Reprint - Earth & Planetary Sciences - University of California, Santa ...

1 Saturn after Cassini/Huygens
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Revised Version April 10, 2009
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Icy Satellites: Geological Evolution and Surface Processes
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Ralf Jaumann 1,2,* , Roger N. Clark 3 , Francis Nimmo 4 , Amanda R. Hendrix 5 , Bonnie J. Buratti 5 6
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Tilmann Denk 2 , Jeff M. Moore 6 , Paul M. Schenk 7 , Steve J. Ostro ✝5 and Ralf Srama 8
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DLR, Institute of Planetary Research. Rutherfordstrasse 2, 12489 Berlin, Germany;
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Freie Universität Berlin, Institute of Geological Sciences, Malteserstr. 74-100, 12249 Berlin, Germany;
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U.S. Geological Survey, Denver Federal Center, Denver CO 80225;
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Dept. Earth & Planetary Sciences, University of California Santa Cruz, Santa Cruz, CA 95064, USA;
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Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA;
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NASA Ames Research Center, MS 245-3 Moffett Field, CA 94035-1000, USA;
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Lunar and Planetary Institute Houston TX 77058, USA;
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Max Planck Institut für Kernphysik, 69117, Heidelberg, Germany.
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Corresponding author (Fax :+493067055 402 ; Email address: ralf.jaumann@dlr.de
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deceased
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Content
Abstract
23 21.1 Introduction
24 21.2 Cassini’s Exploration of Saturn’s Icy Satellites
25 21.3. Morphology, Geology and Topography
26 21.4 Composition and Alteration of Surface Materials
27 21.5 Constraints on top-meter Structure and Composition by Radar
28 21.6 Geologic Evolution
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21.7 Conclusions
31 Abstract
32 The sizes of the Saturnian icy satellites range from ~1500 km in diameter (Rhea) to ~20 km
33 (Calypso), and even smaller 'rocks' of only a kilometer in diameter are common in the system. All
34 these bodies exhibit remarkable, unique features and unexpected diversity. In this chapter, we will
35 mostly focus on the 'medium-sized icy objects' Mimas, Tethys, Dione, Rhea, Iapetus, Phoebe and
36 Hyperion, and consider small objects only where appropriate, whereas Titan and Enceladus will be
37 described in separate chapters. Mimas and Tethys show impact craters caused by bodies that were
38 almost large enough to break them apart. Iapetus is unique in the Saturnian system because of its
39 extreme global brightness dichotomy. Tectonic activity varies widely - from inactive Mimas
40 through extensional terrains on Rhea and Dione to the current cryovolcanic eruptions on Enceladus
41 – and is not necessarily correlated with predicted tidal stresses. Likely sources of stress include
42 impacts, despinning, reorientation and volume changes. Accretion of dark material originating from
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outside the Saturnian system may explain the surface contamination that prevails in the whole
satellite system, while coating by Saturn’s E-ring particles brightens the inner satellites.
45 So far, among the surprising Cassini discoveries are the volcanic activity on Enceladus, the sponge-
46 like appearance of Hyperion and the equatorial ridge on Iapetus - unique features in the solar
47 system. The bright-ray system on Rhea was caused by a relatively recent medium impact which
48 formed a ~40 km crater at 12°S latitude, 112°W longitude, while the wispy streaks on Dione and
49 Rhea are of tectonic origin. Compositional mapping shows that the dark material on Iapetus is
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composed of organics, CO2 mixed with H2O ice, and metallic iron, and also exhibits possible
signatures of ammonia, bound water, H2 or OH-bearing minerals, and a number of as-yet
unidentified substances. The spatial pattern, Rayleigh scattering effect, and spectral properties
53 argue that the dark material on Iapetus is only a thin coating on its surface. Radar data indicate that
54 the thickness of the dark layers can be no more than a few decimeters; this is also consistent with
55 the discovery of small bright-ray and bright-floor craters within the dark terrain. Moreover, several
56 spectral features of the dark material match those seen on Phoebe, Iapetus, Hyperion, Dione and
57 Epimetheus as well as in the F-ring and the Cassini Division, implying that throughout the
58 Saturnian system. All dark material appears to have a high content of metallic iron and a small
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content of nano-phase hematite. However, the complete composition of the dark material is still
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60 unresolved, and additional laboratory work is required. As previously concluded for Phoebe, the dark
61 material appears to have originated external to the Saturnian system.
62 The icy satellites of Saturn offer an unrivalled natural laboratory for understanding the geological
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diversity of different-sized icy satellites and their interactions within a complex planetary system.
65 21.1 Introduction
66 Covering a wide range of diameters, the satellites of Saturn can be subdivided into three major
67 classes: ‘icy rocks,’ with radii < 10 km, 'small satellites' with radii

Siarnaq 17531 895.53 0.2960 46.002 20 2.3 0.0026 20.1R 0.06
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77 Tab. 1 Physical and orbital characteristics of the small (

134 suggested endogenic activities. However, the nature of these activities was not well understood. For
135 a more detailed discussion of our knowledge before Cassini, see Dougherty et al. (2009) and Orton
136 et al. (2009). In terms of geology, the Cassini mission was designed to provide significantly better
137 area coverage, spatial resolution, and spectral range, resulting in numerous new discoveries. A
138 detailed description of the Cassini mission can be found in Dougherty et al. (2009) and Seal et al.
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(2009). Maps of the Saturnian satellites are included in Roatsch et al. (2009).
141 21.2 Cassini’s Exploration of Saturn’s Icy Satellites
142 The Cassini spacecraft is equipped with instruments tailored to investigate the surfaces,
143 environments and interiors of icy satellites. The optical remote sensing (ORS) instrument suite
144 includes cameras and spectrometers designed for high spatial and spectral resolution covering
145 wavelengths between 0.06µm and 1000µm. With the imaging subsystem (ISS) (Porco et al., 2004),
146 morphologic, stratigraphic and other geological surface properties can be observed at spatial
147 resolutions down to a few meters (locally) or a few hundred meters (globally), depending on flyby
148 distances. Moreover, as many as 33 color and polarization filter combinations permit mapping
149 geologically diverse terrain spatially at wavelengths between ~0.3µm and ~1.1µm. The visible and
150 infrared mapping spectrometer (VIMS; 0.4µm to 5.1µm) provides chemical and compositional
151 spectral information with spatial resolutions from a few kilometers (globally) to better than one
152 hundred meters (locally) (Brown et al., 2004; Jaumann et al., 2006). The composite infrared
153 spectrometer (CIRS) determines global and regional surface temperatures and thermal properties on
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a kilometer scale (Flasar et al., 2004). The ultraviolet imaging spectrograph (UVIS) provides
information about thin atmospheres and volcanic plume structures as well as about water ice and
156 other minor constituents on the surface (Esposito et al., 2004), operating in the 60nm - 190nm
157 wavelength range. A suite of magnetosphere and plasma science (MAPS) instruments characterizes
158 the satellites’ environments by in-situ methods. Micron-sized dust grains and neutral molecules
159 released from the surface by active or passive processes (volcanoes/sputtering) carry surface
160 composition information to distances as far as hundreds of kilometers above the surface. Released
161 material affects the plasma surrounding the satellite and is registered by variations in the magnetic
162 field, ion density, and neutral gas and dust density as well as gas and dusty composition.
163 The primary task of the radio science subsystem (RSS) is to determine the mass of the moons (Tab.
164 1) by tracking deviations in Cassini's trajectory. RADAR data can provide unique information about
165 the upper sub-surface (Elachi et al., 2004), though few close-up RADAR SAR observations were
166 performed during satellite flybys. The lack of a scan platform for the remote sensing instruments
167 prevents the radio science, RADAR and remote sensing systems from operating simultaneously
168 during a flyby because the antenna and remote sensing instruments are oriented 90° apart on the
169 spacecraft. Even joint surface scans by the ORS instruments require significant compromises
170 between individual observations due to data rate and integration time constraints. While the ISS
171 instrument needs short 'dwell' times for mosaicking, the VIMS depend on long 'dwells' for
172 improved signal-to-noise ratios. On the other hand, the CIRS and UVIS instruments contain line-
173 scanning devices, which depend on either slow or fast slews to scan the surfaces. Thus, reaching
174 satisfactory compromises is a major challenge in the planning process.
175 Since the Cassini spacecraft orbits the planet rather than individual satellites, the moons can only be
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observed at various distances and illumination conditions, ideally during very close-targeted flybys.
Depending on the distance of a moon to Saturn, the flybys occur at very different velocities. In
178 March 2008, for instance, Enceladus was passed at ~14 km/s, while the Iapetus flyby in September
179 2007 took place at a leisurely 2.4 km/s. The flyby geometry can also vary significantly. The closest-
180 ever approach was during the Enceladus flyby on October 9, 2008, when the spacecraft skimmed as
181 low as 25 km over Enceladus’ surface. Other, more typical targeted flyby altitudes occur between
182 100 km and 2000 km. A flyby can be polar, equatorial, or in between, and at the closest approach,
183 the sub-spacecraft point over a moon can be located either over its illuminated or over its unlit side.
184 The MAPS instruments measure densities or field gradients over time. Sputtering processes lead to
185 high dust and neutral-gas densities above the surface, so that environmental in-situ instruments rely
186 on measurements as close to the surface as possible. Further plasma density variations caused by
187 satellites moving in the magnetosphere are empty flux tubes (wakes) and the drop of plasma along
188 the related L-shell. Crossing those regions is of high value for in-situ plasma investigations that
189 provide indirect information about the satellite surface and the density of its neutral and plasma
190 environment.
191 During the nominal mission, Cassini performed nine targeted as well as numerous close flybys of
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icy satellites (Tab. 2).
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193
Target Orbit Flyby date Altitude
[km]
Solar Phase Angle 2 hrs before,
at, and 2 hrs after C/A [deg]
Phoebe 0 June 11, 2004 1997 85 / 25 / 90 T
Dione B December 15, 2004 72100 108 / 85 / 62 nT
Iapetus B/C December 31, 2004 123399 87 / 94 / 100 nT
Enceladus 3 February 17, 2005 1176 25 / 114 / 158 nT
Enceladus 4 March 9, 2005 500 47 / 42 / 130 T
Hyperion 9 June 10, 2005 168000 25 / 35 / 47 nT
Enceladus 11 July 14, 2005 170 47 / 63 / 132 T
Mimas 12 August 2, 2005 62700 41 / 57 / 91 nT
Tethys 15 September 24, 2005 1503 21 /68 / 155 T
Hyperion 15 September 26, 2005 505 51 / 50 / 127 T
Dione 16 October 11, 2005 500 23 / 66 / 156 T
Iapetus 17 November 12, 2005 415400 89 / 91 / 93 nT
Rhea 18 November 26, 2005 1264 19 / 87 / 161 T
Rhea 22 March 21, 2006 82200 113 / 137 / 157 nT
Iapetus 23 April 11, 2006 602400 123 / 124 / 125 nT
Enceladus 27 September 9, 2006 39900 105 / 118 / 105 nT
Dione 33 November 21, 2006 72300 125 / 144 / 119 nT
Hyperion 39 February 16, 2007 175000 66 / 64 / 64 nT
Tethys 47 June 27, 2007 18400 148 / 116 / 49 nT
Tethys 49 August 29, 2007 55500 135 / 102 / 72 nT
Rhea 49 August 30, 2007 5800 125 / 46 / 41 nT
Iapetus 49 September 10, 2007 1620 139 / 59 / 27 T
Dione 50 September 30, 2007 43400 89 / 49 / 19 nT
Hyperion 51 October 21, 2007 122000 115 / 116 / 115 nT
Mimas 53 December 3, 2007 84200 161 / 139 / 82 nT
Enceladus 61 March 12, 2008 52 115 / 135 / 65 T
Enceladus 80 August 11, 2008 50 108 / 110 / 73 T
Enceladus 88 October 9, 2008 25 108 / 112 / 73 T
Enceladus 91 October 31, 2008 200 108 / 113 / 73 T
Tethys 93 November 16, 2008 57100 79 / 41 / 58 nT
Tethys 94 November 24, 2008 24200 123 / 159 / 84 nT
Tethys 115 July 26, 2009 68400 105 / 89 / 79 nT
Rhea 119 October 13, 2009 40400 140 / 82 / 25 nT
Mimas 119 October 14, 2009 44200 152 / 102 / 44 nT
Hyperion 119 October 16, 2009 127000 155 / 140 / 126 nT
Enceladus 120 November 2, 2009 100 177 / 90 / 5 T
Enceladus 121 November 21, 2009 1600 146 / 87 / 36 T
Rhea 121 November 21, 2009 24400 127 / 58 / 10 nT
Tethys 123 December 26, 2009 52900 136 / 77 / 17 nT
Dione 125 January 27, 2010 45100 158 / 106 / 53 nT
Mimas 126 February 13, 2010 9500 173 / 99 / 21 nT
Rhea 127 March 2, 2010 100 176 / 87 / 3 T
Dione 129 April 7, 2010 500 167 / 79 / 11 T
Enceladus 130 April 28, 2010 100 171 / 93 / 9 T
Enceladus 131 May 18, 2010 200 153 / 108 / 29 T
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Target [T]/Non
Targeted [nT]

