

a b s t r a c t

Article history: Received 13 April 2013 Received in revised form 30 August 2013 Accepted 7 October 2013 Available online 20 November 2013

The Cassini Ion and Neutral Camera measures intensities of hydrogen and oxygen ions and neutral atoms in the Saturnian magnetosphere and beyond. We use the measured intensity spectrum and anisotropy of energetic hydrogen and oxygen ions to detect, qualify, and quantify plasma convection. We find that the plasma azimuthal convection speed relative to the local rigid corotation speed decreases with radial distance, lagging the planetary rotation rate, and has no significant local time dependences. Plasma in the dusk-midnight quadrant sub-corotates at a large fraction of the rigid corotation speed, with the primary velocity being azimuthal but with a distinct radially outward component. The duskside velocities are similar to those obtained from earlier orbits in the midnight-dawn sector, in contrast to the depressed velocities measured at Jupiter using Energetic Particles Detector measurements on the Galileo spacecraft in the dusk-midnight quadrant. We find significant radial outflow in most of the nightside region. The radial component of the flow decreases with increasing local time in the midnight-dawn sector and reverses as dawn is approached. This and previous results are consistent with a plasma disk undergoing a centrifugally induced expansion as it emerges into the nightside, while maintaining partial rotation with the planet. The magnetodisk expansion continues as plasma rotates across the tail to the dawnside. We do not see evidence in the convection pattern for steady state reconnection in Saturn's magnetotail. The outermost region of the magnetodisk, having undergone expansion upon emerging from the dayside magnetopause confinement, is unlikely to recirculate back into the dayside. We conclude that plasma in the outer magnetodisk [at either planet] rotates from the dayside, expands at the dusk flank, but remains magnetically connected to the respective planet while moving across the tail until it interacts with and is entrained into the dawnside magnetosheath flow. This interaction causes plasma in the outer magnetospheric regions of Jupiter and Saturn to decouple from the planet and exhaust tailward down a dawnside low latitude boundary layer. Magnetospheric plasma will also interact with the dayside magnetosheath plasma, moving across the boundary [enhanced by shear instability] and into the magnetosheath, where it is lost to the magnetosphere with the magnetosheath flow. & 2013 Elsevier Ltd. All rights reserved.

Keywords: Saturn Magnetosphere Convection

1. Introduction The magnetospheres of Jupiter and Saturn are rotationally dominated to large distances from the respective planets, in contrast to the case at Earth. A primary signature of this dominance is the azimuthal convection of the plasma detected by preCassini studies of the Saturnian magnetosphere (Bridge et al., 1981, 1982; Richardson, 1986, 1998; Richardson and Sittler, 1990) using data from the Voyager Plasma (PLS) experiments. Instruments detecting higher energy plasma have been successfully used to measure convection using the observed particle anisotropies. For example, ion anisotropies at Jupiter and Saturn were investigated

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(Carbary et al., 1981; 1983) using the ion channels of the Low Energy Charged Particle (LECP) detector (Krimigis et al., 1977) on Voyagers 1 and 2. With the arrival of Cassini at Saturn, numerous orbits provide a wealth of data to perform convection studies under a variety of magnetospheric and interplanetary conditions. We have undertaken such a study using the Ion and Neutral Camera (INCA) detector (see section 2 below) and have demonstrated (Kane et al., 2008) that plasma persistently sub-corotates and appears to partially decouple from (i.e. slips further behind) the planetary rotation near the orbit of Titan (
20 RS) and beyond. The Cassini Plasma Spectrometer (CAPS) (Young et al., 2004; 2005) has been used to calculate plasma convection and other moments as well, and the results of Wilson et al. (2008) in the inner regions, McAndrews et al. (2009) and Thomsen et al. (2010) generally confirm these studies above and are in accord with the present work. The CAPS and INCA instruments do have their

M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

limitations, in their angular coverage for a given orientation of Cassini relative to the local flow direction, their sensitivity, and duty cycle, which affect their capability to perform flow measurements. For this reason, many of the observations from these instruments are complementary and important in forming as complete a picture as possible in determining the convection, composition, and spectral characteristics of Saturn's magnetospheric plasma. As has been determined by analysis of Pioneer data (Simpson et al., 1980; Thomsen et al., 1980), the studies above all confirm the predominance of rotation in the overall plasma convection within the orbit of Titan. In this effort, we extended the previous and ongoing work of Kane et al. (1992, 1993, 1995, 1998, 1999, 2008) by applying the techniques used to quantify the anisotropies of hot ions and calculate from them the velocity of plasma in the magnetosphere of Saturn. To understand dynamic processes in the magnetosphere of Saturn, it is important to know the topology of the system and to understand the extent of influence of the interplanetary medium on that topology. Thus processes such as reconnection may be internally driven by the planetary rotational energy or externally driven by the solar wind energy. In the case of Jupiter and Saturn, most of the plasma is contained within a rotating magnetodisk. With the radial transport of mass from interior sources, this mass causes increasing stress on the magnetic field lines until the plasma can no longer be contained, and is sometimes pinched off to form blobs of plasma called plasmoids. A steady state picture of such a topology (Vasyliunas, 1983) has been used to explain the observed release of plasma by the formation of plasmoids at Jupiter. These decoupled magnetic islands are ejected down the magnetotail away from the Sun. While the release of plasma by this Vasyliunas-cycle reconnection is internally driven, the Dungey cycle (Dungey, 1961) connects solar wind flux to a magnetosphere and releases magnetospheric plasma down the tail. Both types of reconnection could be operating at Jupiter and Saturn.

2. Instruments and data sets This effort utilized principally data from the Magnetospheric Imaging Instrument (MIMI) (Krimigis et al., 2004) on the Cassini spacecraft. The primary dataset consists of observations from the Ion and Neutral Camera (INCA), one of three instruments in the MIMI cluster. INCA has a very wide field of view which allows one to obtain anisotropy information even in stare mode, when the instrument is staring in a particular direction (and the spacecraft is not spinning). Spin mode data is the preferred mode due to its larger nearly all-sky coverage, however. INCA can provide compositional and energy information over a field of view of 1201  901 when staring and 1201  3601 when Cassini is spinning. The Low Energy Magnetospheric Measurement System (LEMMS), a second component of MIMI, measures ions from
30 keV to 410 MeV. A third MIMI dataset that is available is the compositional information of ions in the 7 keV/nuc to 3 MeV/nuc range from the Charge Energy Mass Spectrometer (CHEMS). Where appropriate (for example, to distinguish anisotropies ordered by pitch angle from those ordered by convection), the Cassini magnetometer (MAG) provided us with magnetic field vectors needed for the analysis (Dougherty et al., 2004). LEMMS provided scan data from Saturn orbit insertion (SOI) on Day 183, 2004, through Day 32, 2005, after which the scan stopped operating and LEMMS continued in fixed mode to the present. The INCA and CHEMS instruments (which can provide compositional information) have operated nearly continuously from SOI to the present. 2.1. Analytical approach The INCA detector is an ion and neutral camera so that to acquire ions the high voltage supplied to its charged particle deflection plates