Tethys 132 June 3, 2010 52600 154 / 99 / 45 nT
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Tab. 2 Mission characteristics of the Cassini icy satellites flybys; C/A ... closest approach
197 Three of the targeted flybys were dedicated to Enceladus, two in 2005 and one in 2008. In addition,
198 a very close but non-targeted Enceladus flyby happened in February 2005. As the data provided by
199 the MAPS instruments in March 2005 had raised suspicions about potential internal activity (Spahn
200 et al., 2006; Waite et al., 2006), the July 2005 flyby was lowered from 1000 km to 173 km to allow
201 for increased sensitivity in MAPS measurements. This flyby supplied overwhelming evidence from
202 multiple instruments of current activity on Enceladus, eventually confirmed by direct images of
203 plumes taken during a distant flyby in November 2005. The locations of the plumes were found to
204 coincide with linear tectonic features near the south pole, dubbed 'tiger stripes'. The CIRS
205 instrument also identified these 'tiger stripes' as unusually warm linear features (Spencer et al.,
206 2006); the tiger strips contain relatively large ice particles identified by VIMS (Brown et al., 2006;
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Jaumann et al., 2008).
All other moons had only one targeted flyby during the nominal mission, if any. The Tethys,
209 Hyperion, and Dione flybys in the late summer/early fall of 2005 were performed within 3 weeks of
210 each other during the same orbit of Cassini around Saturn. Due to Hyperion's chaotic rotation, it
211 was impossible to predict which part of its surface would be visible and how the surface might look.
212 By coincidence, the viewing perspective during the approach was almost identical with the best
213 Voyager 1 views. It showed a giant crater, dark material within numerous smaller craters and a very
low mean density of only 0.54g/cm 3 214
. The 'wispy streaks' of Dione were observed at very oblique
215 viewing angles during the October 2005 flyby, and were confirmed to be tectonic in origin (Wagner
216 et al., 2006). The images from this targeted Dione flyby are among the most spectacular pictures of
217 the mission. The Rhea flyby in November 2005 was dedicated to radio science, and the results from
218 that flyby are reported in Matson et al. (2009). However, while the spacecraft itself was inertially
219 pointed with its antenna towards Earth, the cameras slewed across the surface, obtaining the
220 highest-resolved images of Rhea to date.
221 In 2006, no targeted flybys took place, primarily because the spacecraft spent most of its time in
222 highly inclined orbits. The very close but non-targeted Rhea flyby in August 2007 was dedicated to
223 remote sensing and revealed small details of the prominent bright-ray crater and its ejecta
224 environment. This might be the youngest large crater in the Saturnian system. The targeted Iapetus
225 flyby in September 2007 over the anti-Saturnian and trailing hemispheres yielded very high-
226 resolution observations of the highest parts of the equatorial ridge, the transition zone between the
227 dark and the bright terrain, and the bright trailing side. Combined with data from a more distant,
228 non-targeted flyby on December 31, 2004, and other distant observations, these data permitted
229 developing a model of the formation of the global and unique brightness dichotomy of this unusual
230 moon (Denk and Spencer, 2008). For more detailed descriptions of the Cassini mission, see
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Dougherty et al. (2009), Orton et al. (2009) and Seal et al. (2009).
233 21.3. Morphology, Geology and Topography
234 Voyager’s partial global mapping at 1 km resolution suggested that the morphology of the
235 Saturnian satellites was dominated by impact craters (Smith et al., 1981). Indeed, Voyager’s best
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images, 500 m resolution views of Rhea, showed a landscape saturated with craters. Voyager also
found evidence of limited volcanic and tectonic activity in the form of relatively smooth plains and
238 linear features (Smith et al., 1981, 1982; Plescia and Boyce, 1982, 1983; Moore et al., 1985; Moore
239 and Ahren, 1983), especially on Enceladus. The Enceladus discoveries aside, Cassini’s considerably
240 improved resolution and global mapping coverage confirmed the abundance of impact craters on all
241 satellites (Dones et al., 2009) and also revealed surprisingly varied degrees of past endogenic
242 activity on Tethys, Rhea and especially Dione, indicating that these satellites were more dynamic
243 than originally thought. Observed deformations range from linear grooves (Mimas) and nearly
244 global-scale grabens (Tethys, Rhea) to repeated episodes of tectonics and resurfacing (Dione,
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Enceladus). To date, little or no direct evidence of past or present endogenic activity has been
identified on Hyperion, Phoebe or Iapetus (other than its potentially endogenic equatorial ridge).
247 This diversity is in itself a puzzle that has not yet been explained. In the following sections, we
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249
examine the features observed, their morphology, and implications for internal evolution.
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Craters
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251 Comparison with craters on larger icy satellites (e.g. Ganymede) is illuminating because surface
252 gravities differ by a factor of 5 to 10. Central-peak craters are common on Saturnian satellites.
253 Central peaks in larger craters can be as high as 10 km, many of them rising above the local mean
254 elevation. Cassini has revealed little evidence of terrace formation, and floor uplift dominates clearly
255 over rim collapse in the formation of complex craters (Schenk and Moore, 2007). Pancake (formerly
256 called pedestal) ejecta have been observed in smooth plains on Dione. This kind of ejecta deposit,
257 once only known as a feature of the Galilean satellites, now appears to be common on icy satellites
258 in general. The fact that it is not obvious on other satellites is most likely due to the difficulty of
259 identifying it on these rugged, heavily cratered surfaces.
260 A transition to larger, more complex crater landforms (e.g. central pit craters, multi-ring basins) on
261 these satellites is predicted to occur at much larger diameters (>150 km) than those seen on
262 Ganymede and Callisto (Schenk, 1993), due to their weaker gravity. True central-pit craters like
263 those on Ganymede or Callisto (Schenk, 1993) have not been observed, although the central
264 complex of Odysseus features a central depression and closely resembles its counterparts on large
265 icy moons. The dominant morphology at larger diameters (>150 km) is peak-ring or similar, with a
266 prominent central peak surrounded by a ring or an elevated circular plateau. The number of such
267 basins discovered by Cassini on Iapetus and Rhea was greater than expected. Indeed, at least 10
268 basins larger than 300 km across have been seen on Iapetus (e.g. Giese et al., 2008).
269 The most ancient large basins on Rhea and Tethys have been modified by relaxation and long-term
270 impact bombardment (e.g. Schenk and Moore, 2007). The younger basins are generally deep and
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unrelaxed. However, the floor of Evander, is the largest basin identified on Dione, has been relaxed
and uplifted to the level of the original surface, and although its rim and peak remain prominent
273 (e.g. Schenk and Moore, 2007). Relaxation of this basin is consistent with elevated heat flows on
274 Dione, as indicated by the relaxation of smaller (diameters of 25-50 km) craters within the extensive
275 smooth plains to the north of Evander. Relaxation in this basin, which has a low superposed crater
276 density, is consistent with elevated heat flows, as indicated by the smaller relaxed craters nearby and
277 the extensive smooth plains to the north. The relaxation state of impact basins can be determined by
278 comparing their depth to that of unmodified craters and examining their shape. The most ancient
279 large basins on Rhea (e.g. Tirawa) and Tethys are roughly 10 to 20% shallower than predicted and
280 feature prominent domical or conical central uplifts that approach or exceed the local elevation of
281 the satellite (Schenk and Moore, 2007), which suggests they have been modified to some degree by
282 relaxation. The youngest basins appear to be deeper and possibly unrelaxed. The most prominent
283 examples include Herschel on Mimas (diameter ~120 km, depth ~11 km) and Odysseus (diameter
284 ~450 km, depth >8 km). In neither case do the central uplifts rise above the local topography. A
285 prominent exception is the relatively young 400 km wide Evander basin on Dione, which may be
286 either ~3Gyr (Wagner et al., 2006) or less than 2Gyr old (Kirchoff and Schenk, 2009). For a more
287 detailed discussion of surface ages, see Dones et al. (2009).
288 Iapetus’ leading side shows more relief than the other medium-sized Saturnian satellites (Thomas et
289 al., 2007b; Giese et al., 2008). The lack of basin relaxation on Iapetus is consistent with the
290 presence of a thick (50–100 km) lithosphere in Iapetus’ early history. Such a lithosphere might also
291 support Iapetus’ prominent equatorial ridge. The ridge shows a complex morphology, probably the
292 result of impact erosion in some places and of post-formation tectonic modifications in others.
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The abundance of large impact basins, even in Voyager’s limited view, suggested that such basins
could have disruptive effects on these satellites (Smith et al., 1981). Several attempts to discover
295 evidence of these effects (e.g. seismic shaking, global fracturing) have met with limited success (e.g.
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297
Moore et al., 2004), but Cassini-based studies are now in progress.
298 Tectonics
299 The Voyager 1 and 2 spacecraft revealed that the medium-sized icy Saturnian satellites have
300 undergone varying degrees of tectonic deformation (Smith et al. 1981; 1982; Plescia and Boyce
301 1982, 1983). Although such deformation can provide constraints on the interior structure and
302 evolution of a satellite, the imaging resolution was often insufficient to permit sustainable
303 conclusions. Post-Voyager advances include a vastly better understanding of tectonic deformation
304 and thermal evolution in the Jovian system thanks to Galileo, and greatly improved experimental
305 measurements of relevant ice properties (Schmitt et al., 1998). In this section, we combine this
306 improved understanding with the results of the ongoing Cassini mission to discuss the current state
307 of knowledge of icy-satellite tectonics in the Saturnian system. We also examine the extent to which
308 cryovolcanic processes may be relevant. Most of the topics covered in this section are treated in
309
greater detail in Collins et al. (2009).
7

310 Tectonics on icy satellites comprises three major aspects: (1) the various mechanisms by which ice
311 deforms when subjected to tectonic stress; (2) the mechanisms by which tectonic stress may arise;
312 and (3) the observational constraints on these deformation mechanisms.
313 From a tectonics point of view, ice is rather similar to silicate materials: under surface conditions its
314 behavior is rigid or brittle, while at greater depths it is ductile. The rheology of ice is complicated
315 and cannot be discussed in depth here; useful discourses on its elastic, brittle and viscous properties
316 may be found in Petrenko and Whitworth (1999), Beeman et al. (1988) and Durham and Stern
317 (2001), respectively. Although a real ice shell will exhibit all three kinds of behavior, it can be
318 modeled as a simple elastic layer (e.g. McNutt, 1984). This effective elastic thickness depends on
319 the temperature gradient within the ice shell during deformation and can be inferred from
320 topographical measurements (e.g. Giese et al., 2007). Thus, measuring the elastic thickness places a
321 useful constraint on satellite thermal evolution (e.g. Nimmo et al., 2002). Pressures within Saturnian
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323
satellites are generally low (see Tab. 3).
R
(km)
Ms
(10 20 kg)
a
(10 3 km)
e H
(m)
3eH
(m)
8
Tidal
(10 10 W)
Rad.
(10 10 W)
Stress
(kPa)
Mimas 198 0.37 186 .0202 9180 556 78 0.013 8400 7.1 I
Enceladus 252 1.08 238 .0045 3940 53 2 0.038 630 23 III
Tethys 533 6.15 295 0. 7270 0 0 0.22 0 37 II
Dione 562 11 377 .0022 2410 16 0.86 0.39 85 97 II
Rhea 764 23 527 .001 1440 4.3 0.067 0.81 17 124 II
Titan 2575 1346 1222 .0292 254 22 45 47 26 3280 -
Iapetus 736 18 3561 .0283 5 0.4 2.6x10 -5 0.63 1.7 88 I/II
324
325
326
Tab. 3 Data from Yoder (1995) and Thomas et al. (2007). R, Ms, a and e are the radius, mass,
semi-major axis, period and eccentricity, respectively. H is the permanent tidal bulge assuming
327
328
329
330
h2=2.5. 3eH is the approximate magnitude of the diurnal tidal bulge – note that this is likely an
overestimate for small satellites because their rigidity will reduce h2. “Tidal” is the tidal heat
production assuming a homogeneous body with k2=1.5 (again, a likely overestimate for small
satellites) and Q=100. “Rad” is the radiogenic heat production assuming a chondritic rate of
3.5x10 -12 331 W/kg. “Stress” is the approximate diurnal tidal stress given by EeH/R where E is
332
333
Young’s modulus (assumed 9 GPa). Pcen is the central pressure, assuming uniform density. “MG” is
the “metamorphic grade”, modified after Johnson (1998), which is a qualitative assessment of the
334
335
degree of deformation: I=unmodified, II=intermediate, III=heavily modified
336 Thus, although ice transmutes into high-pressure forms at pressures greater than about 0.2GPa (e.g.
337 Sotin et al., 1998), this effect is relevant mainly on Titan.
338 Tectonic features may be used to infer the direction and (sometimes) the magnitude of the driving
339 stresses. Observations can thus place constraints on the mechanisms which gave rise to these
340 tectonic features. We briefly discuss several mechanisms which are potential candidates for
341 producing tectonic features.
342 Satellites experience both rotational and tidal deformation. If the rotation rate or the magnitude of
343
344
the tide changes, the shape of the satellite will change as well, generating global-scale patterns of
stress. A recent compilation of the present-day shapes of the Saturnian satellites is given by Thomas
345 et al. (2007b); a good general description of tidal and rotational satellite deformation is given in
346 Murray and Dermott (1999).
347 If a synchronously rotating satellite’s orbital eccentricity is zero, the tidal bulge will be at a fixed
348 geographic point, and of constant amplitude. The maximum amplitude H of this static (or
349 permanent) tidal bulge depends on the mass and radius of the satellite, the mass of the primary, the
350
351
352
distance between them, and the rigidity of the satellite. The latter is described by the h2 Love
number, where h2 has a value of 5/2 for a fluid body of uniform density which, however, declines as
the rigidity (or shear modulus) µ of the satellite increases (e.g. Murray and Dermott, 1999).
353 In a synchronous satellite with non-zero orbital eccentricity, the amplitude and direction of the
354 static tidal bulge will change slightly, resulting in a time-varying diurnal tide. Causing diurnal
355 stresses, this tide has an amplitude expressed by 3eH, where e is the eccentricity and H is the
356 equilibrium tidal bulge (e.g. Greenberg and Geissler, 2002). Obliquity librations are potentially an
357 additional source of tidal stress (Bills and Ray, 2000); they can be large if their period is a small
358
multiple of the orbital period (Wisdom, 2004).
Pcen
(MPa)
MG

359 Table 3 shows the diurnal tidal stresses of the icy Saturnian satellites at zero rigidity, demonstrating
360 the general decrease in stress as the semi-major axis lengthens. For comparison, the diurnal stresses
361 on Europa thought to be responsible for cycloidal features are a few tens of kPa (Hoppa et al.,
362 1999a). Because the principal stresses rotate during each orbit, contraction, extension and mixed-
363 mode horizontal principal stresses which promote strike-slip motions are all possible as a satellite
364 orbits (Hoppa et al., 1999b). This motion can lead to both shear heating (Nimmo and Gaidos, 2002)
365 and the monotonic accumulation of strike-slip offsets (Hoppa et al., 1999b).
366 A satellite surface may move longitudinally with respect to the static tidal bulge due to tidal
367 (Greenberg and Weidenschilling, 1984) or atmospheric torques (Lorenz et al., 2008) or, in the case
368 of floating ice shells, if the thermally equilibrated ice shell thickness variations are not compatible
369 with rotational equilibrium (Ojakangas and Stevenson, 1989a). This non-synchronous rotation leads
370 to surface stresses (Helfenstein and Parmentier, 1985). Stresses from non-synchronous rotation
371 increase with the angular speed of rotation but will dissipate with time owing to viscous relaxation
372 (Greenberg and Weidenschilling, 1984; Wahr et al., 2008). Maximum tensile stress depends on the
373 amount of reorientation and the degree of flattening, f, which depends in turn on the satellite
374 rotation rate (Leith and McKinnon, 1996). In non-rigid satellites, the stresses due to non-
375 synchronous rotation will exceed the diurnal tidal stresses for a reorientation typically in excess of
376 roughly one degree. In certain circumstances, the surface of a satellite may reorient relative to its
377 axis of rotation (Willemann, 1984; Matsuyama and Nimmo, 2008) and give rise to global fracture
378 patterns (e.g. Schenk et al., 2008). This process of true polar wander is conceptually very similar to
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380
non-synchronous rotation (Melosh, 1980; Leith and McKinnon, 1996), and 'reorientation' will in the
following be understood as implying either non-synchronous rotation or true polar wander. In
381 general, rotation axis motion occurs roughly perpendicular to the tidal axis because this path is
382 energetically favored (Matsuyama and Nimmo, 2007). True polar wander can occur because of ice
383 shell thickness variations (Ojakangas and Stevenson, 1989b), long-wavelength density anomalies
384 (Janes and Melosh, 1988; Nimmo and Pappalardo, 2006), volatile redistribution (Rubincam, 2003)
385 or impacts, either directly (Chapman and McKinnon, 1986) or due to the creation of a new impact
386 basin (Melosh, 1975; Murchie and Head, 1986; Nimmo and Matsuyama, 2007). Like non-
387 synchronous rotation, the maximum stress developed by true polar wander depends on the amount
388 of reorientation and the effective flattening f, with the value of f depending on whether any
389 reorientation of the tidal axis has occurred (Matsuyama and Nimmo, 2008).
390 Satellites that are initially rotating at a rate faster than synchronous will despin (i.e. reduce their
391 rotation rate) to synchronous rotation in timescales that are generally very short compared to the age
392 of the solar system (e.g. Murray and Dermott, 1999). Spin rate reductions result in shape changes
393 and stresses – compressive at the equator as the equatorial bulge collapses, and tensile at the poles
394 as the poles elongate. The maximum differential stress caused by despinning in a hydrostatic body
395 depends mainly on the satellite radius and the difference between the initial and the final angular
396 rotation velocities (Melosh, 1977). Spin-up may occur in some circumstances (e.g. if the satellite
397 undergoes differentiation), and in this case the signs of all the stresses are reversed. Stresses along
398 lines of longitude are always larger than the stresses along lines of latitude. Thus, irrespective of
399 whether a satellite is spinning down or spinning up, equatorial tectonic features are expected to be
400 oriented in a north-south direction. This result is important when considering the origin of the
401
402
global equatorial ridge on Iapetus.
Synchronous satellites, which are evolving outwards due to tidal torques from the primary, will
403 undergo both a reduction in their spin rate and a parallel reduction in the tidal bulge amplitude. This
404 combination of despinning and tidal bulge reduction causes a stress pattern in which a region
405 around the sub-planet point experiences compressive stress, middle latitudes experience horizontal
406 shear stress, and the poles undergo tension (Melosh, 1980; Helfenstein and Parmentier, 1983). The
407 maximum principal stress difference is similar to that produced by despinning alone (Melosh,
408 1980).
409 A large variety of mechanisms can lead to volume changes within a satellite, and thus to extensional
410 or contractional features on the surface (Squyres and Croft, 1986; Kirk and Stevenson, 1987;
411 Mueller and McKinnon, 1988). Volume changes generate isotropic stress fields on the surface. A
412 potentially important source of expansion or contraction is the large density contrast between ice I
413 and water. As water freezes to ice I, large surface extensional stresses may result (Cassen et al.,
414 1979; Smith et al., 1981; Nimmo, 2004a). The reason for this effect is that, on a spherical body, the
415 increased volume of ice relative to water drives the ice shell outwards; and this increase in radius
416
results in extension. A fractional change in radius of 0.1% gives rise to stresses of about 10MPa.
9