Log Intensity (cm2 sr keV sec)-1

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Fig. 1. INCA time of flight channel activity during 2006 day 309. The oscillation seen during part of the interval is due to spacecraft spin. As the spacecraft spins, the detector spins with it and views different directions relative to the convection direction of the population. The intensities then vary due to the Compton–Getting effect. The lower energy channels have higher particle intensity, and thus appear at the top of the line plots. The amplitude of the oscillation (or, the anisotropy) is larger at lower energy, consistent with the trend expected of a particle population convecting by the detector. Slow variation of the averages can indicate time aliasing, thus we attempt to avoid periods with large variations.

is periodically switched off. The INCA detector is able to differentiate oxygen (the other primary species detected by INCA) from hydrogen ions via pulse height analysis as well as to differentiate energies within the species through time of flight in the energy range
5– 500 keV. In ion mode, it typically produces hydrogen and oxygen ion flux data arranged in a 16  16 rectangular pixel array that maps onto a sector of the sky 1201  901 in size. We thus have ion intensities at 256 positions in the sky for eight energy ranges (see Fig. 1 for details) and two species. At oblique angles of incidence, the energy loss of an incident ion penetrating the outer foil of the detector is more uncertain, so that the ions counted at the edges of the detector, the outer rows and columns of pixels, are not used in stare mode. When the instrument (by means of the rotation of the spacecraft) is in spin mode, we may use all columns [parallel to the spacecraft spin axis] of pixels. In this mode, each column of pixels fixed in space is sampled by all columns of the spinning detector pixels, so that the result is an array of counts in pixels fixed in space. We are able to produce a nearly all-sky map (3601  1201 with no coverage within 301 of the spin axis) of counts during spin mode, a mode suitable for anisotropy studies. One typical day of data is shown in Fig. 1. During this day [309 of year 2006] the high voltage was turned off from
hours 8 to 18, allowing the detector to count ions. Those counts were converted into intensities and are plotted for the 8 time of flight channels TOF 1–8. During part of the ion detection mode (
hours 9–13), the spacecraft was spinning, which caused the count rate and therefore the intensity to oscillate at the spin period. Data for our study was acquired during spin periods such as this one. The spin period is usually
1/2 h, so that time aliasing is present in the data. In our case, where we accumulate the data into four 901 (quarter of a spin) sectors, different sectors could be sampling different ion populations. We have chosen periods where the time variation of the average is as small as possible during the spin interval, thus minimizing but not eliminating aliasing from the data and subsequent analysis. Our methodology is similar to that developed previously by Kane et al. (1992, 1993, 1995, 1998, 1999, 2008) for Voyager LECP observations at Jupiter and in the heliosphere, Galileo EPD measurements within the magnetosphere of Jupiter, and recently for the INCA detector on Cassini. We begin with an a priori assumption as to the nature of the hot plasma distribution. In order to keep the analysis as simple as possible, we chose to use for this analysis a convected power law distribution of the form f ðvÞ ¼ C jv  Vj  2ðγ þ 1Þ

ð1Þ

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R~25.4 RS LT~01:39 L~26.4 LAT~10.90 0.6268

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Hour: Minute of Year2006 Day 309 UTC þ

Energy (keV) Fig. 2. Spectrum from LEMMS, INCA, and CHEMS data (H þ channels) fit to a κ distribution function (small triangles). Here the outer magnetospheric ions are sufficiently cold for the calculated temperature to fall below the range of energies detected. In the outer magnetosphere, INCA typically samples a power law spectrum of H þ ions. In this analysis, we use a convected power law distribution.

where C is a normalization constant, v is the particle velocity, V is the convection velocity (assumed parallel to Saturn's equatorial plane, i.e. a 2-D bulk velocity vector), and γ is the spectral index of the distribution function. Justification may be found upon examination of Fig. 2. The intensities from TOF channels are presented as a spectrum during the period 2006 day 309 hours 10–11. We note that the three detectors of the MIMI cluster (INCA, CHEMS, and LEMMS) agree quite well (there is no relative calibration correction). When the spectrum was fit to a convected κ distribution (Vasyliunas, 1971), isotropic in its rest frame, the resulting temperature was lower than the lowest energy channel available. In this case, we are primarily measuring the power law tail of the population (if it may be described by a cold κ distribution). Thus we used a power law distribution for this study. Once a model distribution function is assumed, we generate model counts for each pixel in each energy channel, using the efficiency and passbands of the INCA detector. The model and actual counts are then least squares fit, and the model distribution function parameters (C, V, and γ, V is a 2 component vector) are then varied until a least squares minimum is found. The lower energy channels sense most of the anisotropy generated by the convection of plasma, and are therefore more important in determining flow speeds. Thus we do not use log space in the minimization. We note here that this is a non-linear analysis, i.e. the full distribution of Eq. (1) is fitted to the data. The channel passbands were integrated by energy for accuracy when generating the model intensities and count arrays. All channels used and all pixels are simultaneously least squares fit to the model count arrays. The field of view of each pixel is small so that the pixel centerline is used to generate model counts in a particular pixel from the model intensity incident along that line. As an example of the intensity of ions that INCA is able to measure, Fig. 3 shows an intensity map as a function of energy and look direction. Each small box in the figure is pixelated into a 16  16 array. There is a box for each energy channel and time. Given that INCA has a
901  1201 field of view, each box is a quadrant of the sampled sky encompassing 901 of rotation and 1201 along the spin axis. The spacecraft orientation is [and always

Fig. 3. INCA H ion intensities. For each time period (column) 6 time of flight channel intensity maps are shown. Each small box contains intensities (logs shown aside color bars) in a 16  16 pixel array, viewing in a 1201 (vertical in the plot)  901 sector of the sky. During the period shown, the spacecraft is spinning with the  Z-axis pointing toward the Sun (down in the figure). A full four quadrant scan is completed in one spacecraft spin. The intensities oscillate in a period containing 4 columns (1 spin). Given that in 1 box the intensities are uniform along the vertical axis but oscillate along the horizontal (time) axis, and presupposing convection is the main source of anisotropy, then the convection vector should have a large component perpendicular to the spacecraft spin axis. The lines labeled “30”, “60”, etc., depict pitch angle contours of particles incident on the INCA sensor.