417 This effect is much reduced if there are high-pressure ice phases freezing simultaneously (Squyres,
418 1980; Showman et al., 1997).
419 Similar but smaller effects are caused by thermal expansion or contraction (Ellsworth and Schubert,
420 1983; Hillier and Squyres, 1991; Showman et al., 1997; Nimmo, 2004a). A global temperature
change of 100K will cause stresses on the order of 100 MPa for a thermal expansivity of 10 -4 K -1 421 .
422 Warming may also lead to silicate dehydration (e.g. Squyres and Croft, 1986), which in turn causes
423 expansion.
424 Satellites are quite likely to have non-axisymmetric structures, in which case some of the above-
425 mentioned global mechanisms may lead to local deformation. For instance, local zones of weakness
426 or pre-existing structures can result in enhanced tidal dissipation and deformation (e.g. Sotin et al.,
427 2002; Nimmo, 2004b) and/or the alteration of local stress trajectories. Convection (thermal or
428 compositional) and buoyancy forces due to lateral shell thickness variations can generate local
429 deformation. They are discussed at greater length in Collins et al. (2009). We will not discuss these
430 processes further here, as they do not tend to produce globally organized tectonic structures.
431 Large impact events can also potentially create global fracture networks. This mechanism is implicit
432 in the model presented by E. Shoemaker in Smith et al. (1981) for the catastrophic breakup and re-
433 accretion of the medium-sized icy satellites of Saturn. Presumably, impacts not quite large enough
434 to completely disrupt a satellite could do considerable damage in the form of global fracturing, and
435 post-Voyager studies of these satellites focused in part on attempts to find evidence of an incipient
436 break-up of this kind. For instance, the surface of Mimas is crossed by a number of linear-to-
437
438
439
arcuate, sub-parallel troughs (Fig. 1) that are thought to be the consequence of global-scale fractures
during the Herschel impact event (McKinnon, 1985; Schenk, 1989a).
440 Fig. 1 Linear troughs on Mimas, extending east to west. These troughs are interpreted as fractures,
441 possibly related to the Herschel crater. Orthographic map projection at 400m/pixel is centered on
442 the antipode to Herschel, the largest impact basin on Mimas.
443
444 Voyager and Cassini mapping indicates that these features, if related to Herschel, do not form a
445 simple radial or concentric pattern. Nor has it been demonstrated that any of these features form on
446 a plane intersecting Herschel, or whether they should do so. It thus remains possible that these are
447 random fractures formed during freeze expansion of the Mimas interior. Detailed mapping remains
448 to be done to test these hypotheses. Mimas does not otherwise exhibit endogenic landforms and
449 will, therefore, not be discussed further.
450 Enceladus shows the greatest tectonic deformation of any of the Saturnian satellites, and very
451 significant spatial variations in surface age (Smith et al., 1982; Kargel and Pozio, 1996; Porco et al.,
452 2006; Jaumann et al., 2008). Spencer et al. (2009) discusses the tectonic behavior of Enceladus in
453 much greater detail; only a very brief summary is given here.
454 Three different terrain types exist on Enceladus. Ancient cratered terrains cover a broad band from
0 to 180 o longitude. Centred on 90 o and 270 o 455
longitude there are younger, deformed terrains roughly
90 o wide at the equator. The southern polar region south of 55 o 456
S is heavily deformed and almost
457 uncratered. The most prominent tectonic features of this region are the linear depressions called
458 ‘tiger stripes’ (Porco et al., 2006). These tiger stripes are typically ~500m deep, ~2 km wide, up to
459
460
~130 km long and spaced ~35 km apart. Crosscutting fractures and ridges with almost no
superimposed impact craters characterize the area between the stripes. The tiger stripes are
461 apparently the source of the geysers observed to emanate from the south polar region (Spitale and
462 Porco, 2007). The energy source of these geysers is unknown, but it may be shear heating generated
463 by strike-slip motion at the tiger stripes (Nimmo et al., 2007). The coherent (but latitudinally
464 asymmetrical) global pattern of deformation on Enceladus is currently unexplained, but it is most
465 probably due to shape changes, perhaps related to a hypothetical episode of true polar wander
466 (Nimmo and Pappalardo, 2006).
467
468
Nearly encircling the globe, Ithaca Chasma is the most prominent feature on Tethys (Fig. 2).
469 Fig. 2 Tethys and Ithaca Chasma. Topography was derived from Cassini stereo images and has a
470 horizontal resolution of 2-5 km and a vertical accuracy of 150-800m. Profiles have superimposed
471
472
model flexural profiles (in red). Te is elastic thickness. (Reproduced from Giese et al. (2007))
473 It extends approximately 270° around Tethys and is not a complete circle. It is narrowly confined to
474 a zone which lies along a large circle whose pole is only ~20° from the center of Odysseus, the
475
relatively fresh largest basin on the satellite which is 450 km wide (Smith et al., 1982; Moore and
10

476 Ahern, 1983). Smith et al. (1981) suggested that Ithaca Chasma was formed by freeze-expansion of
477 Tethys's interior. They noted that if Tethys was once a sphere of liquid water covered by a thin solid
478 ice crust, freezing in the interior would have produced an expansion of the surface comparable to
479 the area of the chasm (~10% of the total satellite surface area). This hypothesis fails to explain why
480 the chasm occurs only within a narrow zone; besides, it is difficult to account for the geophysical
481 energy needed for the presence of a molten interior. Expansion of the satellite's interior should have
482 caused fracturing over the entire surface in order to effectively relieve stresses in a rigid crust
483 (Moore and Ahern, 1983).
484 The nearly concentric geometrical relationship between Odysseus and Ithaca Chasma prompted
485 Moore and Ahern (1983) to suggest that the trough system was an immediate manifestation of the
486 impact event, perhaps caused by a damped, whole-body oscillation of the satellite. Based on the
487 deep topography of Odysseus, Schenk (1989b) also suggested that it had not undergone substantial
488 relaxation, and that Ithaca Chasma was the equivalent of a ring graben formed during the impact
489 event by a prompt collapse of the floor involving a large portion of the interior. In a Cassini-era
490 study, Giese et al. (2007) doubted the Odysseus-related hypothesis (Moore and Ahern, 1983;
491 Schenk, 1989b) on the basis of crater densities within Ithaca Chasma relative to Odysseus, from
492 which they concluded that Ithaca Chasma is older than Odysseus and therefore could not have
493 influenced its formation. They went on to measure photogrammetrically the raised flanks of Ithaca
494 Chasma, which they attributed to flexural uplift. Assuming that these raised flanks were indeed due
495 to flexural uplift, they derived a mechanical lithospheric thickness of ~20 km and surface heat
fluxes of 18-30 mW/m 2 at assumed strain rates of 10 -17 to 10s -1 496
497
. They concluded that Ithaca Chasma
is an endogenic tectonic feature but could not explain why it is confined to a narrow zone
498 approximating a large circle.
499 Several Voyager-era studies of Dione (Plescia, 1983; Moore, 1984) noted a near-global network of
500 tectonic troughs and the presence of a smooth plain crossed by both ridges and troughs located near
501 the 90°W longitude and extending eastward. Cassini coverage revealed the westward extension of
502 this plain (Wagner et al., 2006) and demonstrated that the fractured and the cratered dark plains are
503 spectrally distinct (Stephan et al., 2008). Like the eastern portion observed by Voyager, the western
504 plain also exhibits both ridges and troughs (Fig. 3).
505
Fig. 3 (a) Image mosaic of Dione, centered on 180 o 506
longitude. Red and white arrows point to Amata
507 and Turnus impact structures, respectively. (b) Topography of Dione, from Schenk and Moore
508 (2007). An ancient degraded impact basin centered just west of Amata can be identified by the two
509 concentric rings in the DEM (white arrows). DEM has a dynamic range of 8 km (red is highest,
510
511
magenta is lowest).
512 However, the ridge system is much more pronounced. The western ridges are planimetrically
513 narrower and morphologically better expressed (e.g. their relative lack of degradation). These
514 western ridges also conform much more closely to the boundary between the plain and the older
515 cratered rolling terrain to the west. The linear to arcuate troughs interpreted to have been created by
516 tectonic extension form a complex network that divides Dione’s surface into large polygons (Smith
517 et al., 1981, 1982; Plescia, 1983; Moore, 1984). The troughs often form parallel to sub-parallel sets,
518
519
somewhat reminiscent of the grooved terrain on Ganymede. The global orientation of Dione’s
tectonic fabric is non-random (Moore, 1984) and exhibits patterns consistent with a decline in
520 oblateness due to either despinning or orbital recession (Melosh 1977, 1980).
521 The topography derived from Cassini images of Dione shows one clear case of tectonic orientation
522 regionally influenced by a preexisting but very degraded large impact basin centered on the smaller
523 crater Amata (Schenk and Moore, 2007 (Figs. 3 and 4)).
524 Fig. 4 Fracture zones on Dione. Global base map of Cassini imaging has a resolution of 400m. The
525 circle represents the outer ridge corresponding to the rim of the degraded basin near Amata shown
526
527
in Fig. 3.
528 Moore et al (2004) identified another example of tectonic lineament orientation associated with an
529 unambiguously identified impact feature, 90 km Turnus. Such crater-focused tectonics has been
530 identified on Ganymede from Galileo data (e.g. Pappalardo and Collins, 2005). This crater-
531 lineament relationship on Dione is the most pronounced among any of the middle-sized icy
532 satellites except Enceladus (Spencer et al., 2009). Moore (1984) proposed a tectonic history of
533
extension followed by (at least regional) compression, the latter being responsible for the ridges.
11

534 Cassini-era images of the tectonic troughs (mostly grabens) show them to be morphologically fresh,
535 implying that extension continued into geologically recent times (Wagner et al., 2006).
536 After Voyager, Rhea was held to be the least endogenically evolved satellite larger than 1000 km in
537 diameter. Moore et al. (1985) pointed out a few parallel lineaments and noted several large ridges or
538 scarps (which they termed ‘megascarps') and speculated that these features were formed by a period
539 of extension followed by compression. Cassini images, and digital terrain models derived from
540 them, of the trailing hemisphere of Rhea seen from both middle northern and far southern latitudes
541 clearly show that the wispy arcuate albedo markings seen in Voyager images are associated with a
542 graben/extensional fault system dominantly trending north-south (Moore and Schenk, 2007;
543
544
Wagner et al. 2007) (Fig. 5a).
545 Fig. 5 (a) Cassini image of the northern trailing hemisphere of Rhea showing the bright scarps of
546 the northern extension of the graben system and their relationship to 'wispy' terrain. (b) Mosaic of
547 Cassini images showing curvilinear grabens (black arrows) and an orthogonal set of ridges and
548
549
troughs (white arrow). Base mosaic resolution is 1 km/pixel.
550 This is the first identification of unambiguous regional endogenic activity on Rhea. The digital
551 terrain models of Rhea derived from Cassini images also show a set of roughly north-south trending
552
553
ridges (Figs. 5b, and 6) that correspond to the 'megascarps' of Moore et al. (1985).
554
555
556
Fig. 6 Preliminary stereo topography map, from Schenk and Moore (2007). White arrows indicate
graben systems. The DEM has a dynamic range of 10 km (red is highest, magenta is lowest).
557 If these ridges are in fact due to compressional tectonics, their geometrical relation to the
558 graben/scarp system may also indicate a genetic relationship. However, the poor morphological
559 presentation of these ridges implies that they are old, while the 'wispy terrain' grabens are fresh and
560 presumably much more recently formed.
561 While the troughs and ridges suggest mainly extensional and (minor) compressional tectonics
562 (Thomas, 1988), the en-echelon pattern of the scarps and troughs on the trailing hemisphere
563 suggests shear stress (Wagner et al., 2007). This pattern may have been influenced by the possible
564 presence of a large, degraded impact basin (basin C of Moore et al., 1985).
565
566
The equatorial ridge of Iapetus (Fig. 7) is the most enigmatic feature on any of the medium-sized
567 Fig. 7 Topography of Iapetus as derived from stereo Cassini imaging. Note the sharply defined
568 equatorial ridge. The DEM at left is referenced to a sphere, the DEM at right to a biaxial ellipsoid.
569
570
(Adapted from Giese et al., 2008)
571 icy satellites (Porco et al., 2005a; Giese et al., 2008); it extends across more than half of Iapetus'
572 circumference (Giese et al., 2008). Different parts show very different morphologies, from
573 trapezoidal to triangular cross-sections with steep walls and a sharp rim to isolated mountains with
574 moderate slopes and rounded tops (Giese et al., 2008). Although the ridge is variable in height, it
575 rises up to 10 km above the local mean but also has a gently sloping flank. The ridge is abundantly
576
577
cratered, indicating that it is comparable in age to other terrains on Iapetus (Schmedemann et al.,
2008). Although the equatorial location of the ridge might be due to despinning (Castillo-Rogez et
578 al., 2007), it represents such a large mass that, if it had formed in another location, it would almost
579 certainly have reoriented to the equator anyway. However, as already mentioned above, irrespective
580 of whether a satellite is spinning down or spinning up, equatorial tectonic features can be expected
581 to be oriented north-south, making it difficult to explain the ridge with despinning. One possibility
582 is that despinning stresses combined with lithospheric thickness variations due to convection may
583 have resulted in equatorial extension and diking (Roberts and Nimmo, 2009; Melosh and Nimmo,
584 2009). Other proposals suggest that the morphology of the ridge indicates that the surface was
585 warped up by tectonic faulting (Giese et al., 2008), that it is a result of extensional forces acting
586 above an ascending current of solid-state convection from a two-cell convection pattern
587 (Czechowski and Leliwa-Kopystyński, 2008), or that it was formed by the impact of an ancient
588 debris disk on Iapetus (Ip, 2006). The great apparent elastic thickness suggested by the survival of
589 the ridge may be used to place constraints on the thermal evolution of Iapetus (Castillo-Rogez et al.,
590 2007).
591 Although the global shape of Iapetus does suggest an earlier, more rapid spin rate (Castillo-Rogez et
592
al., 2007), despinning cannot explain the equatorial ridge (see above): stresses in the north-south
12

593 direction are always smaller than those in the east-west direction, and will result in N-S oriented
594 features (Melosh, 1977). Thus, the origin of the equatorial ridge is currently a mystery.
595 The deformation of Saturnian satellites varies wildly, from currently active and heavily deformed
596 Enceladus to nearby Mimas, which is tectonically uninteresting. Nor does the degree of deformation
597 correlate in any straightforward way with predicted tidal stresses. However, there are at least two
598 global yet somewhat contradictory aspects, which may be identified. On the one hand, coherent
599 global tectonic patterns are evident on Dione, Tethys and Enceladus and to a lesser extent on Rhea
600 and Iapetus. On the other, several satellites, especially Tethys and Enceladus, show very
601 pronounced spatial variations in the extent of deformation. These two aspects suggest that global
602 sources of stress are common, but also that lateral variations in the mechanical properties of the
603 satellites’ ice shells play an important role.
604 Regarding this second point, localization of deformation (e.g. at the south pole of Enceladus or at
605 Ithaca Chasma) may arise naturally in ice shells (e.g. Sotin et al., 2002; Nimmo, 2004b) but is
606 currently poorly understood and makes modeling tectonic deformation much harder. As far as
607 global deformation patterns are concerned, several satellites (Enceladus, Iapetus, perhaps Dione)
608 show patterns with simple symmetries about the rotational and tidal axes. This symmetry may be
609 due either to feature reorientation or the operation of mechanisms (tides, despinning etc.), which
610 have a natural symmetry. On Dione, Rhea and Iapetus at least, diurnal tidal stresses are rather small
611 (Tab. 3), and most of the tectonic features are apparently ancient, which suggests that despinning,
612 volume changes or reorientation are more likely to be responsible for the observed deformation than
613
614
present-day tidal stresses.
As in the Galilean satellites, extensional features are generally more common than compressional
615 features on the other Saturnian satellites. This observation is readily explained if we assume that the
616 satellites once contained subsurface oceans which progressively froze (e.g. Smith et al., 1981;
617 Nimmo, 2004a); alternative explanations are that compressional features such as folds (Prockter and
618 Pappalardo, 2000) are less easy to recognize than extensional features, or that larger stresses are
619
620
required to form compressional features (see Pappalardo and Davis, 2007).
621 Cryovolcanism
622 Cryovolcanism can be defined as 'the eruption of liquid or vapor phases (with or without entrained
623 solids) of water or other volatiles that would be frozen solid at the normal temperature of the icy
624 satellite’s surface' (Geissler, 2000). Cryovolcanism therefore includes effusive eruptions of icy
625 slurries and explosive eruptions consisting primarily of vapor (Fagents, 2003). Voyager-era data
626 suggested that such activity might be common in the outer solar system, resulting in a flurry of
627 theoretical papers (e.g. Stevenson, 1982; Allison and Clifford, 1987; Crawford and Stevenson,
628 1988; Kargel. 1991; Kargel et al., 1991). Higher-resolution images from Galileo generally revealed
629 tectonic rather than cryovolcanic resurfacing processes, although some features may be due to
630 cryovolcanic activity (e.g. Schenk et al., 2001; Fagents et al., 2000; Miyamoto et al., 2005). In the
631 following, we discuss the current state of knowledge about cryovolcanism in the Saturnian
632 satellites. Enceladus, which is definitely cryovolcanically active (Porco et al., 2006), and Titan,
633 which may be (Lopes et al., 2007; Nelson et al., 2009a,b), are discussed by Spencer et al. (2009),
634 Jaumann et al. (2009) and Sotin et al. (2009).
635
636
The principal difference between cryovolcanism and silicate volcanism is that in the former, the
melt is denser than the solid (by ~10%). This means that special circumstances must exist for an
637 eruption of non-vapor cryovolcanic materials to occur (though of course on Earth it is not
638 uncommon for denser basaltic lavas to erupt through lighter granitic material). Such circumstances
639 include one or more of the following: (1) the melt includes dissolved volatiles, which exsolve and
640 reduce the bulk melt density (Crawford and Stevenson, 1988); (2) the melt is driven upwards by
641 pressure differences caused by tidal effects (Greenberg et al., 1998), surface topography (Showman
642 et al., 2004) or freezing (Fagents, 2003; Manga and Wang, 2007); (3) compositional differences
643 render the ice shell locally dense or the melt less dense (e.g. by the addition of ammonia) (Kargel
644 1992, 1995). The latter mechanism can be nothing more than a local effect, since if the bulk ice
645 shell is denser than the underlying material; it will sink if the ice shell is disrupted.
646 Erupted material consisting primarily of water plus solid ice will simultaneously freeze and boil, the
647 latter process ceasing once a sufficiently thick ice cover is formed (Allison and Clifford, 1987). The
648 resulting flow velocity will depend on the viscosity of the ice-water slurry (Kargel et al., 1991) as
649 well as on flow thickness and local slope (Wilson et al., 1997; Miyamoto et al., 2005). If the erupted
650 material is dominated by the vapor phase, the vapor plume will expand outwards (Tian et al., 2007)
651
and escape entirely if the molecular thermal velocities exceed the escape velocity. Any associated
13