is for this study]  Z (in spacecraft coordinates) toward the Sun. On the plot,  Z points downward. In one spin, we are thus sampling all the space except for cones within 301 of the spin axis (sunward or anti-sunward). For clarity, we restate the following, described above. The quadrants represent intensities fixed in space. The spacecraft is spinning and each detector pixel is actually sampling a small band of the sky continually throughout the 901 rotation. There is onboard processing, however, that uses the efficiencies and geometric factors and ions counted for each detector pixel to generate sky pixel arrays fixed in space. Thus each spatial pixel in a quadrant is constructed based on counts accumulated in all 16 detector pixels in a given row, since during the 1/4 rotation all 16 detector pixels in a row will at various times sample a given spatial (fixed in space) pixel. In this case, inspection of the boxes in Fig. 3 reveals that there is only weak variation in the intensities as the viewing angle relative to the spin axis (vertical within each box in the plot) varies. Rather, there is a significant intensity variation in time (horizontal). This is produced by the rotation of the spacecraft. Given the usual spectral form (intensities increase with decreasing energy), intensity is enhanced when a pixel is sampling ions whose gyrocenter is approaching the spacecraft and diminished when receding (i.e. the Compton–Getting effect). Therefore this set of intensities is produced by a plasma having a bulk flow velocity with the primary component perpendicular to the spacecraft spin axis. When a subset of the data (2006:309:10:29–56) in Fig. 3 is fit with our procedure to a convected distribution, we find an azimuthal velocity of 180 km/s and an outward radial velocity of 75 km/s. At this position (
25 RS, 10.901 lat) the rigid speed would be 25.4 RS  9.76 km/s  RS  cos(10.901)¼243 km/s, yielding a fractional azimuthal speed of 180/243 ¼ 0.74. This event occurred during Rev 32 (Fig. 4) when the spacecraft was in the pre-dawn sector. Events such as this that have spin-generated anisotropies are common near local midnight, signifying a large azimuthal component to the flow. When the spacecraft is near dawn or dusk, one sees anisotropies that are significant in the vertical direction of

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Rev 8-14,16,18-27,121-123,125-128 Trajectory 20

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XSZS (RS) Fig. 4. Rev 32. The activity detected by the INCA instrument during hour 10, day 309 of 2006 (Figs. 1–3) occurred while the spacecraft was spinning and its [  Z] spin axis pointed toward the Sun (to the right in this figure). The coordinates used are SZS (see beginning of Section 3 for details). The magnetospheric local time is
01:38. The convection vector derived from fitting the observed H þ anisotropies during minutes 29–56 of hour 10 is shown as an arrow at the spacecraft location. This vector is 22.61 outward from the corotation direction and nearly perpendicular to the spacecraft spin axis. This is consistent with our assumption based on the intensities shown in Fig. 3 that the magnetospheric plasma velocity had a large component in the direction perpendicular to the spacecraft spin axis. Typical locations of the magnetopause and bow shock are shown for this and subsequent figures and unless specified were taken from Arridge et al. (2006) and Masters et al. (2008).

Fig. 3, roughly parallel to the spacecraft spin axis, which again is in the general direction of corotation. Quite frequently the situation is mixed, so that a full quantitative analysis of the energy dependent anisotropies must be done. However, events may also be analyzed qualitatively using the intensity anisotropy plots similar to Fig. 3, drawing conclusions about the convection using the reasoning presented with Figs. 3 and 4. For our quantitative analysis, the efficiencies were determined by assuming the plasma was not flowing during a period of time when the spacecraft was sufficiently near the nose of and within the magnetosphere. The spacecraft Z-axis (the spin axis) was pointed toward the Sun during this time. A flat fielding approach was used, where the array of efficiencies for each of the TOF channels available was adjusted to yield uniform intensities. It was clear during the analysis that the anisotropies were small here, and subsequent analysis confirms that this period was a good choice to perform the calibration. The responses of the highest and the lowest energy channels are not well determined and are therefore not used. We also have an internal calibration that compares hydrogen and oxygen ions. The two species have different bandpasses and will respond differently to convection due to their mass difference. We have assumed the same geometrical factors for both species. We are able to perform the analysis independently for both species and compare the results in some cases where oxygen channels contain sufficient counts and yield reliable results. In those cases, we have found good agreement between velocities derived using the two species. One significant point that deserves elucidation is that sampled results are similar when we used either a convected κ distribution or a convected power law distribution. The κ distribution is a generalized spectrum that reduces to a near power law when the energetic ion "temperature" is cold (Kane et al., 1992). In our analysis of the magnetosphere, we have fit sample data to both spectra and the resulting calculation produces similar convection velocities. We emphasize that the a priori choice of the distribution function is

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Xszs (RS) Fig. 5. Orbit trajectories in the SZS coordinate system (see beginning of Section 3 for details) and data coverage for this convection study. The markings along the orbits indicate where data has been analyzed. The shaded boundary is a model magnetopause (Kanani et al., 2010). The study spans the nightside local time.

not of primary importance in the determination of the primary parameters of interest (convection velocity, pressure, composition) as long as that function describes the data well (i.e. the least squares error is small). Though the κ distribution has been remarkable in its ability to describe well observations of energetic ions encountered in planetary magnetospheres, objections to its use as an empirically based function with no physical justification deserve consideration. Thus we realize that our choice in using such representations does not necessarily reveal analytically, for example, the nature of the acceleration and transport mechanism that produces the energetic ion population in its observed state.

3. Plasma velocity determinations For our study of nearly 1400 spins of the Cassini spacecraft, we used the orbits whose trajectories in the SZS X–Y plane are illustrated in Fig. 5. The SZS coordinate system is formed by the Saturn spin axis (ZSZS), the cross product of ZSZS and the Saturn— Sun vector (YSZS) and XSZS ¼YSZS  ZSZS. Essentially, the X–Y plane is the Saturn equatorial plane with X lying in the local noon meridional plane. The time spanned the period from 2005 day 143 to 2010 day 83, with data analyzed on 163 different days. All orbits analyzed are equatorial orbits, and all points were within the magnetosphere. 3.1. Azimuthal flow In our analysis, we assume the convection direction is parallel to the equatorial plane of Saturn. Thus we derive an azimuthal and