652 solid material will follow ballistic trajectories, ultimately forming deposits centered on the eruption
653 vents (Fagents et al., 2000).
654 The plume properties measured on Enceladus by Cassini imply considerably smaller velocities for
655 icy dust grains than for vapor. On hypothesis is that the dust grains condense and grow in vertical
656 vents of variable width. Repeated wall collisions and re-acceleration by the gas in the vents lead to
657 effective friction, which explains the observed plume properties (Schmidt et al., 2008). ). Most dust
658 grains are growing by direct condensation of the ascending water vapor within the ice channels and
659 contain almost pure water ice; they are free of atomic sodium. In 2008 the CDA instrument while
660 analyzing the composition of E-ring particles discovered a second type of icy dust particles. These
661 grains are rich in sodium salts (0.5 - 2% in mass) and there origin is explained by the existence of
662 an ocean of liquid water below the surface of Enceladus and which is in direct contact with its rock
663 core (Postberg et al., 2009). This shows that cryovolcanism can provide valuable information about
664 geological conditions of the interior.
665 The plains of Dione are another candidate for cryovolcanic resurfacing (Fig. 8). Several Voyager-
666
667
era studies concluded that these plains were
668 Fig. 8 Cassini view of smooth plains on Dione. Older heavily cratered terrains lie to the west. Note
669 the prominent north-south ridge and the subtle linear features to the east. Mosaic resolution is
670
671
400m/pixel.
672
673
formed in this way (e.g. Plescia and Boyce, 1982; Plescia, 1983; Moore, 1984). Moore (1984)
suggested that several branching broad-rimmed troughs north of the equator at 60°W may have
674 been the source vents. Surprisingly, the plains unit may be topographically higher than the cratered
675 terrain (Moore and Schenk, 2007). This observation seems a little difficult to reconcile with the
676 plains being emplaced as a low-viscosity cryovolcanic flow. One possibility is that the plains
677 material was emplaced as a deformable plastic unit (like terrestrial glaciers), but this, or any other,
678 explanation needs both further evidence and thought. The ridges discussed above bound (or at least
679 occur near the edges) the plains unit in the west and southeast and may be related to their formation.
680 A plains unit on Tethys is located antipodal to the very large Odysseus impact feature. Moore and
681 Ahern (1983) proposed that the plains were formed by cryovolcanism. Alternatively, Moore et al.
682 (2004) considered the possibility that it was extreme seismic shaking caused by energy focused
683 there by the Odysseus impact that formed the antipodal plains. If seismic shaking was responsible,
684 there would probably be a significant transition zone between that unit and the adjacent hilly and
685 cratered terrain. Digital terrain models derived from Cassini images indicate that the transition
686 between the cratered terrain and the plains is abrupt and the plains themselves are level, supporting
687 the hypothesis that they are endogenic (Moore and Schenk, 2007).
688 Both Tethys and Dione show evidence for weak plasma tori (Burch et al., 2007) somewhat
689 analogous to the much more marked feature at Enceladus, which is caused by geysers (Porco et al.,
690 2005). It is not yet clear whether surface activity or some other process is generating these tori.
691 Similarly, Clark et al. (2008a) report a possible, but currently unconfirmed, detection of material
692 around Dione, which may also indicate some surface activity. The Cassini magnetometer has also
693 measured small amounts of mass loading near Dione (Burch et al., 2007).
694
695
So far, none of the digital terrain models of Rhea developed in the Cassini era shows any regions
that would qualify as plains, supporting the perception that Rhea shows no record of cryovolcanic
696 resurfacing. However, unlike Dione, the diffuse nature of the wisps (the 'wispy terrain') on Rhea has
697 yet to be explained as numerous smaller bright scarps associated with fracture and faulting, and may
698 yet be shown to be an expression of cryo-pyroclastic mantling (Stevenson, 1982). Furthermore,
699 Rhea, surprisingly, exhibits evidence for faint particle ring (Jones et al., 2008), the origin of which
700 is unclear but may be related to endogenic or impact processes. On the other hand, Cassini VIMS
701 observations found no evidence of a water vapor column around Rhea (Pitman et al., 2008). On
702 Iapetus, no evidence of cryovolcanism has been detected thus far. If volcanism ever operated on
703
704
these satellites, it occurred in ancient times and its evidence is now unseen.
705 21.4 Composition and Alteration of Surface Materials
706 The composition of the satellites of Saturn is determined by remote spectroscopy, by sampling
707 materials that were sputtered from the satellite surfaces or, in the case of Enceladus, injected into space
708 where they can be directly sampled by spacecraft instruments.
709 It has long been known that the surfaces of Saturn’s major satellites, Mimas, Enceladus, Tethys,
710
Dione, Rhea, Hyperion, Iapetus and Phoebe, are predominantly icy objects (e.g. Fink and Larsen,
14

711 1975; Clark et al., 1984, 1986; Roush et al., 1995; Cruikshank et al., 1998a; Owen et al., 2001;
712 Cruikshank et al., 2005; Clark et al., 2005, 2008a; Filacchione et al., 2007, 2008). While the
713 reflectance spectra of these objects in the visible range indicate that a coloring agent is present on
714 all surfaces except Mimas, Enceladus and Tethys (Buratti 1984). Only Phoebe and the dark
715 hemisphere of Iapetus display near-IR spectra markedly different from very pure water ice (e.g.
716 Cruikshank et al., 2005; Clark et al., 2005).
717 The major absorptions in icy satellite spectra shown in Figs 9 a-e are due to water ice and can be
718
719
observed at 1.5µm., 2.0µm., 3.0µm., and 4.5µm.
720 Fig. 9 VIMS spectra of the Saturnian icy satellites: (a) The major absorptions in icy satellite spectra
721 are due to water ice and appear at 1.5, 2.0, 3.0, and 4.5µm; (b) Representative (A) bright and (B)
722 dark region Dione spectra (the gaps in the bright region spectra (A) near 1µm are due to sensor
723
724
725
saturation); the inset shows an area of high concentration of CO2 (after Clark et al., 2008a); (c)
Spectra of Hyperion showing an increase in blue reflectivity and absorptions of water ice and CO2
(after Clark et al., 2008a); (d) Spectra of Iapetus showing a range of blue peak intensities,
726
727
absorptions of different intensities by water ice and CO2 absorption bands (after Clark et al.,
2008a); (e) Spectra of Iapetus showing a range of blue peak intensities and absorptions of different
728
729
intensities by water ice and CO2 absorption bands (after Clark et al., 2008a). However, no CO2 was
found on the small satellites Atlas, Pandora, Janus, Epimetheus, Telesto, and Calypso (Buratti et
730
731
al., 2009a).
732
733 Weaker ice absorptions appear at 1.04µm and 1.25µm in the spectra of bright regions where the ice
734
735
is purer. Trapped CO2 was first detected in Galileo near-infrared mapping spectrometer (NIMS)
spectra of the Galilean satellites (Carlson et al., 1996) and in the Saturnian system on Iapetus
736 (Buratti et al, 2005), Phoebe (Clark et al., 2005), Hyperion (Cruikshank et al., 2007), and Enceladus
737
738
(Brown et al., 2006). Weak CO2 absorptions were also detected in VIMS data on Dione, Tethys,
Mimas and Rhea (Clark et al., 2008a).
739 Organic compounds were detected directly in the Enceladus plume by the INMS (Waite et al.,
740 2006) and in trace amounts by the VIMS on Phoebe (Clark et al., 2005), Hyperion (Cruikshank et
741 al., 2007) and Iapetus (Cruikshank et al., 2008), and possibly Enceladus (Brown et al., 2006). While
742 the spectral signatures indicate aromatic hydrocarbon molecules, exact identifications have remained
743 elusive, partly because the detections have low signal-to-noise ratios. Another spectral feature
744 observed in VIMS data of dark material on Phoebe, Iapetus and Dione appears to be an absorption
745 centered near 2.97µm (Fig. 9 d). The detection of this feature is complicated by the fact that VIMS
746 has an order-sorting filter change at this wavelength. Clark et al. (2005) presented evidence that this
747 feature is real and maps to geological patterns on Dione and Iapetus (Clark et al., 2008a, 2009).
748 However, a better understanding of the instrument response is needed to increase confidence in the
749 identification, including confirmation by a different instrument. If correct, the feature strength
750 indicates approximately 1% ammonia in the dark material.
751 The position shift of amorphous ice absorptions from their crystalline counterparts (Mastrapa et
752 al., 2008 and references therein) can be used to determine whether amorphous or crystalline ice is
753
754
present on a surface. Traditionally, amorphous ice displays almost no 1.65µm feature and the
3.1µm Fresnel peak shifts, broadens and weakens. Newman et al. (2008) used variations in the
755 3.1µm peak and the 1.65µm absorption to argue for amorphous ice around the 'tiger stripes' on
756 Enceladus. However, Clark et al. (2009) showed that scattering and sub-micron grain sizes also
757 influence the intensities of the 3.1-µm peak and the 1.65-µm absorption, and can be confused with
758 temperature and crystallinity effects. Further examination of the ice feature positions in spectra of
759 Enceladus and other satellites in the Saturnian system showed that positions and strengths are
760 consistent with those of fine-grained crystalline ice (Clark et al., 2009). Newman et al. (2009)
761 showed that water ice on Dione exists exclusively in crystalline form.
762 The spectra of the satellites in the visible wavelength range are smooth with no sharp spectral features,
763 but all show an ultraviolet absorber affecting the blue end of the visible-wavelength spectral slope.
764 Visible spectra can be classified by their spectral slope: from bluish Enceladus and Phoebe to redder
765 Iapetus, Hyperion and Epimetheus. In the 1µm to1.3µm range, the spectra of Enceladus, Tethys,
766 Mimas and Rhea are characterized by negative slope, consistent with a surface mainly dominated by
767 water ice, while the spectra of Iapetus, Hyperion and Phoebe show considerable reddening, pointing
768 out the relevant role played by the darkening materials present on the surface (Filacchione et al., 2007,
769
2008). In between these two classes are Dione and Epimetheus, which have a relatively flat global
15

770 spectrum in this range. At wavelengths less than about 0.5µm, all satellite spectra display a UV
771 absorber, which might be due to charge transfer in oxides such as nano-phase hematite (Clark et al.,
772 2008a), with the apparent strength of the absorber generally indicating the amount of contaminant.
773 Trends in the UV absorber are complicated by ice dust contributions to satellite surfaces transmitted
774 from Enceladus via the E-ring. Satellites close to the rings, such as Janus, Atlas and Prometheus, show
775 an intermediate red tint between that of the rings and Enceladus. Mimas and Tethys show a red tint
776 between that of Janus, Atlas and Prometheus and that of Enceladus on the other (e.g. Cuzzi et al.,
777 2009; Buratti et al., 2009a).
778 Disk-integrated ultraviolet observations by the Cassini UVIS (Esposito et al., 2004; Hendrix and
779
780
Burrati, 2009) (Fig. 10) show Tethys and Dione exhibit marked UV
781 Fig. 10 Cassini UVIS disk-integrated spectra of the icy moons, all at phase angles between 10° and
782
783
13.6°.
784 leading-trailing differences whereas Mimas, Enceladus and Rhea do not, suggesting compositional
785 variations among the satellites as well as spectrally active components in the ultraviolet. UVIS
786 spectra of all the icy moons of Saturn are dominated by a pronounced water-ice absorption edge at
787 ~0.165µm. Its strength varies with the abundance of water ice, and its wavelength depends on grain
788 size. Nearly all UVIS spectra display a small absorption feature at ~0.185µm (Hendrix and Hansen,
789 2008b) whose origin is not clear at present. It is possible that this feature is due to water ice as it is
790
791
seen in published reflectance spectra of water ice (Pipes et al., 1974; Hapke et al., 1981). Published
optical constants (Warren, 1982) do not exhibit the feature, but the constants have not been well
792 measured in the ~0.17µm to 0.4µm range.
793 Cassini spectral coverage of the icy moons is complete in the 0.4µm to 5.2µm and in the 0.1µm to
794 0.19µm region, but there is a gap in the 0.19µm to 0.4µm region (except for the broad band spectral
795 measurements of ISS). This is an important wavelength region for non-water ice absorptions, and the
796 Cassini UVIS data in the far UV provide a critical tiepoint between the far UV and the visible band that
797 is helpful in understanding the nature of the non-ice species present in the surfaces of the moons. The
798 icy moons are dark in the far UV compared to their spectra in the Vis-NIR. For instance, the disk-
799 integrated albedo of Enceladus at 0.19µm is ~0.4, while at 0.439µm (Verbiscer et al., 2005) it is ~1.2.
800 Water ice is thought to have a relatively flat spectral signature in this wavelength range, which suggests
801 that some other species darkens Enceladus in the FUV. Ammonia, which exhibits an abrupt absorption
802 edge at 0.2µm (Pipes et al., 1974), is a candidate (and ammonia hydrate has been tentatively detected at
803 Enceladus and Tethys in Earth-based spectra (Verbiscer et al., 2005; Verbiscer et al., 2008)), though at
804
805
this time other species cannot be ruled out other materials.
806 Dark Material
807 Cassini (1672) was the first to infer the presence of dark material in the Saturnian system, which
808 was verified by Murphy et al. (1972) and Zellner (1972). Dark material appears throughout the
809 system, most predominantly on Dione, Rhea, Hyperion, Iapetus and Phoebe. The nature of the dark
810 material especially on Iapetus has been studied by numerous authors, often with conflicting
811 conclusions, including Cruikshank et al. (1983), Wilson and Sagan (1995, 1996) Vilas et al. (1996),
812
813
Jarvis et al. (2000), Owen et al. (2001), Buratti et al. (2002) and Vilas et al. (2004).
The dark material tends to darken and redden the satellites’ spectra, particularly in the visible. In
814 dark-region spectra, additional absorptions are seen at 2.42µm and 4.26µm, with the latter due to
815
816
CO2, as well as stronger water ice absorption near 3µm (Fig. 21.5.1d). The 2.42µm feature was first
tentatively identified as a CN overtone band by Clark et al. (2005) in spectra of Phoebe, but
817 evidence of CN fundamentals in the 4µm to 5µm region has not been confirmed. At present, the
818 2.42µm feature is believed to be due to OH or hydrogen (Clark et al., 2008a, b, 2009), but
819 complications with the VIMS calibration have limited any detailed definition of the absorption.
820 Diverse spectral features of the dark material on Dione match those seen on Phoebe, Iapetus, Hyperion
821 and Epimetheus as well as in the F-ring and the Cassini Division, implying that its composition is the
822 same throughout the Saturnian system (Clark et al., 2008a, 2009). Models of the dark material on
823 Iapetus have used tholins to explain the reddish visible spectral slope (e.g. Cruikshank et al., 2005).
824 However, VIMS data, which better isolate the spectral signature of the dark material from that of icy
825 regions thanks to the high spatial resolution afforded by close fly-bys, show that the dark material has
826 a remarkably linear spectrum (Clark et al., 2008a, 2009) (Fig. 12) not explained by tholins which have
827 curved spectral response. Tholins also have strong CH absorbers in the 3-4-μm region as well as other
828
marked absorptions at wavelengths longer than 2 μm that are not seen in, for example, spectra of the
16

829 dark regions of Iapetus. The problem of matching one spectrum might be accomplished with a
830 specific mixture of many compounds, but with spatially resolved spectra on the satellites we now
831 see a large range of mixtures, so models must show consistency across that range of observed
832 mixtures, and many models do not.
833 The source of the dark material in the Saturnian system is unclear. Several researchers (e.g. Clark et
834 al., 2005) suggest a source outside the system, perhaps the Kuiper belt or comets. Owen et al.
835 (2001) suggested that an impact on Titan created a spray of dark, organic-rich material, part of
836 which was swept up by Iapetus’ leading hemisphere, another part being transported to the inner
837 system to accrete onto the surfaces of the inner icy moons. The Dione VIMS spectra provide a key to
838 solving some of the mysteries of the dark material. Clark et al. (2008a) showed that a pattern of
839 bombardment by fine particles below 0.5 μm impacted Dione from the trailing-side direction. Several
840 lines of evidence point to an external origin of the dark material on Dione, including its global spatial
841 pattern, local patterns including crater and cliff walls shielding implantations on slopes facing away
842 from the trailing side, exposing clean ice, and slopes facing the trailing direction that show higher
843 abundances of dark material.
844 A blue scattering peaks is seen in the spectrum of the dark material and was initially observed in the
845 spectra of Tethys, Dione and Rhea by Noland et al. (1974). Clark et al. (2008a) suggest that the
846 blue peak is caused by particles less than 0.5μm in diameter embedded in the ice. The VIMS resolved
847 areas on Phoebe, Iapetus and Hyperion also show blue peaks (Fig. 9 c, d, e), indicating that this is a
848 common property of satellites in the Saturnian system (Clark et al., 2008a, 2009). The blue scattering
849
850
851
peak with a strong UV/visible absorption was modelled with ice plus 0.2 μm diameter carbon grains
(to provide the Rayleigh scattering effect) plus nano-phase hematite (Fe2O3) to produce the UV
absorber (Fig. 11). But carbon is inconsistent with the linear-red slope seen in the purest dark material
852 spectra from 0.4µm to 2.5µm. Metallic iron is the only known compound to display this spectral
853 property and nano-phase metallic iron plus nano-phase hemtatite mixed with ice show consistency
854
855
with the mixture spectra observed on the icy satellites and in Saturn’s rings (Clark et al., 2009).
856 Fig. 11 Spectrum of Dione compared to a spectrum of silicon carbide. (b) Spectrum of Dione
857 compared to the spectrum of a laboratory mixture of 99.25wt% H2O ice, 0.5wt% carbon black, and
858 0.25wt% nano-crystalline hematite (reagent by Sigma-Aldrich). The carbon black grains are 0.2µm
859 in diameter. Both the carbon black and hematite grains create Rayleigh scattering in the sample, but
860 at wavelengths shorter than about 0.5µm, absorption by hematite causes a downturn similar to that
861
862
seen in the Dione spectra (after Clark et al., 2008a).
863 Previous models of the dark material on Iapetus typically used tholins (e.g. Cruikshank et al., 2005).
864 However, VIMS data, which better isolate the spectral signature of the dark material from that of icy
865 regions thanks to the high spatial resolution afforded by close fly-bys, show that the dark material has
866
867
a remarkably linear spectrum (Clark et al., 2008a, 2009) (Fig. 12).
868 Fig. 12 Spectrum of the dark material of Iapetus compared to laboratory spectra of mixtures of iron
869
870
powder, H2O, CO2 and NH3. The small Rayleigh scattering peak and UV absorber can be explained
by traces of iron oxide nano-powder. The 3-micron absorption is best explained by a combination of
871
872
873
bound water, ice and trace ammonia. Finer-grained iron would require higher abundances of water,
ammonia and carbon dioxide to match the spectrum of Iapetus.
874 Tholins have strong CH absorbers in the 3-4-μm region as well as other marked absorptions at
875 wavelengths longer than 2μm that are not seen in, for example, spectra of the dark regions of Iapetus.
876 Using VIMS data, Clark et al. (2008a) found the general composition of the dark material in the
877
878
Saturnian system likely contains bound water, CO2 and, tentatively, ammonia. More recently, Clark et
al. (2008b, 2009) argued that the signature of iron identifies the major coloring agent on the icy
879 surfaces in the Saturnian system. Metallic iron does not exhibit sharp diagnostic spectral features,
880 but it does have unique spectral properties in that its absorption increases linearly with decreasing
881 wavelength (Clark et al., 2009).
882 The theory of Clark et al. (2008b, 2009) envisages small-particle meteoroid dust contaminating the
883 system as well as meteoroids with a high metallic-iron content. As one moves closer to the rings,
884 the dark material diminishes (e.g. Mimas and Enceladus have little dark material). However, the
885 competing process of accretion by Saturn’s E-ring, which originates on Enceladus, also contributes
886 to the greater abundance of high-albedo material on Mimas and Enceladus (Buratti et al. 1990,
887
Vebiscer et al., 2007), The main rings do not show any strong signature of metallic iron particles,
17

888 but the Cassini Division displays spectra that closely match those of Iapetus in Iapetus’ transition
889 zone from dark to light material, implying mixtures of the same materials. The ring spectra are
890 much redder in the 0.3µm to 0.6µm region than the spectra of dark material on Phoebe, Iapetus and
891 Hyperion. Cosmic rays impacting the huge surface area of the rings break up water molecules, thus
892 creating an atmosphere of hydrogen and oxygen. The hydrogen escapes, leaving an atmosphere of
893 oxygen around the rings (Johnson et al., 2006). Oxygen ions are very reactive, and when they
894
895
encounter iron particles, they oxidize them and turn them into nano-phase hematite (Fe2O3). The
spectra of nano-phase hematite mixed with ice and dark particles closely match the UV red slope
896 seen in the spectra of Dione (Clark et al., 2008a) (Fig. 11). To explain the dark material in the
897 Saturnian system, Clark et al. (2008b, 2009) suggest that oxidized sub-micron iron particles from
898 meteorites, mixed with ice, might explain all the colors and spectral shapes observed in the system.
899 A small amount of carbon might also be identified in the mixture, which would change the slope of
900 iron and further lower the albedo. Sub-micron particles also explain the origin of the observed
901
902
Rayleigh scattering. Trace amounts of CO2, ammonia and organics would not significantly
influence the visible and UV spectra down to 0.3µm. Finally, the inner satellites may be exposed to
903 as much iron as the outer satellites, but E-ring ice particles mix with and cover up the iron particles.
904 Species found by INMS, CAPS and MIMI in the magnetosphere are candidates for surface
905 components as the magnetospheric material originates either directly from satellite surfaces or from
906 ring particles. Compositional analyses of Enceladus’ plume are described in Spencer et al. (2009).
907 The idea of iron in the Saturn system is supported by the Cassini CDA instrument, which has
908
909
910
discovered small (

947 late 2004 suggest external material emplacement (e.g. dark material on ram-facing crater walls at
948 high latitudes) (Porco et al., 2005a), but this was later found to be due to photometric effects (Denk
949 et al., 2008). The initial theory of an exogenically-created dark pattern (Cook and Franklin, 1970)
950 suggested that pre-existing dark material was uncovered by meteoritic bombardment; this idea was
951 extrapolated by Wilson and Sagan (1996). Researchers theorized (Soter, 1974) that the dark
952 material is exogenically emplaced on Iapetus’ leading hemisphere as material is lost from the moon
953 Phoebe (Burns, 1979). Retrograde Phoebe dust from a distance of 215 Saturn radii would travel
954 inward from the effects of Poynting-Robertson drag and impact the leading hemisphere of Iapetus
955 orbiting at 59 Saturn radii. However, Phoebe is spectrally neutral at visible wavelengths, while the
956 Iapetus dark material is reddish (Cruikshank et al., 1983; Squyres et al., 1984). If the material does
957 come from Phoebe, then some sort of chemistry or impact volatilization must occur to change the
958 color and darken the material (Cruikshank et al., 1983; Buratti and Mosher, 1995). Another
959 possibility is that the exogenic source of the dark material is Hyperion (Matthews, 1992; Marchi et
960 al., 2002), Titan (Wilson and Sagan, 1995; Owen, 2001), Iapetus itself (Tabak and Young, 1989) or
961 the debris remnants of a collision between a former outer moon and a planetocentric object (Denk
962 and Neukum, 2000). Matthews (1992) suggested that the hypothetical impact that disrupted
963 Hyperion created a debris cloud that subsequently hit Iapetus. Both Hyperion and Titan dark
964 materials are spectrally reddish (Thomas and Veverka, 1985; McDonald et al., 1994; Owen et al.,
965 2001), though not as dark as those of Iapetus. Buratti et al. (2002; 2005) suggested that both
966 Hyperion's and Iapetus’ leading hemispheres are impacted by dark, reddish dust from recently
967
968
discovered retrograde satellites exterior to Phoebe. Clark et al. (2008a) offer a solution for the color
discrepancy. Red, fine-grained dust (less than one-half micron in diameter) when mixed in small
969 amounts with water ice produces a Rayleigh scattering effect, enhancing the blue response. This
970 blue enhancement appears just enough on Phoebe to create an approximately neutral spectral
971 response. An examination of the spectral signatures of Dione, Hyperion, Iapetus and Phoebe by
972 Clark et al (2008a) showed a continuous mixing space on all these satellites, demonstrating variable
973 mixtures of fine-grained dark material in the ice. Thus, the colors observed are simply variable
974 amounts of dark material, with the redder slopes indicating greater abundance.
975 Ground-based radar observations at 13cm (Black et al., 2004) and Cassini RADAR data at 2.2cm
976 (Ostro et al., 2006) indicate that the dark terrain on Iapetus must be quite thin (one to several
977 decimeters); an ammonia-water ice mixture may be present at a depth of several decimeters below
978 the surface on both the leading and trailing hemispheres of Iapetus. The RADAR results appear to
979 rule out any theories of a thick dark material layer (Matthews, 1992; Wilson and Sagan, 1996) and
980
981
are consistent with the spectral coating indicated by VIMS data as described by Clark et al (2008a).
3
Cassini radio science results (Rappaport et al., 2005) indicate a bulk density for Iapetus of 1.1g/cm ,
982 from which it can be inferred that the moon is composed primarily of water ice.
983 Because the large craters within the dark terrain appear evenly colored by the dark material in
984 Voyager data, no craters significantly break up the dark material to expose bright underlying terrain
985 (Denk et al., 2000). This suggests that the emplacement of dark material by whatever mechanism is
986 relatively new or ongoing.
987 With the arrival of Cassini at Saturn and the first Iapetus flyby (Dec 2004), the idea of thermal
988 segregation developed (Spencer et al., 2005; Hendrix and Hansen, 2008a). This idea was originally
989
990
mentioned as one option by Mendis and Axford (1974). The September 2007 close flyby of Iapetus
permitted testing of this hypothesis, resulting in the following theory. A small amount of exogenic
991 dust from an unknown source created a 'global color dichotomy': the leading side became slightly
992 redder and darker than the trailing with fuzzy boundaries close to the sub-Saturn and anti-Saturn
993 meridians (Denk et al., 2008). A runaway thermal segregation process might have been initiated
994 primarily at the low and middle latitudes of the leading side as well as in some local areas at low
995 latitudes on the trailing side, especially on equatorward-facing slopes (Denk et al., 2008). Indeed,
996 poleward-facing slopes even within the dark terrain (at middle latitudes) are bright. On these slopes
997 the temperature never raises enough to allow the thermal process to act faster than the destruction
998 process by micrometeorite gardening, which transports bright sub-surface material to the surface. In
999 those locations that are now dark, the originally bright surface material becomes warm enough for
1000 volatiles (mainly water ice) on the top part of the surface to grow unstable and migrate towards cold
1001 traps, leaving behind residual dark material that was originally intimately mixed with the water ice
1002 as a minor constituent. The slow 79.3day rotation is a crucial boundary condition for this process to
1003 work in the Saturnian system. ISS images suggest that significant cold traps might be located at
1004
1005
high latitudes on the trailing side, where lower-resolution images show bright white polar caps.
19

1006 Compositional data from the UVIS and the VIMS (Hendrix and Hansen, 2009a; Hendrix and
1007 Hansen, 2009b; Clark et al., 2009) indicate variable amounts of ice and volatiles in the dark
1008 material. The water-ice absorption edge in the UV at 0.165µm was shown to increase with latitude,
1009 consistent with decreasing temperatures and increasing numbers of visibly bright regions (Hendrix
1010 and Hansen, 2009a). This is consistent with thermal segregation as a possible process explaining the
1011 overall appearance of Iapetus and, more particularly, the observed correlations of the dark/bright
1012 boundary with geographical location and topography (Fig. 15) and the small-scale dramatic
1013
1014
brightness variations (Fig. 15).
1015 Fig. 14 Cassini ISS images from the September 2007 flyby, demonstrating the large- and small-scale
1016 albedo effects of thermal segregation. (Left) The central longitude of the trailing hemisphere is 24°
1017 to the left of the mosaic's center. (Right) The mosaic consists of two image footprints across the
1018 surface of Iapetus. The view is centered on terrain near 42° southern latitude and 209.3° western
1019 longitude, on the anti-Saturn facing hemisphere. Image scale is approximately 32m/pixel; the crater
1020
1021
in the center is about 8 km across.
1022 The 3µm absorption observed in VIMS data (Clark et al., 2009) might be due to ice, ammonia or
1023 bound water. As the 3µm absorption is not been well matched by those due to ice, bound water and
1024 ammonia, a combination seems to be indicated, but a model showing a good match has not yet been
1025 developed (Clark et al., 2009). The amount of ice contained in the dark materials from VIMS data
1026
1027
1028
could be on the order of 10%, but is estimated to be

1065 visible wavelength by the Hubble space telescope (Verbiscer et al., 2007). Their brightness changes
1066 with their distance from the E-ring, which is consistent with the assumption that their albedo results
1067 from E-ring coating. Hamilton and Burns (1994) predicted that the relative velocities between the
1068 E-ring particles and the icy satellites might explain the large-scale longitudinal albedo patterns on
1069 the icy satellites. Because of their higher relative velocity, the satellites exterior to the densest point
1070 in the E-ring have brighter leading hemispheres, while those interior to it have brighter trailing
1071 hemispheres because of the higher velocity of the particles. Indeed, Mimas and Enceladus are both
1072 slightly darker on their leading than on their trailing hemispheres, while the leading hemispheres of
1073 Tethys, Dione and Rhea are brighter than their trailing hemispheres (Buratti et al., 1998). The
1074 visible orbital phase curve of Enceladus shows that the leading hemisphere is ~1.2 times darker than
1075 its trailing hemisphere. At visible wavelengths, the leading hemisphere of Tethys is ~1.1 times
1076 brighter than its trailing hemisphere, Dione’s leading hemisphere is up to ~1.8 times brighter and
1077 Rhea’s leading hemisphere is ~1.2 times brighter than the trailing hemisphere (Buratti and Veverka,
1078 1984).
1079 The bombardment of icy surfaces by charged particles causes both structural and chemical changes
1080 in the ice (Johnson et al., 2004). Effects include the implantation of magnetospheric ions, sputtering
1081 and grain size alteration. Low-energy electrons can cause chemical reactions on the surface, while
1082 high-energy electrons and ions tend to cause radiation effects at depth. The latter can result in bulk
1083 chemical changes in composition as well as structural damage. A primary effect of energetic
1084 particle bombardment on ice is to cause defects in it (Johnson, 1997; Johnson and Quickenden,
1085
1086
1997; Johnson et al., 1998). Extended defects, in the form of voids and bubbles, affect the lightscattering
properties of the surface. They also act as reaction and trapping sites for gases, which can
1087 produce spectral absorption features. Incident ions deposit energy over a small region around their
1088 path through the ice, causing local damage. Depending on excitation density and cooling rate, this
1089 can result in either local crystallization or amorphization (Strazzulla, 1998).
1090 Energetic ions and electrons have been seen to brighten laboratory surfaces in the visible by
1091 producing voids and bubbles (Sack et al., 1992). Depending on the size and density of these
1092 features, a clear surface can become brighter due to enhanced scattering in the visible. The plasma-
1093 induced brightening competes with annealing, which operates efficiently in regions where the ice
1094 temperature exceeds ~100K and can make equatorial regions less light-scattering despite the plasma
1095 bombardment.
1096 A significant effect of the ion bombardment of ice is the creation of new species through radiolysis.
1097
1098
Bombarding particles can decompose H2O ice, creating molecular hydrogen and oxygen. Miniatmospheres
of trapped volatiles have been suggested to form in voids in the ice (Johnson and
1099
1100
Jesser, 1997). As a result, species such as O3, H2O2, O2 and oxygen-rich trace species have been
detected as gases trapped in the regoliths of the icy Galilean satellites (Nelson et al., 1987; Noll et
1101 al., 1996; Hendrix et al., 1998, 1999; Spencer et al., 1995; Carlson et al., 1996, 1999; Hibbitts et al.,
1102 2000, 2003). On these satellites, the hydrogen and oxygen produced by radiation decomposition
1103 form tenuous atmospheres; in contrast, atmospheres have not been detected on satellites in the
1104 Saturnian system (with the exception of the gaseous plume of Enceladus). There are several
1105 possible reasons for this. Lower gravities might support the loss of volatiles; temperatures are
1106
1107
1108
lower, which might slow down the breakup of H2O and the loss of H; the particle environment is
different (discussed below); the electron density and temperature is not effective in exciting and
'lighting up' any gases around the moons; and sputtering is not as intense as in the Jupiter system
1109 due to differences in the net charge and energy of the particle environment. On Saturn, the primary
1110 surface alterations found thus far are due to the E-ring, dust (e.g. on Dione) and thermal processing
1111 (especially apparent on Iapetus, as previously discussed).
1112 Neutrals, rather than ions as on Jupiter, dominate the Saturnian system. The energetic heavy ions
1113 (~10keV to MeV) that are critical at Europa and Ganymede are lost by charge exchange to the
1114 neutral cloud of the Saturnian system as they diffuse inwards from beyond ~10Saturn radii
1115 (Paranicas et al., 2008). However, although the energy fluxes due to trapped plasma and the solar
1116 UV are relatively low on the icy Saturnian satellites (Paranicas et al., 2008), radiation processing
1117 appears to be occurring. Radiation fluxes are known to cause decomposition in ice, possibly
producing ozone in the icy surfaces of Dione and Rhea (Noll et al., 1997) and O2 + 1118
1119
in Saturn’s
plasma (Martens et al., 2007). In addition, the role of low energy ions has been recently re-
1120 examined using Cassini CAPS data (Sittler et al., 2008), indicating that sputtering by this
1121
1122
component is not negligible (Johnson et al., 2008). Spencer (1998) searched for but did not find O2
inclusions on Rhea and Dione. Such inclusions have been seen on Ganymede in the low latitudes of
1123
the trailing hemisphere, and it was suggested that they might be related to charged-particle
21