M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

0.6

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Since we compute flow in a plane parallel to the Saturn equatorial plane, we have radial speeds available as well and plot their values in Fig. 7. While there is considerable scatter in the results, with the plasma sometimes reversing direction toward radially inward, there is a distinct bias toward outward flow [and transport] of plasma. This is consistent with a source of plasma for the inner magnetodisk, with the ions being generated from neutrals by electron impact, photoionization, and other processes, and subsequently transported outward by centrifugal force as part of the flux tube interchange mechanism. The measured bulk flow is then the result of net outward mass transport, which at our time resolution (
34 min per spin period) averages out the motion of flux tubes [and their varying mass content] that are participating in interchange. While there is no significant radial dependence observed in the radial flow speed, there is a local time dependence. The plasma in the dusk-midnight quadrant is evidently being transported outward at a higher rate of speed than the average. We might expect such a result if the plasma emerging from the dayside in the dusk flank expands rapidly outward under centrifugal force due to reduced pressure from the magnetopause, although we do not find a peak at or near dusk. Given that the azimuthal speed of the plasma across the magnetotail does not change systematically, the plasma we are sampling is mostly coupled to the planet (but subcorotating). If significant quantities of plasma detached near dusk and moved in a nearly rectilinear motion, we should expect to see the azimuthal speed decrease with local time as the local time azimuth direction increasingly diverts from its direction at dusk. It is possible that we are not sufficiently distant to see detached plasma “blobs”, if the duskside detachment occurs well beyond Titan's orbit. Examination of Fig. 7 reveals that the averaged radial speed is constant (given there is a large degree of scatter) until local midnight. At this point, the averaged radial speeds begin to

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Distance from Spin Axis of Saturn (RS) Fig. 6. Azimuthal normalized plasma speeds at Saturn. Data from 1376 spins of the Cassini spacecraft from all local times yielded the azimuthal speeds, here shown as a fraction of the local rigid corotation speed. Fits to data with average deviations above 0.3 were eliminated from the plot (leaving 603 spins). This deviation is the average (over all pixels) of the fractional difference between actual and model count rate for each pixel, after all parameter values were determined by the least squares minimization of the pixel count rate differences.

radial component to the flow. Our analysis is limited by the sensitivity of the INCA detector, which tends to saturate in ion mode near the equatorial plane at
12 RS. We can, however, derive speeds deep into the magnetotail, where intensities are very low. We can also derive speeds when the spacecraft (and therefore the detector) is pointed in various positions in both stare and spin mode. For this study, though, we have chosen times when the spacecraft  Z-axis is pointed toward the Sun (which occurs most frequently). Our analysis produced the azimuthal plasma speeds shown in Fig. 6, normalized to the local rigid corotation speed. We note that our results indicate a recovery from the flow disruption seen in the inner magnetosphere by Voyagers 1 and 2 (e.g., Richardson, 1998) and modeled by Saur et al. (2004). In their model, radial transport, ion-neutral friction, and ion pickup all work to produce dips in azimuthal speed just outside the orbits of Dione (6.4 RS) and Rhea (8.8 RS). The steadily increasing departure from rigid corotation seen in Fig. 6 has been detected at Jupiter as well using analysis of Voyager 2 LECP measurements (Kane et al., 1992, 1995) and Galileo EPD measurements (Kane et al., 1999). In the outer magnetosphere, we would expect ion-neutral friction and ion pickup to be less important than radial mass transport. In the corotation dominated magnetospheres of Jupiter and Saturn, magnetodisk plasma from the dayside expands into the nightside region at the dusk flank, released from the pressure of the dayside magnetopause (albeit not into a vacuum). The effectiveness of radial transport in disrupting the flow (relative to the ever increasing rigid corotation value, essentially causing the plasma to lag local rigid corotation speeds) should be important. As the plasma subsequently moves dawnward across the magnetotail and outward, the ionosphere and magnetic field lines will attempt to enforce corotation, within the bounds of the available ionospheric conductivity. The results indicate that, as expected from work at Jupiter (Hill, 1979), corotation enforcement becomes increasingly difficult with increasing distance from Saturn. Our results are generally in accord with those derived from analysis of cold plasma moments by CAPS both in the solar wind prior to orbital insertion (Crary, private communication, 2005) and in the magnetosphere (McAndrews et al., 2009; Thomsen et al., 2010). We note that in the present study there was no specific local time effect found in the nightside azimuthal flow pattern. The values shown in Fig. 6 have a large degree of scatter. There are points in the 13–26 RS range that indicate superrotation. While

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some are outliers and not well fit, others could indicate fast flows possibly associated with dynamic effects in the magnetosphere such as dipolarization after plasmoid release. Other points are indicating little or no corotational motion, and could be disruptions in the flow induced by injection events, plasma sheet flapping (causing the spacecraft to sample higher latitude noncorotating plasma), or other dynamic events.

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10–60 RS) from INCA measurements. The speeds are centered at about 60 km/s pre-midnight but decline post-midnight, shifting to radially inward as dawn is approached. Averages binned in 1 h local time increments and their standard deviations are shown (black dots and error bars). Some of the pre-midnight plasma sampled, given its trajectory, would likely escape the magnetosphere.

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decline until, near the dawn flank, the speeds are predominantly negative (the plasma is moving radially inward). We believe that this trend toward radially inward speeds results from magnetopause pressure at the dawn flank. Given the significant radial speeds, the outermost magnetospheric plasma sampled in the pre-midnight region in this study is likely to escape the system (also, see McAndrews et al., 2009). However, the dawnside plasma sampled (that subset which is moving radially inward) is plasma likely returning for another traversal of the dayside. There are significant exceptions, where plasma is seen moving radially inward with large velocities. Some of these

may be associated with transient events such as reconnection or injections. Give the radial speeds in Fig. 7, with the average prior to the duskside transition at
60 km/s, and the model of Bagenal and Delamere (2011), this would imply a high mass transport rate at Saturn near the 250 kg/s line (in their Fig. 10). 3.3. Duskside Beginning at
2009 day 285 (during Rev 119), the Cassini spacecraft began to explore the low latitude dusk magnetosphere, magnetopause, and magnetosheath. The INCA H þ (13.5–149 keV)

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Fig. 8. Flow velocities of magnetospheric plasma in the dusk-midnight quadrant at Saturn. Velocity vectors averaged over several hours are shown at daily intervals (arrows), with line segments denoting full (rigid) corotation. The flow is assumed to be parallel to the equatorial plane. Results from two of the orbits analyzed reveal that the plasma flow is primarily azimuthal and sub-corotating (but with a significant radial component) even at large radial distances from Saturn, where one might expect flows parallel to the magnetopause.

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Energy (keV) Fig. 9. Oxygen ion (O þ ) spectra measured by INCA (left panel) at the locations shown (arrows) in the Rev 49 trajectory (right panel). Two hydrogen ion (H þ ) spectra are also shown for comparison. The H þ spectra are typically power law (as shown) in the outer magnetosphere. During Rev 49, the INCA detector measured O þ ion intensities in the magnetosphere, magnetosheath, and solar wind. There are significant abundances of O þ in the dusk magnetosheath, originating from Saturn's magnetosphere via diffusion, Kelvin–Helmholtz instability-driven transport at the magnetopause boundary, or other dayside reconnective processes.