1124
1125
1126
bombardment or UV photolysis. The lack of O2 on Rhea and Dione suggests that low temperatures
preclude the formation of O2 bubbles in the ice. Hydrogen peroxide (H2O2) has been tentatively
VIMS data (Newman et al., 2007) of the south pole tiger-stripes region of Enceladus, although it is
1127 not clear whether exogenic processing is needed to create peroxide.
1128 It is important to distinguish between hemispheric dichotomies in color, albedo, texture and
1129 macroscopic structure because each exogenic alteration mechanism leaves a different signature on
1130 the physical and optical properties of the surface. Photometric methods have been employed over
1131 the past two decades to characterize the roughness, compaction state, particle size and phase
1132 function, and single scattering albedo of planetary and satellite surfaces, including the Saturnian
1133 satellites (Buratti, 1985; Verbiscer and Veverka, 1989, 1992, 1994; Domingue et al. 1995). An
1134 indication of the scattering properties of the satellites’ surfaces can be obtained by extracting
1135 photometric scans across their disks at small solar phase angles. Dark surfaces such as that of the
1136 Moon, in which single scattering is dominant, are not significantly limb-darkened, whereas bright
1137 surfaces in which multiple scattering is important could be expected to show limb darkening.
1138 Physical properties such as roughness, compaction state and particle size give clues as to the past
1139 evolution of the satellites’ surfaces and reveal the relative importance of exogenic alteration
1140 processes. Accurate photometric modeling is also necessary to define the intrinsic changes in albedo
1141 and color on a planetary surface; much, and in many cases most of the change in intensity is due to
1142 variations in the radiance viewing geometry and is not intrinsic to the surface.
1143 Planetary surface scattering models have been developed (e.g. Horak, 1950; Goguen, 1981; Lumme
1144
1145
and Bowell, 1981a, 1981b; Hapke, 1981, 1984, 1986, 1990; Shkuratov et al., 1999, 2005) which
express the radiation reflected from the surface in terms of the following physical parameters: single
1146 scattering albedo, single particle phase function, the compaction properties of the optically active
1147 portion of the regolith, and the scale and extent of macroscopically rough surface features.
1148 Observations at small solar phase angles (0-10°) are particularly important for understanding the
1149 compaction state of the upper regolith: this is the region of the well-known opposition effect, in
1150 which the rapid disappearance of shadowing among surficial particles causes a nonlinear surge in
1151 brightness as the object becomes fully illuminated. At the smallest phase angles (40°) are necessary for
1153 investigating the macroscopically rough nature of surfaces, ranging in size from clumps of several
1154 particles to mountains, craters and ridges. Such features cast shadows and alter the local incidence
1155 and emission angles. Two formalisms have been developed to describe roughness for planetary
1156 surfaces: a mean-slope model in which the surface is covered by a Gaussian distribution of slope
1157 angles with a mean angle of θ (Hapke, 1984; Helfenstein and Shepard, 1998), and a crater model in
1158 which the surface is covered by craters with a defined depth-to-radius ratio (Buratti and Veverka,
1159 1985). The single particle phase function, which expresses the directional scattering properties of a
1160 planetary surface, is an indicator of the physical character of individual particles in the upper
1161 regolith, including their size, size distribution, shape, and optical constants. Small or transparent
1162 particles tend to be more isotropically scattering because photons survive to be multiply scattered;
1163 forward scattering can exist if photons exit in the direction away from the observer. The single
1164 particle phase function is usually described by a 1 or 2-term Henyey-Greenstein phase function defined
1165 by the asymmetry parameter g, where g = 1 is purely forward-scattering, g = -1 is purely
1166
1167
backscattering, and g = 0 is isotropic. The opposition surge of airless surfaces has been described by
parameters that define the compaction state of the surface and the degree of coherent backscatter of
1168 multiply scattered photons (Irvine, 1966; Hapke, 1986, 1990). One final physical photometric
1169 parameter is the single scattering albedo (ώ), which is the probability that a photon will be scattered
1170 into 4π steradians after one scattering.
1171 Three other important physical parameters are the geometrical albedo (p), the phase integral (q) and the
1172
1173
Bond albedo (AB). The geometrical albedo is the flux received from a planet at a solar phase angle of
0º compared with that of a perfectly diffusing disk of the same size, also at 0º. The phase integral (q) is
1174 the flux, normalized to unity at a solar phase angle of 0º, integrated over all solar phase angles. The
1175 Bond albedo is p·q, and when integrated over all wavelengths it yields the bolometric Bond albedo,
1176 which is a measure of the energy balance on the planet or satellite. Except for the bolometric Bond
1177 albedo, all these parameters are wavelength-dependent.
1178 Buratti (1984) and Buratti and Veverka (1984) performed photometric analyses of visible-
1179 wavelength Voyager data of the icy Saturnian moons based on Voyager disk-resolved data sets, and
1180 disk-integrated Voyager observations at large phase angles to supplement existing ground-based
1181 observations at smaller phase angles. These studies investigated limb darkening with implications
1182
for multiple scattering (significant limb darkening suggests that multiple scattering is important).
22

1183 Physical photometric modeling of Saturnian moons based on Voyager and ground-based data was
1184 done by Buratti (1985), Verbiscer and Veverka (1989, 1992, 1994), Domingue et al. (1995) and
1185 Simonelli et al. (1999). In general, the surfaces of the Saturnian satellites are backscattering and
1186 exhibit roughness comparable to that of the Moon, indicating that impact processes produce similar
1187 morphological effects on rocky and icy bodies at low temperatures. The photometric and physical
1188 properties of the satellites determined so far from both Voyager and Cassini data are summarized in
1189
1190
Table 4.
Parameter Mimas Enceladus Tethys Dione Rhea Phoebe
Mean slope 30 (2)
13±5 (4) 31±4 (6)
angle θ (º)
Asymmetry
parameter g
30±1 (5) 6±1 (3)
33±3 (7)
Single
scattering
albedo ώ
Phase
integral q
Visual
geometric
albedo pv
Bond
albedo AB
-0.3±0.05 (2)
-0.213±0.005
(5)
0.93±0.03 (2)
0.951±0.002
(5)
0.82±0.05 (1)
0.78±0.05 (5)
-0.35±0.03
(2)
-0.399±0.005
(3)
0.99±0.02 (2)
0.998±0.001
(3)
0.86±0.1 (1)
0.92±0.05 (3)
0.77±0.15 (1) 1.04±0.15 (1)
1.41±0.03 (8);
0.6±0.1 (1)
0.5±0.1 (5)
0.9±0.1 (1)
0.91±0.1 (3)
0.7±0.1 (1)
0.80±0.15
(1)
0.6±0.1 (1)
23
0.8±0.1 (1)
0.55±0.15
(1)
0.45±0.1
(1)
-0.287±0.007 (4) -0.24 (6)
-0.3±0.1 (plus
a forward
component)
(7)
0.861±0.008 (4) 0.068 (6)
0.07±0.01 (7)
0.70±0.06 (1) 0.24±0.05 (6)
0.78±0.05 (4) 0.29±0.03 (7)
0.65 (1) 0.081±0.002
(6);
0.082±0.002
(7);
0.45±0.1 (1) 0.020±0.004
0.49±0.06 (4) (6)
0.023±0.007
(7)
1191
1192 Tab. 4 Photometric and physical parameters of the icy Saturnian satellites (references: (1) Buratti
1193 and Veverka, 1984; (2) Buratti, 1985; (3) Verbiscer and Veverka, 1994; (4) Verbiscer and Veverka
1194 1989; (5) Verbiscer and Veverka, 1992; (6) Simonelli et al., 1999; (7) Buratti et al., 2008; (8)
1195 Verbiscer et al., 2005
1196
1197 Opposition surge data from 0.2 to 5.1µm were obtained for Enceladus, Tethys, Dione, Rhea and
1198
1199
Iapetus. Fig. 16 shows the opposition curve of Iapetus in between 0.35µm and 3.6µm
1200 Fig. 15 (Above) Cassini visual infrared mapping spectrometer (VIMS) disk-integrated observations
1201 of the opposition surge on Iapetus between 0.35 and 3.6µm. (Lower left) The slope of the solar phase
1202 curve (the 'phase coefficient') between 1º and 8º decreases as the albedo increases. This effect is
1203 expected for shadow hiding, since bright surfaces have partly illuminated primary shadows. (Lower
1204 right) On the other hand, the amplitude of the 'spike' under 1º (shown as a fractional increase above
1205 the value of the phase curve at 1º) increases with the albedo, suggesting that it is a multiple scattering
1206
1207
effect, such as coherent backscatter.
1208 (Buratti et al., 2009b; Hendrix and Buratti, 2009). The curve is divided into two components: at
1209 phase angles greater than 1º, brightness increases linearly with decreasing phase angle, while at
1210 phase angles smaller than one degree the increase is exponential. Moreover, the albedo-dependence
1211 of the two types of surge is different. For the first type, the magnitude of the surge decreases as the
1212 albedo increases, while for the second type of surge the reverse is true. This result suggests that
1213 different phenomena are responsible for the two types of surge. At angles greater than 1º, the surge
1214
1215
is caused by mutual shadowing, in which shadows cast among particles of the regolith rapidly
disappear as the face of the satellite becomes fully illuminated to the observer. Since multiply
1216 scattered photons, which partly illuminate primary shadows, become more numerous as the albedo
1217 increases, this effect becomes less pronounced for bright surfaces. On the other hand, the huge
1218 increase at very small solar phase angles can be explained by coherent backscatter, a phenomenon
1219 in which photons following identical but reversed paths in a surface interfere constructively in exactly
1220
the backscattering direction, causing brightness to increase by up to a factor of two (Hapke, 1990).

1221 The Cassini flyby in June 2004 prior to Saturn orbit insertion afforded views at large phase angles –
1222 important for modeling roughness - and over a full excursion in geographical longitudes. Phoebe
1223 exhibits a substantial forward-scattering component to its single particle phase function. This result
1224 is consistent with a substantial amount of fine-grained dust on the surface of the satellite generated
1225 by particle infall, as suggested by Clark et al. (2005, 2008a). As observed in Cassini images (Porco
1226 et al., 2005a), Phoebe also shows extensive macroscopic roughness, hinting at the violent collisional
1227 history predicted by Nesvorny et al. (2003).
1228 ISS images of the low-albedo hemisphere of Iapetus were fitted to the crater roughness model
1229 (Buratti and Veverka, 1985), yielding a markedly smooth value for the degree of macroscopic
1230 roughness on this side with a mean slope angle of only a few degrees (Lee et al., 2009). This result
1231 suggests that small-scale rough features on the dark side have been filled in (features which
1232 probably dominate the photometric effects; see Helfenstein and Shepard, 1998), which is consistent
1233 with the idea of Spencer and Denk (2009) that thermal migration of water ice leaves behind a fluffy
1234 residue of dark material. The low-albedo side of Iapetus shows a roughness similar to Enceladus,
1235 which is coated with micron-sized particles from its plume and the E-ring (Verbiscer and Veverka,
1236
1237
1994; see Table 21.4.1). Model-fits for the low albedo hemisphere are shown in Fig. 14.
1238 Fig. 16 A macroscopic roughness model for the low-albedo hemisphere of Iapetus, based on the
1239 crater roughness model by Buratti and Veverka (1985). The best-fit roughness function corresponds
1240 to a depth-to-radius of 0.14, or a mean slope angle of 11º, which is much lower than the ~30º typical
1241
1242
1243
of other icy satellites (see Tab. 4). For these fits, a surface phase function f(α) of 0.023 was derived.
From Lee et al., 2009
1244 Finally, observations at very large solar phase angles (155º-165º) have been used to search for
1245 cryovolcanic activity on satellites. By comparing the brightness of Rhea to Enceladus in this phase
1246 angle range at 2.02µm, Pitman et al. (2008) were able to place an upper limit on water vapor
column density based on a possible plume of 1.52 x 10 14 to 1.91 x 10 15 cm -2 1247 , two orders of
1248
1249
magnitude below the observed plume density of Enceladus (Fig. 17).
1250 Fig. 17 Cassini VIMS disk-integrated brightness as a function of solar phase angle at 2.23µ.
1251 Symbols: Normalized Rhea and Enceladus VIMS brightness (every fifth data point plotted). Solid
1252 line: Third-order polynomial fit to Rhea data. Dashed lines: (polynomial fit to Rhea data).
1253 Enceladus's plume can be seen as a peak at a solar phase angle of 159°. Adapted from Pitman et al.
1254
1255
(2008)
1256 The microphysical structure of the E-ring grain coatings of the inner icy moons embedded in the E-
1257 ring can be probed by investigating the opposition surge. Verbiscer et al. (2007) found that the
1258 amplitude and angular width of the opposition effect on the inner icy moons are correlated with the
1259
1260
moons’ position relative to the E-ring.
1261 21.5 Constraints on the Top-Meter Structure and Composition by Radar
1262 The wavelength of Cassini's 13.8GHz RADAR instrument, 2.2cm, is about six times shorter than
1263
1264
the only ground-based radar wavelength available to study the satellites (13cm, at Arecibo) and 22
times longer than the millimeter wavelengths at the limit of the CIRS.
1265 The echoes result from volume scattering, and their strength is sensitive to ice purity. Therefore,
1266 they provide unique information about near-surface structural complexity as well as about
1267 contamination with non-ice material. This section summarizes the most basic aspects of Cassini-
1268 and ground-based radar observations of the satellites, the theoretical context for interpreting the
1269 echoes, and the inferences drawn about subsurface structure and composition.
1270 Cassini measures echoes in the same linear polarization as transmitted, whereas most ground-based
1271 radar astronomy consists of Arecibo 13cm or 70cm observations and Goldstone 3.5cm observations
1272 that almost always use transmission of a circularly polarized signal and simultaneous reception of
1273 echoes in same-circular (SC) and opposite-circular (OC) polarizations. A radar target's radar albedo
1274 is defined as its radar cross section divided by its projected area. The radar cross section is the
1275 projected area of a perfectly reflective isotropic scatterer which, if observed at the target's distance
1276 from the radar and using the same transmitted and received polarizations, would return the observed
1277 echo power. The total-power radar albedo is the sum of the albedos in two orthogonal polarizations,
1278 TP = SL + OL = SC + OC. Here, we will use 'SL-2' to denote the 2cm radar albedo in the SL
1279
polarization, 'TP-13' to denote the 13cm total-power albedo, etc.
24

1280 For a smooth, homogeneous target, single back reflections would be entirely OC when the
1281 transmission is circularly polarized, or entirely SL when the transmission is linearly polarized, so
1282 the polarization ratios SC/OC and OL/SL would equal zero. Saturn's icy satellites and the icy
1283 Galilean satellites not only have radar albedos an order of magnitude greater than those of the Moon
1284 and most other solar system targets but also unusually large polarization ratios.
1285 These strange radar signatures were discovered in observations of Europa, Ganymede and Callisto
1286 in the mid-1970s. It took a decade and a half to realize that the signatures are the outcome of
1287 'coherent backscattering,' which is phase-coherent, multiple scattering within a dielectric medium
1288 that is both heterogeneous and non-absorbing (e.g. MacKintosh and John, 1989; Hapke, 1990).
1289 During the next decade, Peters (1992) produced a vector formulation of coherent backscattering that
1290 led to predictions of radar albedos and polarization ratios, and Black (1997) and Black et al. (2001b)
1291 built on Peters' work with models of scatterer properties.
1292 Extremely clean ice is insufficient for anomalously large radar albedos and circular polarization
1293 ratios, which arise from multiple scattering within a random, disordered medium whose intrinsic
1294 microwave absorption is very low. Thus, perfectly clean ice that is structurally homogeneous
1295 (constant density at millimeter scales) would not give anomalous echoes. The regoliths on solar
1296 system objects are structurally very heterogeneous, with complex density variations from the
1297 distribution of particle sizes and shapes produced by impacts of projectiles of different sizes and
1298 impact energies.
1299 It is to be expected that the uppermost few meters of Saturn's icy satellite regolith are basically
1300
1301
similar to those of the icy Galilean satellites or the Moon's in terms of particle size distribution,
porosity and density as functions of depth, and lateral/vertical heterogeneity. It is the combination
1302 of naturally complex density variations and the extreme transparency of pure water ice that gives
1303 icy regolith its exotic radar properties.
1304 The icy Galilean satellites were observed in dual-circular-polarization experiments from the ground
1305 at 3.5, 13, and 70cm; experiments measuring SL-13 and OL-13 were done only for Europa,
1306 Ganymede and Callisto. Arecibo obtained estimates of SC-13 and OC-13 for Enceladus, Tethys,
1307 Dione, Rhea and Iapetus. The Cassini radar measured SL-2 for Mimas, Hyperion and Phoebe. The
1308 satellites' albedos show different styles and are location- and wavelength-dependent, as discussed
1309 below. The results described in this chapter make use of 73 Cassini SL-2 measurements, more than
1310 twice the number reported by Ostro et al. (2006).
1311 If a radar target is illuminated by a uniform antenna beam (that is, if the antenna gain is constant
1312 across the disk), then it is a simple matter to use the appropriate 'radar equation' to convert the
1313 measured echo power spectrum (power as a function of Doppler frequency) into a radar cross
1314 section (e.g. Ostro, 1993). For Cassini measurements, except for the SAR imaging of Iapetus
1315 discussed below, the antenna beam width is within a factor of several times the target's angular
1316 width, so the gain varies across the target's disk, usually significantly. Data reduction yields an
1317 estimate of echo power, which is the sum of contributions from different parts of the surface, with
1318 each contribution weighted by the two-way gain. However, the distribution of echo power as a
1319 function of location on the target disk is not known a priori. Thus, we do not know how much the
1320 echo power from any given surface element has been 'amplified' by the antenna gain. To estimate a
1321 radar albedo from our data we must make an assumption about the homogeneity and angular
1322
1323
dependence of the surface's radar scattering. For the purpose of obtaining a radar albedo estimate
from an echo power spectrum, one assumes uniform, azimuthally isotropic, cosine scattering laws,
1324 and uses least squares to estimate the value of a single parameter that defines the echo's incidence-
1325
1326
angle dependence. In general, only modest departures from Lambert limb darkening are found.
1327 Radar-Optical Correlations
1328
1329
Fig. 18 shows all measurements of the satellites' individual SL-2 albedos along with their
1330 Fig. 18 Individual Cassini measurements of SL-2 (small symbols) and average optical geometrical
1331 albedos, plotted for the satellites vs. their distances from Saturn. Optical albedos of Mimas,
1332 Enceladus, Tethys, Dione and Rhea are from Verbiscer et al. (2007); those of Iapetus and Phoebe
1333 are from Morrison et al. (1986) and involve a V filter; Hyperion’s is from Cruikshank and Brown
1334 (1982). Note the similar distance dependences. Hyperion's radar albedo is large and appears to
1335 break the pattern, perhaps (Ostro et al. 2006) because Hyperion's optical coloring may be due to a
1336 thin layer of exogenously derived material that, in contrast to the other satellites, has never been
1337
1338
worked into the icy regolith, which remains very clean and radar-bright.
25