M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

channel measurements were analyzed using the techniques described above. Results were averaged in daily intervals (several hours of coverage per day in spin mode). The velocity vectors in the equatorial (SZS) plane are shown in Fig. 8 next to vectors pointing in the azimuthal direction scaled to the local rigid corotation speed. Despite the proximity to the magnetopause and the distance from Saturn, plasma is still apparently bound to Saturn and rotates around the planet, with some outward flow evident, but with no persistent duskside flow disruption such as that reported at Jupiter (Krupp et al., 2001). During the inbound portion of Rev 49 (Fig. 9), Cassini passed through a large region of solar wind, passed completely through the magnetosheath and entered the magnetosphere. The INCA measured spectra in Fig. 9 showed very small fluxes of O þ (possibly H þ contamination at these very low levels) in the solar wind, but much more intense flux in the magnetosheath region. Some of the intensities rivaled those in the magnetosphere. It is likely that much of the O þ in the magnetosheath had been within the magnetosphere (Sergis et al., 2012) and crossed the magnetopause upstream in the post-noon quadrant, and that, given the intensities, this transport is a

Log Intensity (cm2 sr keV sec)-1

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-2 -3 286.0

286.5

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288.5

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Fig. 10. Oxygen and hydrogen intensities. Major intensity transitions are INCA high voltage switching. The lower intensity data are neutrals. While Cassini was executing Rev 16, the dawnside magnetopause was crossed (outbound) during day 288 of 2006. Peak intensities of O þ in the magnetosheath rival peaks (horizontal line in figure) inside the magnetosphere.

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significant factor in the escape of magnetospheric plasma (Krimigis et al., 2009). 3.4. Dawnside The Cassini spacecraft began exploration of the dawnside of the magnetosphere early in the mission, beginning with orbital insertion. After Rev 16 it began to explore the nightside dawn flank of the magnetosphere, probing deeper into the magnetotail and culminating in Rev 22–23. We have examined the anisotropy and rate data during this interval in conjunction with the magnetic field data and magnetopause crossings derived from analysis of low energy plasma data (Hansen, personal communication, 2009). This region is of particular interest since flows can oppose each other at this boundary, where magnetosheath flows are tailward and the corotation direction here is mainly sunward. Where large velocity shear is present, Kelvin– Helmholtz instability can generate waves which travel along the magnetopause and can produce plasma exchange between the magnetosphere and magnetosheath in vorticies (Masters et al., 2012, 2009; Desroche et al., 2013; Delamere et al., 2013; Wilson et al., 2012). We note with interest that evidence (from the studies above) is accumulating that shows the KH instability can operate equally well in both the pre-noon and post-noon sectors at Saturn. Beginning with Rev 16, we examined the anisotropies for flow signatures near the magnetopause. Some of the anisotropies found during these orbits were either not well ordered or were not available in ion mode near the transition to/from the magnetosheath. After crossing the magnetopause outbound, a significant oxygen flux was typically found, streaming tailward in the magnetosheath flow. In Fig. 10 below, we show the intensities of O þ and H þ ions within two time-of-flight channels measured by INCA. Peak O þ intensity enhancements in the sheath rivaled those in the magnetosphere. Oxygen ions are therefore escaping the magnetosphere upstream of this position. Their source could be from Kelvin–Helmholtz driven transport, other forms of diffusion, or dayside reconnection that populates the upstream solar wind. This enhancement was also seen at the dusk flank as noted above, and thus together would comprise a significant escape route for magnetospheric ions. Oxygen spectra were obtained in the magnetosphere, sheath, and solar wind during Rev 16 and 17 at the locations shown in Fig. 11. The sheath is populated with O þ ions far above solar wind

1.0E+02

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0

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1.0E+03

-20

-10

0

10

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30

XSZS (RS)

Fig. 11. Oxygen ion (O þ ) spectra measured by INCA (left panel) obtained at times (arrows, right panel) during Rev 16 (right panel). During Rev 16, the INCA detector measured O þ ion intensities in the magnetosphere, magnetosheath, and solar wind. There are significant intensities of O þ in the magnetosheath, evidently leaked from Saturn's magnetosphere upstream of the local position via diffusion, the Kelvin–Helmholtz instability at the magnetopause boundary, or other dayside reconnective processes, all in the pre-noon sector.

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M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

levels seen later (day 297 is actually at the beginning of Rev 17, but that orbit is not shown for clarity). The green arrow is the approximate location at which the solar wind spectrum was obtained. Thus the magnetosheath on both flanks is serving as an exhaust for magnetospheric ions. The detection of such ions sunward of the terminator is an indication that the dayside magnetopause is one upstream source for these ions. The spacecraft executed a more distant sweep of the dawn magnetotail flank during Rev 22 and 23 and spent a considerable amount of time in the vicinity of the magnetopause. Elements of Rev 22 and 23 appear in Fig. 12, with the red arrow pointing at 2006, day 105 where, in the vicinity of the magnetopause, tailward

Rev 22-23 2006 90-130

0

YSZS (RS)

-10 -20 -30 -40

flowing ions were detected inside the magnetosphere (Fig. 13). It is clear that plasma in the magnetotail crossing from dusk to dawn will, at some distance from Saturn, no longer return with the corotational flow to the dayside. Some of that plasma will encounter the dawnside magnetosheath flow and be entrained, to flow down the dawn flank and out of the Saturn system. This flow will accumulate more plasma from cross-tail flow at greater distances downstream. That may account for our lack of detection of similar flows in earlier orbits. We note that, even at such a great distance from Saturn, sunward (or corotational) flows were detected. Given the observed cross-tail flows well beyond Titan's orbit on the nightside, we would expect that the plasma which could not return to the dayside would form a region of tailward flow that would increase in prominence further downstream (see discussion section below). The thumbnail images in Fig. 13 are actually a collection of 16  16 pixel intensities measured by INCA. Each small box is an array of intensities in a given channel and time. The channel energies increase upward (by rows), and each column is a detection within a given time interval. The orientation of the spacecraft is  Z to Sun (after the third frame), where Z is the usual spacecraft spin axis. The images are oriented so that the bottom of each small box is sensing particles from the  Z direction, essentially from the Sun. There is a transition from stare mode to spin mode at 21:11 and a transition to  Z to Sun during the first three frames. During both stare and spin mode, there is a strong anisotropy produced by a more intense particle flux coming from the –Z direction (from the Sun), essentially indicating tailward flow.