1339 average optical geometrical albedos, plotted vs. distance from Saturn. A very pronounced radar-
1340 optical correlation means that variations in optical and radar albedos involve the same or related
1341 contaminant(s) and/or surface processes. Absolutely pure water ice at satellite surface temperatures
1342 is extremely transparent to radar waves. Contamination of water ice with no more than trace
1343 concentrations of virtually any other substance (including earth rocks, chondritic meteorites, lunar
1344 soil, ammonia, metallic iron, iron oxides and tholins) can decrease its radar transparency (and hence
1345 the regolith's radar albedo) by one to two orders of magnitude. Phenomena that depend on the
1346 distance from Saturn include primordial composition, meteoroid flux and the influx of dark material
1347 (perhaps from several sources), the balance between the influx of ammonia-containing particles and
1348 the removal of ammonia by a combination of ion erosion and micrometeoroid gardening, variation
1349 in E-ring fluxes, and the systematics of thermal volatiles redistribution.
1350 Modeling of icy Galilean satellite echoes (Black et al., 2001b) suggests that penetration to about
1351 half a meter should be adequate to produce the range of measured values of SL-2.
1352 The explanation by Clark et al. (2008b, 2009) for dark Saturnian system material, which suggests
1353 that it consists of sub-micron iron particles from meteorites which oxidized and became mixed with
1354 ice, would also suffice to explain, at least qualitatively, the overall radar-optical albedo correlation
1355 in Fig. 18 if the contamination of the ice extends some one to several decimeters below the optically
1356 visible surface, as expected by Clark et al. (2008b, 2009). The tentative detection of ammonia in
1357 VIMS spectra by Clark et al. (2008a) opens up the possibility that contamination by this material
1358 may play an important role in inter- and intra-satellite radar albedo variations. If so, this
1359
1360
contamination would affect the radar but not the optical albedo.
Almost all Cassini radar data on the icy satellites consist of real-aperture scatterometry, which
1361 means that the beam size determines the resolution. Cassini has no real-aperture resolution of
1362 Mimas, Enceladus, Hyperion or Phoebe, but does have abundant measurements for Tethys, Dione,
1363 Rhea and Iapetus with beams smaller than, and in many cases less than half, the size of the disk
1364 (Ostro et al., 2006).
1365 The SL-2 distributions of the individual satellites can be summarized briefly as follows. As regards
1366 Mimas, Herschel may or may not contribute to the highest value of the disc-integrated albedo.
1367 Enceladus shows no dramatic leading/trailing asymmetry and no noteworthy geographical
1368 variations except for a low albedo in an observation of the Saturn-facing side. Tethys' leading side
1369 is brighter both optically and in SL-2. Dione's highest-absolute-latitude views give the highest radar
1370 albedos. Rhea's leading side has a lower SL-2 in the north, where there is much cratering, and
1371 higher in the south, where there are optically bright rays. Its trailing and Saturn-facing sides, which
1372 look younger, show a relatively high SL-2. Iapetus displays an extremely marked correlation
1373 between optical brightness and SL-2.
1374 In general, a first-order understanding of the SL-2 distribution is that higher SL-2 means cleaner
1375 near-surface ice, but the nature of the contaminant(s) and the depth of the contamination are open
1376 questions. The latter can be constrained by measuring radar albedos at more than one wavelength.
1377 Fortunately, there are Arecibo 13-cm observations of Enceladus, Tethys, Dione, Rhea (Black et al.,
1378 2004) and Iapetus (Black et al., 2007).
1379 Also fortunately, the wavelength dependence observed for Europa, Ganymede and Callisto (EGC)
1380 gives us a context for discussing the wavelength dependence observed for Saturn's satellites.
1381
1382 Wavelength Dependence
1383 EGC total-power radar albedos and circular polarization ratios at 3.5cm and 13cm are
1384 indistinguishable. This means that their regolith structure (distribution of density variations) and ice
1385 cleanliness look identical at the two wavelengths, for each object. That is, near-surface structures
1386 are 'self-similar' at 3.5cm and 13cm scales; there are no depth-dependent phenomena, like
1387 contamination levels or structural character, perceptible in EGC's 3.5cm and 13cm radar signatures.
1388 At 70cm, moreover, EGC SC/OC ratios show little wavelength dependence, so subsurface multiple
1389 scattering remains the primary source of the echo even on this longer scale. However, the radar
1390 albedos (TP-70) plummet for all three objects. The implication is that at 70cm, scattering
1391 heterogeneities are relatively sparse, and the self-similarity of the regolith structure disappears on
1392 the larger scale. On Europa at least, with the most striking albedo drop from TP-13 to TP-70, the
1393 regolith may simply be too young for it to be heterogeneous at 70cm.
1394 Within the framework of the Black et al. (2001b) model, this implies that at 70cm 'the scattering
1395 layer is disappearing in the sense that there are too few scatterer at this scale and they are confined
1396 to a layer too shallow to permit numerous multiple scatterings. At even longer wavelengths, of the
1397
order of several meters, the reflectivity from the surfaces of these moons, and Europa in particular,
26

1398 may drop precipitously or even vanish as the layers responsible for the centimeter-wavelength radar
1399 properties become essentially transparent.'
1400 Now, given that the 3.5cm and 13cm polarization properties of EGC show no significant
1401 wavelength dependence, and that the 13cm OL/SL ratios of those objects are very similar, let us
1402 assume that the linear polarization properties of the Saturnian satellites are wavelength-invariant.
1403 Let us then use EGC's mean OL/SL value to estimate total-power 2cm albedos: TP-2 = (1.52 +/-
1404 0.13) SL-2 and compare them to the Arecibo 13cm total-power albedos (TP-13).
1405
1406
Fig. 19 shows relevant Cassini and Arecibo information.
1407 Fig. 19 Wavelength dependence of total-power radar albedos. Individual symbols are estimates of
1408 TP-2 obtained from individual measurements of SL-2 using TP-2 = 1.52 SL-2 +/- 0.13 SL. Arecibo
1409
1410
13-cm albedos are horizontal lines. Dashed lines represent EGC.
1411 Moving outward from Saturn, the magnitude of 2cm to-13cm wavelength dependence is
1412 unremarkable for Enceladus' leading side, extreme for Enceladus' trailing side, very large for
1413 Tethys, large for Dione, unremarkable for Rhea, large for Iapetus' leading side, and extreme for
1414 Iapetus' trailing side.
1415 Whereas compositional variation (ice cleanliness) almost certainly does not contribute to the
1416 wavelength dependence seen in EGC, it may well play a key role in the wavelength dependence
1417 seen in Saturn's satellites. In this complicated system, the tapestry of satellite radar properties may
1418
1419
involve many factors, including inter- and intra-satellite variations in E-ring flux, variations in the
meteoroid or ionizing flux, variations in the concentration of radar absorbing materials (including
1420 optically dark material and possibly ammonia) as a function of depth, the systematics of thermal
1421 volatile redistribution (especially for Iapetus), and/or constituents with wavelength- or temperature-
1422 dependent electrical properties. Here, we will present what we consider reasonable hypotheses
1423 about what the most important factors are.
1424 The unremarkable 2- to-13-cm wavelength dependence seen on Enceladus' leading side and also on
1425 Rhea is reminiscent of Europa's 3.5- to-13-cm wavelength invariance. The simplest interpretation is
1426 that at least the top few decimeters are uniformly clean and old enough for meteoroid bombardment
1427 to have created density heterogeneities that are able to produce coherent backscattering and extreme
1428 radar brightness.
1429 Enceladus' leading-side 2cm and 13cmcm TP albedos resemble Europa's 3.5cm and 13cm TP
1430 albedos, whereas Rhea's 2cm and 13cm TP albedos resemble Ganymede's 3.5cm and 13cmcm TP
1431 albedos, reflecting much greater contamination of Rhea's upper regolith.
1432 Dione's and Rhea's TP-2 albedos are similar, but Dione is dimmer at 13cm. Dione VIMS data show
1433 evidence of a bombardment of apparently dark particles on the trailing side (Clark et al., 2008a). If
1434 the dark material contains ammonia, and the ejecta systematics of meteoroid bombardment
1435 efficiently distribute the dark material over the surface of the object, multiple scattering will be cut
1436 off at a depth that may be too shallow for TP-13 to be enhanced.
1437 The striking wavelength dependence of Enceladus' trailing side (TP-2 >>TP-13) mirrors Europa's
1438 (TP-13 >>TP-70). It is easily understood if we assume that the top meter or so of ice was deposited
1439 so recently that there has not been time for the regolith to be matured by meteoroid bombardment.
1440
1441
That is, there are small-scale heterogeneities that produce coherent backscattering at 2cm, but none
of the larger-scale heterogeneities needed to produce coherent backscattering at 13cm.
1442 This idea is consistent with the fact that Enceladus' trailing side is optically brighter than its leading
1443 side, apparently because of the greater E-ring flux it receives (Buratti, 1988; Showalter et al.,
1444 1991b; Buratti et al., 1998). So all of Enceladus is very clean, but the uppermost layer of one or
1445 several meters on the trailing side is also extremely young, resulting in a striking hemispheric
1446 dichotomy at 13cm that is the most intense seen at that wavelength in the icy-satellite systems of
1447 either Jupiter or Saturn.
1448 At least some of the decrease in average SL-2 (Fig. 18) and the decrease in average wavelength
1449 dependence outward from Enceladus (Fig. 19) is probably due to the outward decrease in E-ring
1450 flux (Verbiscer et al., 2007).
1451 Finally, we get to Iapetus, where TP-2 is strongly correlated with the optical albedo: low for the
1452 optically dark leading-side material and high for the optically bright trailing-side material.
1453 However, Iapetus' TP-13 values show much less of a dichotomy and are several times lower than
1454 TP-2, being indistinguishable from the weighted mean of TP-13 for main-belt asteroids, 0.15 ± 0.10
1455
1456
(Tab. 5).
27

1457 Dark/Leading Bright/Trailing
1458 2.2-cm 0.32 + 0.09 0.58 + 0.18
1459 13-cm 0.13 + 0.04 0.17 + 0.04
1460
1461
Tab. 5 Total-power radar albedos
1462 Therefore, whereas Iapetus is anomalously bright at 2cm (similar to Callisto) and mimics the optical
1463 hemispheric dichotomy, at 13cm it looks like a typical main-belt asteroid, with minimal evidence of
1464 mimicking the optical hemispheric dichotomy.
1465 The Iapetus results are understandable if (1) the leading side's optically dark contaminant is present
1466 to depths of at least one to several decimeters, so that SL-2 can see the optical dichotomy, and (2)
1467 ammonia and/or some other radar-absorbing contaminant is globally much less abundant within the
1468 upper one to several decimeters than at greater depths, which would explain the wavelength
1469 dependence, that is, the virtual disappearance of the object's 2-cm signature at 13cm. The first
1470 conclusion is consistent with the dark material being a shallow phenomenon, as predicted from the
1471 thermal migration hypothesis discussed elsewhere in this chapter.
1472 Moreover, ammonia contamination may be uniform within the top few meters, so that the
1473 wavelength dependence arises because ammonia's electrical loss is much greater at 13cm
1474 (2.38GHz) than at 2cm (13.8GHz). Unfortunately, the likelihood of this possibility cannot yet be
1475 judged because measuring the complex dielectric constant of ammonia-contaminated water ice as a
1476 function of concentration, frequency and temperature is so difficult that, although needed, it has not
1477
1478
been done yet. This is a pity, because the lack of this information undermines the interpretation of
the 2cm to13cm wavelength dependence seen in other Saturnian satellites as well as of the possible
1479 influence of ammonia and its 'modulation' by magnetospheric bombardment (which varies with
1480 distance from Saturn) on the dependence of SL-2 on the distance from Saturn.
1481 Verbiscer et al. (2006) suggested that, while their spectral data do not contain an unambiguous
1482 detection of ammonia hydrate on Enceladus, their spectral models do not rule out the presence of a
1483 modest amount of it on both hemispheres. However, they note that the trailing hemisphere is also
1484 exposed to increased magnetospheric particle bombardment, which would preferentially destroy the
1485
1486
ammonia. In any case, ammonia does not appear to be a factor in Enceladus’ radar albedos.
1487 Iapetus Radar Image
1488 During the Iapetus 49 flyby, RADAR obtained an SAR image at 2km to12km surface resolution
1489 that covers much of the object's leading (optically darker) side. From a comparison with optical
1490 imaging (Fig. 20), we see that, as expected, the SAR reliably reveals the surface features seen in
1491
1492
ISS images.
1493 Fig. 20 RADAR SAR image of the leading side of Iapetus (Sept. 2007, left) and the ISS mosaic
1494 PIA 08406 (Jan. 2008), using similar projections and similar scales (the large crater (on the right in
1495
1496
the ISS image is about 550 km in diameter).
1497 The pattern of radar brightness contrast is basically similar to that seen optically, except for the
1498 large basin at 30°N, 75°W, which has more SL-2-bright than optically bright highlights. This
1499
1500
suggests that the subsurface of these highlights at decimeter scales may be cleaner than in other
parts of the optically dark terrain, which is consistent with the idea that the contamination by dark
1501 material might be shallow, at least in this basin. Conversely, many optically bright crater rims are
1502
1503
SL-2-dark, so the exposure of clean ice cannot be very deep.
1504 21.6 Geological Evolution
1505 The densities of the Saturnian satellites constrain their composition to be about 2/3 ice and 1/3 rock.
1506 Accretional energy, tidal despinning and radioactive decay might have heated and melted the
1507 moons. Due to the very low melting temperatures of various water clathrates, it is likely that the
1508 larger bodies differentiated early in their evolution while smaller satellites may still be composed
1509 and structured in their primitive state. For a more detailed discussion of the aspects of origin and
1510 thermal evolution, see Johnson et al. (2009) and Matson et al. (2009).
1511 The heavily cratered solid surfaces in the Saturnian system indicate the role of accretional and post-
1512 accretional bombardement in the geological evolution of the satellites. Observed crater densities are
1513 summarized in Dones et al. (2009). However, the source as well as the flux of bombarding bodies is
1514 still under discussion, and Dones et al. (2009) came to the conclusion that, since no radiometric
1515
sample data exist from surfaces in the outer solar system and the size-frequency distribution of
28

1516 comets is less well understood than that of asteroids, the chronology of the moons of the giant
1517 planets remains in a primitive state. Therefore, we refer the reader to Dones et al. (2009) for the
1518 crater chronology discussion. In interpreting the geological evolution of the Saturnian satellites, we
1519 will rely on relative stratigraphic relations. Nevertheless, absolute ages are given in Dones et al.
1520 (2009).
1521 In a stratigraphic sense, three major surface unit categories can be distinguished on all icy satellites:
1522 (1) Cratered plains, which are characterized by dense crater populations that are superimposed by
1523 all other geological features and, thus, are the relatively oldest units.
1524 (2) Tectonized zones of narrow bands and mostly sub-parallel fractures, faults, troughs and ridges
1525 of global extension indicating impact-induced or endogenic crustal stress. These units dissect
1526 cratered plains and partly dissect each other, indicating different relative ages.
1527 (3) Resurfaced units consisting of terrain that has undergone secondary exogenic and endogenic
1528 processes such as mass wasting, surface alterations (due to micrometeorite gardening, sputtering,
1529 thermal segregation and crater ray formation by recent impacts), contamination by dark material
1530 and cryovolcanic deposition mostly resulting in smoothed terrain. Superimposed on cratered plains
1531 and tectonized regions, these are the youngest surface units.
1532
1533
Heavily cratered surfaces are found on all the icy rocks and small satellites (Fig. 21) as well as on
1534 Fig. 21 Phoebe, 213 km in average diameter, was the first one of the nine classic moons of Saturn to
1535 be captured by the ISS cameras aboard Cassini in a close flyby about three weeks prior to Cassini’s
1536
1537
Saturn Orbit Insertion (SOI) in July 2004. The mosaic shows a densely cratered surface. Higherresolution
data revealed that the high frequency of craters extends with a steep slope down to craters
1538 only tens of meters in diameter (Porco et al., 2005). Image source: http://photojournal.jpl.nasa.gov,
1539
1540
image identification PIA06064.
1541 the medium-sized moons Mimas, Enceladus, Tethys, Dione, Rhea, Titan, Hyperion, Iapetus and
1542 Phoebe, indicating an intense post-accretional collision process relatively early in the geological
1543 history of the satellites. Cassini observations added to the number of large basins and ring structures
1544 mainly on Rhea, Dione and Iapetus. Differences in the topographical expression of these impact
1545 structures suggest post-impact processes such as relaxation, which may reflect differences in the
1546 thermal history and/or crustal thickness of individual satellites (e.g. Schenk and Moore, 2007; Giese
1547 et al., 2008). Although age models are based on different impactor populations, the cratered
1548 surfaces on Iapetus seem to be the oldest in the Saturnian system, whereas cratered plains on the
1549 other satellites seem to be younger (see Dones et al., 2009). One large (48 km in diameter)
1550
1551
stratigraphically young ray crater was found on Rhea (Wagner et al., 2007) (Fig.22).
1552 Fig. 22 A prominent bright ray crater with a diameter of 48 km is located approximately at 12.5° S,
1553 112° W on Saturn’s second-largest satellite Rhea. The crater is stratigraphically young, indicated by
1554 the low superimposed crater frequency on its floor and continuous ejecta (Wagner et al., 2008). The
1555 high-resolution mosaic (right) of NAC images, shown in context of a WAC frame, was obtained by
1556 Cassini ISS during orbit 049 on Aug. 30, 2007. The global view to the left showing the location of
1557 the detailed view is a single NAC image from orbit 056.
1558
1559 While the small satellites lack tectonic features, the medium-sized bodies experienced rotation-,
1560 tide- and impact-induced deformation and volume changes, which created extensional or
1561 contractional tectonic landforms. Parallel lineaments on Mimas seem to be related to the Herschel
1562 impact event (McKinnon, 1985; Schenk, 1989a), although freezing expansion cannot be ruled out as
1563
1564
the cause (Fig. 23).
1565 Fig. 23 The anti-Saturnian hemisphere of Mimas, centered approximately at lat. 22° S, long. 185°
1566 W, is shown in a single image (NAC frame N1501638509) at a resolution of 620 m/pixel, taken
1567 during a non-targeted encounter in August 2005. The densely cratered landscape implies a high
1568
1569
surface age. Grooves and lineaments infer weak tectonic on this geologically unevolved body.
1570 Enceladus exhibits the greatest geological diversity in a complex stress and strain history that is
1571 predominantly tectonic and cryovolcanic in origin, with both processes stratigraphically correlated
1572 in space and time, although the energy source is so far not completely understood (e.g. Smith et al.,
1573 1982; Kargel and Pozio, 1996; Porco et al., 2006; Nimmo and Pappalardo, 2006; Nimmo et al.,
1574
2007; Jaumann et al., 2008). Enceladus is discussed in detail in Spencer et al. (2009).
29