-50 -60

-50

-40

-30

-20

-10

0

4. Discussion

XSZS (RS) Fig. 12. Dawn magnetotail encounter during Rev 22–23. During this distant sweep of the dawn flank of Saturn's magnetotail, INCA detected a variety of flow signatures. Cassini crossed the magnetopause outbound during 2006, day 101 and inbound near the beginning of day 105 (arrow). Later the same day, tailward flowing ions were detected (see Fig. 13).

Our quantitative study reveals that nightside convection is predominantly azimuthal to
50 RS on the nightside, but with a significant radial outflow component. Even the outer magnetospheric plasma sampled is only partially decoupled from the ionosphere of Saturn. There is more ionospheric slippage of the field-aligned

R~63.3 RS LT~03:24 L~63.3 -0.2160 Oxygen 231-332 keV

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169-231 keV

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To Sun

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20:41

20:47

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20:59

21:05

21:11

21:17

21:24

21:31

Hour: Minute of Year 2006 Day 105 UTC Fig. 13. Tailward flowing O þ ions in the dawn magnetotail. The INCA instrument detected a region of tailward flow at the flank of the dawn magnetotail at
63 RS. Within each box are a 16  16 pixel array with intensities (logs shown at right) deduced from measured count rates, the total field of view being 1201 (vertical)  901 (horizontal). The intensities are enhanced within most of the lower energy boxes toward the bottom of each box, where the pixels point toward the sun (the  Z direction in spacecraft coordinates, Z is the spin axis during part of this period).

M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

currents that enforce corotation (Hill, 1979) at larger distances, but this distant nightside plasma is still [loosely] bound to Saturn and continues to rotate with the magnetospheric plasma sheet until it

Fig. 14. Voyagers 1 and 2 trajectories at Jupiter. In this figure [Sun is at left], adapted from Figure 23 of Krimigis et al. (1981), the Voyagers 1 and 2 trajectories are plotted in ecliptic coordinates and first order anisotropy vectors (line segments) projected onto the ecliptic plane are shown as a representation of the convection direction. The nominal magnetopause locations are represented by the dashed curves in the figure. Voyager 2 crossed the dawn magnetopause further downstream than did Voyager 1. Voyager 2 encountered a region where anisotropies are directed anti-sunward (highlighted), both inside and outside the magnetopause. Particle populations were similar to those seen in the magnetodisk, presumed to be their source.

9

approaches the dawn magnetopause. Given the presence of a radial outflow in much of the nightside plasma sheet, we conclude that some of the nightside plasma beyond
20 RS cannot return to the dayside, in accordance with the conclusion of McAndrews et al. (2009). Given the radial flow dependence seen in our study (Fig. 7), some of the plasma that expanded well beyond the orbit of Titan in the nightside during passage across the tail is transported inward near the dawn flank. However, a significant amount of nightside plasma must fully decouple at the dawn flank and is entrained into the magnetosheath flow down the dawn flank of Saturn's magnetosphere. Such a scenario is similar to the Voyager 2 detection of a distant dawnside boundary layer of tailward streaming plasma in Jupiter's magnetotail (Fig. 14), magnetospheric in origin (Krimigis et al., 1981). We have measured the intensity, spectra, velocity, and composition of this [Jovian] structure and have found it to be very similar to that encountered in Jovian plasma sheet crossings but having convection speeds typical of magnetosheath ions. The Voyager 2 spacecraft encountered Jupiter in 1979 and subsequently moved deep down the dawnside magnetotail, encountering the dawn magnetopause beginning at
200 RJ. An examination of Fig. 15 reveals that after closest approach during day 222 of 1979 the spacecraft periodically encountered high fluxes associated with the tilted current sheet of Jupiter, where motion of the current sheet past the spacecraft was assumed to be enforcetropies is predominantly azimuthal during the periodic encounters with the current sheet plasma (Krimigis et al., 1981; Carbary et al., 1981; Kane et al., 1995; Krupp et al., 2001), but begins to transition after
140 RJ and eventually becomes predominantly tailward as the magnetopause is approached. Tailward flow continues on both sides of the dawn magnetopause.

Fig. 15. Voyager 2 encounter with the dawnside boundary region (horizontal line) at Jupiter, adapted from Figure 13 of Krimigis et al. (1981). Ion distributions were very similar to those in the plasma sheet encounters. The existence of this substantial region of tailward flowing magnetospheric plasma provides a means for plasma produced in the inner region (at Io) to escape the system. This observation motivated Cheng and Krimigis (1989) to propose an asymmetric convection model for Jupiter. The corresponding existence of tailward flowing magnetospheric ions at Saturn's dawn flank and substantial O þ ions in the sheath motivates us to propose that Jupiter and Saturn have similar configurations and have significant losses of ions along the magnetopause boundary.

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This observation is consistent with the existence of a dawn boundary layer of tailward flowing plasma, likely entrained into the magnetosheath flow by a viscous interaction (including enhanced transport at Kelvin–Helmholtz instability structures) at the magnetopause. The aperiodic fluxes are explained by a latitudinal broadening of the intense plasma region near the magnetopause beyond that of the outer magnetodisk. Here particles are no longer constrained by closed and pinched (current sheet-like) field lines; the plasma is decoupled from the planet. The tilted rotating current sheet would interact with the dawnside magnetopause in the range of latitudes subtended by the current sheet as the planet completes a rotation, essentially broadening in latitude magnetospheric ions at the boundary layer. Though in our observations at Saturn we detected tailward flowing magnetospheric ions in the dawn flank, the feature is not prominent as at Jupiter. It is possible that our observations were not far enough downstream at the dawn flank of Saturn's magnetosphere to collect enough particle flux from ions at large distances crossing the tail, then turning tailward and escaping the system. With this observation of Voyager 2, Cheng and Krimigis (1989) proposed that the Jovian convection system is asymmetric (Fig. 16), with an Earth-like tail comprised mainly of solar wind ions flowing anti-sunward interacting with magnetospheric cross-tail flow. At the dawn flank, cross-tail flowing ions then interact with magnetosheath ions and are entrained into the flow, creating a “kink” in the flow. Subsequent observations (e.g., Krupp et al., 2001) during the Galileo mission detected azimuthally convecting ions of magnetospheric origin as far as 142 RJ down the magnetotail. Thus convection in the magnetotail is not generally planetward, with the possible exception of occasional reconnection. Nevertheless, the explanation in Cheng and Krimigis (1989) of a boundary layer in the dawn flank produced by entrainment of cross-tail flow is still compelling and consistent with some observations at Jupiter and Saturn. Our adaptation is to move the “kink” where magnetospheric plasma turns and becomes entrained into the magnetosheath flow further downstream, so that plasma well beyond the standoff distance in the nightside would be entrained by the dawn magnetosheath, while less distant plasma would be forced inward by the presence of the dawnside magnetopause and rotate again around the dayside. Undoubtedly some of the return flow (to the dayside), particularly that near the magnetopause, would encounter instability regions and may be transported into the magnetosheath to flow down either flank. In fact, this transport must occur at the dayside magnetosheath to produce the magnetospheric ion flux Jovian Ions in Dawn Magnetosheath and Magnetospheric Wind Dawn Magnetospheric Wind Boundary Layer

Voyager 2 Turbulent

Cross-tail Flow Solar Wind Earth-like Boundary Layer Dusk Fig. 16. Asymmetric model of Jovian convection. After the Voyager encounters with Jupiter, an asymmetric model of the magnetosphere was proposed by Cheng and Krimigis (1989). Their model shown here, adapted from Figure 3 of their work, was motivated by Voyager observations of corotational flow in the dawn magnetotail, followed by tailward flow of magnetospheric ions on approach to the magnetopause and during a series of encounters with that boundary.

seen in the magnetosheath. Therefore we view the magnetopause boundary as a continuously [in time and space] “leaky” boundary where significant quantities of ions escape from Saturn (and Jupiter). At Saturn, we presented evidence (Fig. 13) for tailward flowing plasma in the distant dawn flank. The Saturn magnetosphere differs in an important way from that at Jupiter. At Saturn, the magnetic field tilt is nearly zero. Thus the dawnside boundary region where tailward flow exists could be more limited in latitude than that at Jupiter, since the current sheet where transport into the magnetosheath flow occurs would not “flap” with the dipole tilt. This could make detection of the region less likely, particularly since Saturn's spin axis is tilted from the ecliptic and the plasma sheet shape is distorted by the solar wind away from the equatorial plane of Saturn and hence the spacecraft location during the observation in question. However, given the mystery of strong periodicities in Saturn plasma fluxes, magnetic field strength, SKR, etc., Khurana et al. (2009) have proposed that the solar wind coupled with planetary rotation drives the plasma sheet, containing an inner magnetic anomaly, into oscillation, producing periodicities that mimic those produced by the Jovian current sheet tilt. This could negate the lack of current sheet tilt at Saturn and broaden the current sheet-magnetopause interaction region at Saturn, although the relative effectiveness of this [compared to a magnetic dipole tilt] is uncertain. There is also, in the particle data from INCA, a possible indication that during one orbit an outer edge of the magnetodisk was traversed. As shown in Fig. 17, Cassini was in a distant region in the midnight-dawn quadrant during a
12 day period in 2006 and detected a dropout in oxygen and hydrogen ions in the INCA detector. The dropout occurred over many days so that any periodic current sheet oscillation would have ample time to manifest itself as an increase in ion intensity. Examining Fig. 15, there is a dropout in periodic enhancements [seen more clearly in the medium ion trace] during the Voyager 2 encounter with Jupiter at
140 RJ (at the highlighted boundary). When scaled from the Jovian system at the nominal standoff distance ratio of 60/25 ¼2.4), this distance corresponds to 140/2.4 ¼58 RS. The dropout seen in Fig. 17 occurs between 67.8 RS outbound and 56.4 RS inbound. Knowledge of the plasma convection in the outer magnetosphere of Saturn is particularly important for exploring the signatures of the reconnection process in a rotation-dominated magnetosphere (e.g., Cowley et al., 2005; Badman et al., 2005). This process, which may be dominated by internal Vasyliunas-type reconnection (Vasyliunas, 1983) was introduced as a steady state global structure of the Jovian magnetosphere but may in fact exist only as part of a time dependent sporadic process in corotation dominated magnetospheres (Hill et al., 2008; Kivelson and Southwood, 2005). Our global convection study places important constraints on the nature of such processes operating at Saturn. For example, our observations (and in accord with similar findings by McAndrews et al. (2009) and Thomsen et al. (2010)) of a persistent sub-corotational flow to distances well beyond Titan's orbit, with the azimuthal component predominant, would suggest that disruption of the flow generally associated with large scale reconnection and subsequent repolarization is limited. The location of such reconnection could be further down the tail, well beyond our observations, but Fig. 17 suggests there may not be sufficient mass flux at unexplored distant regions to create a significant sink for magnetospheric particles to escape. Based on our observations and calculations in both the Jovian and Saturnian magnetosphere and other sources, the global convection pattern at both planets can be represented by Fig. 18. Solar wind responding to magnetic and particle pressure near the respective planet forms a bow shock and magnetosheath. Magnetosheath flow

M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

drives magnetospheric plasma [that crosses the magnetopause all along the boundary] downstream. Shear instability and other processes transport plasma across the boundary. Plasma typically moves outward, averaged over the small scale flux tube interchange mechanism operating at both planets, as it rotates about the planet, greatly expanding (Fig. 18, Regions II and III) in the nightside. Some of the

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3

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164

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168

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172

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Fig. 17. INCA derived intensities of hydrogen and oxygen ions during 2006. During the 12 day period bounded by the highlighting in the figure, particle fluxes where unusually low, consistent with entry into a lobe-like environment beyond the magnetodisk.

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plasma well beyond the standoff distance in the nightside (Region III) approaches the dawn flank and is turned tailward by the influence of the magnetosheath flow. A boundary layer is created where tailward flowing plasma inside the magnetopause is driven by the magnetosheath flow downstream (Kane et al., 2009). This plasma undergoes small scale reconnection (Delamere and Bagenal, 2010), and is subsequently lost down the flanks of the magnetotail. Thus the primary interaction with the solar wind leading to plasma escape is viscous (particle driven) rather than driven by large scale magnetic reconnection such as the Dungey cycle in the high-β plasmas in the Jovian (Mauk et al., 2004) and Saturnian (Sergis et al., 2009) magnetodisk. Other plasma not as distant (Region II) will contract at the dawn flank and rotate to the dayside. Internal reconnection and plasmoid formation (Vasyliunas, 1983) does occur in the magnetotail, but it is not the dominant release mechanism for magnetospheric plasma at Jupiter and Saturn (Bagenal and Delamere, 2011). The conceptual model presented in Fig. 18 is scaled to the magnetopause standoff distance and applies to both Saturn and Jupiter (in its compressed state). This scaling is feasible by considering the effect of the magnetopause on plasma in Region II and III that has crossed the magnetotail. The dawnside magnetopause, magnetosheath [via viscous interaction] and the planetary rotation [via the ionosphere] exert stresses that essentially bifurcate the nightside flow beyond the standoff distance into two pieces: that which returns to the dayside and tailward driven flow. The magnetopause position is important for determining where distant nightside plasma is transported, thus it is a logical quantity to use as a scaling factor in the convection pattern. Quantitatively, we have also discovered that such a scaling factor relates trends in the speed of plasma and the temperature of hot ions seen at Jupiter and Saturn (addressed in another work). We note that in this model the magnetodisk has a discrete outer edge in the nightside that would limit magnetodisk plasma available for large scale reconnection and mass loss beyond this edge in the magnetotail.

5. Summary I

II III

Fig. 18. Empirical convection model for Jupiter and Saturn. Based on the analysis described herein and previous observations and analytical work at Jupiter during the Voyager and Galileo missions, this empirically derived model contains an inner rotating core extending roughly to the magnetopause standoff distance (Region I), an intermediate more distant plasma on the nightside (Region II, the velocities peak here) which returns to the dayside, and more distant nightside plasma (Region III) that is turned and entrained downstream by the dawn magnetopause and magnetosheath flow, forming a low latitude boundary layer (purple). Sporadic reconnection in the tail (red  ) will pinch off blobs of plasma that will move across and/or tailward. Plasma diffuses into the magnetosheath (red arrows), assisted by the Kelvin–Helmholtz instability at the magnetopause boundary, to populate this region with magnetospheric plasma. That diffusion is enhanced on the dawnside (dark red arrows) by small scale reconnective processes due to the large velocity shear and the presence of plasma propagating dawnward toward the magnetopause boundary there. The magnetospheres are “leaky” at the magnetopause boundary and small scale reconnection (black  ) allows plasma to escape down the dawn flank and in the magnetosheath as well as in the upstream region. (For interpretation of the references to color in this figure caption, the reader is referred to the web version of this article.)

We have used the anisotropies of hot ions measured by the INCA detector to quantify plasma convection in Saturn's magnetodisk. The plasma azimuthal convection speed, measured relative to the local corotation speed, decreases with radial distance and lags the planetary rotation rate. Unlike Jupiter, there are not significant local time asymmetries. There is a distinct radially outward component in the convection in the pre-midnight sector, but in the post-midnight sector the pattern gradually reverses to radially inward. The magnetosheath contains significant intensities of O þ from the magnetosphere. The particles are transported across the dayside magnetopause into the magnetosheath, assisted by the Kelvin–Helmholtz instability. This plasma is a significant sink of particles to be lost from the magnetosphere. We do not see evidence in the convection pattern for steady state reconnection in Saturn's magnetotail. Plasma in the outer regions of Saturn's nightside magnetodisk will likely interact with magnetosheath ions to be entrained into their flow and exhausted down the dawn flank of the magnetotail. When the magnetodisks of Jupiter and Saturn are scaled by their respective standoff distances, similar convection patterns describe the plasma transport at each planetary magnetosphere.

Acknowledgments The authors would like to thank Martha Kusterer at JHU/APL for her tireless efforts in providing the INCA data interface and

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M. Kane et al. / Planetary and Space Science 91 (2014) 1–13

1.0E+08

H+ Scalar C

1.0E+07 1.0E+06 1.0E+05 1.0E+04 1.0E+03 1.0E+02 5

10

15

20

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30

35

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45

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Distance from Spin Axis of Saturn (RS) Fig. A1. Intensity Scaling Factor. C (Eq. (1)) is shown vs. radial distance.

4.5

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Distance from Spin Axis of Saturn (RS) Fig. A2. Spectral Index. The power law slope γ (Eq. (1)) is shown vs. radial distance. The trend is toward softer spectra in the outer magnetosphere.

assisting with numerous computer and software issues. This work was supported in part by NASA grant NNX10AF18G to Harford Research Institute and by contract NAS5‐97271 between NASA Goddard Space flight Center and JHU/APL.

Appendix A Our large scale analysis of the middle to outer magnetosphere of Saturn was carried out using a forward model assuming the hydrogen plasma could be adequately described by a convected power law (Eq. (1)). There are four free parameters in the model: two convection directions in a plane parallel to Saturn's spin equator, a normalization constant, and a power law index representing the slope of the function in log–log space. For completeness we have added the normalization constant and slope, both plotted as a function of the distance in Saturn radii from the spin axis (Figs. A1 and A2). In both plots, the same subset was excluded (where the average deviation between model and actual count rates in each pixel and five TOF channels of the INCA detector exceeded 0.30 in a given spin of the spacecraft, also see Fig. 6). References Arridge, C.S., et al., 2006. Modeling the size and shape of Saturn’s magnetopause with variable dynamic pressure. J. Geophys. Res. 111, A11227, http://dx.doi.org/ 10.1029/2005JA01574. Badman, S.V., Bunce, E.J., Clarke, J.T., Cowley, S.W.H., Gerard, J.-C., Grodent, D., Milan, S.E., 2005. Open flux estimates in Saturn's magnetosphere during the January 2004 Cassini-HST campaign, and implications for reconnection rates. J. Geophys. Res. 110, A11216, http://dx.doi.org/10.1029/2005JA011240. Bagenal, F., Delamere, P.A., 2011. Flow of mass and energy in the magnetospheres of Jupiter and Saturn. J. Geophys. Res. 116, A05209, http://dx.doi.org/10.1029/ 2010JA016294. Bridge, H.S., et al., 1982. Plasma observations near Saturn: initial results from Voyager 2. Science 215, 563–570.

Bridge, H.S., et al., 1981. Plasma observations near Saturn: initial results from Voyager 1. Science 212, 217–224. Carbary, J.F., Mauk, B.H., Krimigis, S.M., 1983. Corotation anisotropies in Saturn's magnetosphere. J. Geophys. Res. 88, 8937–8946. Carbary, J.F., Krimigis, S.M., Keath, E.P., Gloeckler, G., Axford, W.I., Armstrong, T.P., 1981. Ion anisotropies in the outer Jovian magnetosphere. J. Geophys. Res. 86, 82. Cheng, A.F., Krimigis, S.M., 1989. A model of global convection in Jupiter's magnetosphere. J. Geophys. Res. 94, 12003–12008, http://dx.doi.org/10.1029/ JA094iA09p12003. Cowley, S.W.H., Badman, S.V., Bunce, E.J., Clarke, J.T., Gerard, J.-C., Grodent, D., Jackman, C.M., Milan, S.E., Yeoman, T.K., 2005. Reconnection in a rotationdominated magnetosphere and its relation to Saturn's auroral dynamics. J. Geophys. Res. 110, A0221, http://dx.doi.org/10.1029/2004JA010796. Delamere, P.A., Wilson, R.J., Eriksson, S., Bagenal, F., 2013. 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