1575 Located in densely cratered terrain, the nearly globe-encircling Ithaca Chasma rift system on Tethys
1576 (Fig. 2) is 100 km wide, consists of at least two narrower branches towards the south and was
1577 caused either by expansion when the liquid-water interior froze (Smith et al., 1981) or by
1578 deformation by the large impact which formed Odysseus (Moore and Ahern, 1983). However,
1579 crater densities derived from Cassini images suggest Ithaca Chasma to be older than Odysseus
1580 (Giese et al., 2007), so that the rift system may be an endogenic feature.
1581 Dione exhibits a nearly global network of tectonic features such as troughs, scarps, ridges and
1582
1583
lineaments partly superimposed on each other (Fig. 24), which correlate with smooth plains
1584 Fig. 24 This image of Tethys, taken by the Cassini ISS NA camera in May 2008 in visible light,
1585 shows the old, densely cratered surface of Tethys and a part of the extensive graben system of
1586 Ithaca Chasma. The graben is up to 100 km wide and is divided into two branches in the region
1587 shown in this image. Using digital elevation models infers flexural uplift of the graben flanks and a
1588 total topography ranging several kilometers from the flanks to the graben floor (Giese et al., 2007).
1589 Image details: ISS frame number N1589080288, scale 1 km/pixel, centered at lat. 42° S, long. 50°
1590
1591
W in orthographic projection.
1592 (e.g. Plescia, 1983; Moore, 1984; Stephan et al., 2008). The global orientation of Dione’s tectonic
1593 pattern is non-random, suggesting a complex endogenic history of extension followed by
1594 compression. Cassini images indicate that some of the grabens are fresh, which implies that tectonic
1595
1596
activity may have persisted into geologically recent times (Wagner et al., 2006). Cassini also
revealed a crater-lineament relationship on Dione (Moore et al., 2004; Schenk and Moore, 2007),
1597 indicating exogenic-induced tectonic activity in the early geological history of the moon. Thus,
1598 Dione’s crust not only experienced stress caused by different exogenic and endogenic processes,
1599 such as ancient impacts, despinning and volume change over a long period of time, but also more
1600 recent tidal stress. As on Dione, wispy accurate albedo markings on Rhea are associated with an
1601 extensional fault system (Moore and Schenk et al., 2007; Wagner et al., 2007), although Rhea
1602
1603
seems less endogenically evolved than the other moons (Fig. 25).
1604 Fig. 25 Map of Dione’s trailing hemisphere showing an age sequence of troughs, graben and
1605 lineaments, inferred from various crosscutting relationships (Wagner et al., 2009). Three age groups
1606 of troughs and graben can be distinguished: Carthage and Clusium Fossae are oldest, cut by the
1607 younger systems of Eurotas and Palatine Chasmata, which in turn are truncated by Padua Chasmata
1608 which are the youngest. Several systems of lineaments with various trends, either parallel or radial
1609 (the latter previously assumed to be a bright ray crater named Cassandra) can be mapped which
1610
1611
appear to be younger than the troughs and graben.
1612 Parallel north-south trending lineaments are thought to be the result of extensional stress followed
1613 by compression (Moore et al., 1985) and appear in Cassini observations as poorly developed ridges.
1614 This suggests that they are older than the grabens identified in the wispy terrain that appear to have
1615 formed more recently (Wagner et al., 2006).
1616 Iapetus’ equatorial ridge (Porco et al., 2005a; Giese et al., 2008) (Fig. 7) is densely cratered, making
1617
1618
it comparable in age to any other terrains on this moon (Schmedemann et al., 2008). Models explain
the origin of the ridge either by endogenic shape and/or spin state changes (e.g. Porco et al., 2005a;
1619 Giese et al., 2008) or exogenic causes, with the ridge accumulating from ring material left over
1620 from the formation of proto-Iapetus (Ip, 2006). On Iapetus, diurnal tidal stress is rather low,
1621 indicating that the ridge is old, and its preservation suggests the formation of a thick crust early in
1622 the moon's history (Castillo-Rogez et al., 2007).
1623 Enceladus, where cryovolcanism is definitely active (Porco et al., 2006), is discussed in Spencer et
1624 al. (2009). Titan, where some evidence of cryovolcanism exists (Lopes et al., 2007a; Nelson et al.,
1625 2009a,b) is discussed in Jaumann et al. (2009) and Sotin et al. (2009).
1626 There is no direct evidence of cryovolcanic processes (Geissler, 2000) on the other Saturnian
1627 satellites. Although it has been suggested in pre-Cassini discussions that the plains on Dione
1628 (Plescia, and Boyce, 1982; Plescia, 1983; Moore, 1984) and Tethys (Moore and Ahern, 1984) are
1629 endogenic in origin, Cassini has not directly observed any cryovolcanic activity on these satellites.
1630 The plains correlate with tectonic features and exhibit distinct boundaries with more densely
1631 cratered units (Moore and Schenk, 2007), which indicate that the plains are a stratigraphically
1632 younger formation. In addition, there is evidence of material around Dione, Tethys and even Rhea
1633
(Burch et al., 2007; Jones et al., 2008; Clark et al., 2008a), but its origin is unclear. There is no
30

1634 evidence of cryovolcanism on Iapetus. Although the plains on Dione and Tethys are younger than
1635 the surrounding terrain, cryovolcanic activity must have occurred in ancient times, assuming that
1636 such processes caused these surface units.
1637 All medium-sized icy satellites appear to have a debris layer, as may be inferred from (1) the non-
1638 pristine morphology of most craters; (2) the presence of groups of large boulders and debris lobes
1639 seen in the highest-resolution ISS images; and (3) derived thermal-physical properties consistent
1640 with loose particulate materials. With the exception of Enceladus, which is addressed in Spencer et
1641 al. (2009), these debris layers were generated by impacts and their ejecta on the other satellites. The
1642 upper limiting case for a given satellite is crustal disruption from extreme global impact-induced
1643 seismic events. A way to evaluate the depth of fracture in such events is to assume that these events
1644 were energetic enough to open fractures in the brittle ice (brittle on the timescale of these events)
1645 composing the interiors of the target satellite, and that, during the impact event, the lithostatic
1646 pressure everywhere was less than the brittle tensile strength of cold ice. Brittle fracture caused by
1647 global seismicity, all else being equal, extends deeper into smaller satellites, given that larger
1648 satellites are sheltered from dynamic rupture by their interior pressure. These being transient events,
1649 a simple estimate of the realm where brittle fracture may be expected to be induced, and visibly
1650 sustained, can be obtained from global seismicity by applying an expression (e.g. Turcott and
1651
1652
Schubert, 2002) of the lithostatic pressure at radius Rp, and comparing this to the brittle tensile
strength of water ice at timescales in which extreme global impact-induced seismic events produce
1653 energy sufficient to break ice mechanically. The determination of the fracture depth is based on
1654
1655
laboratory-derived dynamic values of ~1MPa (Lange and Ahrens, 1987; Hawkes and Mellor, 1972).
-3
With ρ the (bulk) density of the body (approximated as 1000kg/m ), G the gravitational constant,
1656 and R the radius of the satellite, the depth of the fracture can be constrained. Based on this, the
1657
1658
value of R - Rp for Mimas is ~20 km (~10% of its radius), ~10 km (~2% of radius) for Tethys, and
~5 km (~0.5% of radius) for Rhea. This is a measure of the depth to which it is reasonable to expect
1659 the surface layer (i.e., the megaregolith) to be intensely fractured. Of course, a body may be
1660 fractured to greater depths by large impacts, and the resulting pore space may not easily be
1661 eliminated. Deep porosity is most easily maintained on smaller and presumably colder bodies such
1662 as Mimas (Eluszkiewicz, 1990). Thus, for Mimas we expect that relatively recent major impact
1663 events, such as Herschel, might exploit pre-existing fracture patterns in the 'megaregolith' which
1664 might penetrate deep into the body rather than focus energy, as has been suggested for the
1665 formation of the plains of Tethys (e.g. Bruesch and Asphaug, 2004; Moore et al., 2004).
1666 Consequently, the lack of antipodal focusing or axial symmetry about a radius through Herschel in
1667 the lineament pattern on Mimas may be unsurprising, if that pattern is indeed an expression of
1668 tectonic movements initiated by the Herschel impact.
1669 The impact-gardened regolith of most middle-sized icy satellites can be assumed to be many meters
1670 thick if the regolith in the lunar highlands is any guide. Calculations of impact-generated regolith,
1671 based on crater densities observed in Voyager images, yielded mean depths of 500m for Mimas,
1672 740m for Dione, 1600m for Tethys and 1900m for Rhea (Veverka et al., 1986).
1673 This regolith will be further distributed by gravity-driven mass wasting (non-linear creep, i.e.
1674 disturbance-driven sediment transport that increases nonlinearly with the slope due to granular
1675 creep (e.g. Roering et al., 2001; Forsberg-Taylor et al., 2004)) which reduces high-standing relief on
1676
1677
craters (e.g. lowering of crater rims and central peaks) and tectonic features (e.g. the brinks of
scarps) as well as filling depressions (e.g. Howard, 2007). The process of non-linear creep, when
1678 not competing with other processes, produces a landscape, which has a repetitive consistent
1679 appearance. Impact-induced shaking and perhaps diurnal and seasonal thermal expansion facilitate
1680 such mass wasting. Note that this common expression of mass wasting on regolith-covered slopes
1681 assumes no topography-maintaining precipitation of frosts on local topographical highs, as
1682 discussed below.
1683 Bedrock erosion might be caused by the sublimation of mechanically cohesive interstitial volatiles.
1684 The icy bedrock would crumble, move down slope through nonlinear creep and form mantling
1685 slopes, perhaps leaving only topographical highs (crests) composed of bedrock exposed. This
1686 situation would occur if the initial bedrock landscape disaggregates to form debris at a finite rate, so
1687 that exposure of bedrock gradually diminishes, lasting longest at local crests. Such a situation might
1688 explain the landscape of Hyperion. If there is re-precipitation of volatiles in the form of frosts on
1689 local topographical highs, these frosts would armor the topographical highs against further
1690 desegregation. It has been proposed that landscapes evolved in this fashion on Callisto (Moore et al,
1691 1999; Howard and Moore, 2008). Thermally driven global and local frost migration and
1692
redistribution may occur on Iapetus (Spencer et al., 2005; Giese et al., 2008), as previously
31

1693 discussed. This surficial frost redistribution may occur regardless of whether volatiles in the
1694
1695
bedrock of Iapetus sublimate and contribute to bedrock erosion.
1696 21.7 Conclusions
1697 The icy satellites of Saturn show great and unexpected diversity. In terms of size, they cover a range
1698 from ~1500 km (Rhea) to ~270 km (Hyperion) and even smaller 'icy rocks' of less than a kilometer
1699 in diameter. The icy satellites of Saturn offer an unrivalled natural laboratory for understanding the
1700 geology of diverse satellites and their interaction with a complex planetary system.
1701 Cassini has executed several targeted flybys over icy satellites, and more are planned for the future.
1702 In addition, there are numerous other opportunities for observation during non-targeted encounters,
1703 usually at greater distances. Cassini used these opportunities to perform multi-instrument
1704 observations of the satellites so as to obtain the physical, chemical and structural information
1705 needed to understand the geological processes that formed these bodies and governed their
1706 evolution.
1707 All icy satellites exhibit densely cratered surfaces. Although there have been some attempts to
1708 correlate craters and potential impactor populations, no consistent interpretation of the crater
1709 chronology in the Saturnian system has been developed so far (Dones et al., 2009). However, most
1710 of the surfaces are old in a stratigraphical sense. The number of large basins discovered by Cassini
1711 was greater than expected (e.g. Giese et al., 2008). Differences in their relaxation state provide
1712 some information about crustal thickness and internal heat flows (Schenk and Moore, 2007). The
1713
1714
absence of basin relaxation on Iapetus is consistent with a thick lithosphere (Giese et al., 2008),
whereas the relaxed Evander basin on Dione exhibits floor uplift to the level of the original surface
1715 (Schenk and Moore, 2007), suggesting higher heat flow.
1716 The wide variety in the extent and timing of tectonic activity on the icy satellites defies any simple
1717 explanation, and our understanding of Saturnian satellite tectonics and cryovolcanism still is at an
1718 early stage. The total extent of deformation varies wildly, from Iapetus and Mimas (barely
1719 deformed) to Enceladus (heavily deformed and currently active), and there is no straightforward
1720 relationship to predicted tidal stresses (e.g. Nimmo and Pappalardo, 2006; Matsuyama and Nimmo,
1721 2007, 2008; Schenk et al., 2008; Spencer et al., 2009). Extensional deformation is common and
1722 often forms globally coherent patterns, while compressional deformation appears rare (e.g. Nimmo
1723 and Pappalardo, 2006; Matsuyama and Nimmo, 2007, 2008; Schenk et al., 2008; Spencer et al.,
1724 2009). These global patterns are suggestive of global mechanisms, such as despinning or
1725 reorientation, which probably occurred early in the satellites’ histories (e.g. Thomas et al., 2007);
1726 present-day diurnal tidal stresses are unlikely to be important, except on Enceladus (e.g. Spencer et
1727 al., 2009) and (perhaps) Mimas and Dione (e.g. Moore and Schenk, 2007). Ocean freezing gives
1728 rise to isotropic extensional stresses; modulation of these stresses by pre-existing weaknesses (e.g.
1729 impact basins) may explain some of the observed long-wavelength tectonic patterns (e.g. Moore
1730 and Schenk, 2007). Except for Enceladus, there is as yet no irrefutable evidence of cryovolcanic
1731 activity in the Saturnian system, either from surface images and topography or from remote sensing
1732 of putative satellite atmospheres. The next few years will hopefully see a continuation of the flood
1733 of Cassini data, and will certainly see a concerted effort to characterize and map in detail the
1734 tectonic structures discussed here. Understanding the origins of these structures and the histories of
1735
1736
the satellites will require both geological and geophysical investigations and probably provide a
challenge for many years to come.
1737 Although the surfaces of the Saturnian satellites are predominately composed of water ice,
1738 reflectance spectra indicate coloring agents (dark material) on all surfaces (e.g. Fink and Larsen
1739 1975; Clark et al., 1984, 1986; Roush et al., 1995; Cruikshank et al., 1998a; Owen et al., 2001;
1740 Cruikshank et al., 2005; Clark et al., 2005, 2008a, 2009; Filacchione et al. 2007, 2008). Recent
1741 Cassini data provide a greater range in reflected solar radiation at greater precision, show new
1742 absorption features not previously seen in these bodies, and allow new insights into the nature of the
1743 icy-satellite surfaces (Buratti et al., 2005b; Clark et al., 2005; Brown et al., 2006; Jaumann et al.,
1744
1745
2006; Cruikshank et al., 2005, 2007, 2008; Clark et al., 2008a, b, 2009). Besides H2O ice, CO2
absorptions are found on all medium-sized satellites (Buratti et al., 2005b; Clark et al., 2005;
1746
1747
Cruikshank et al., 2007; Brown et al., 2006; Clark et al., 2008a), which make CO2 a common
molecule on icy surfaces as well. Various spectral features of the dark material match those seen on
1748 Phoebe, Iapetus, Hyperion, Dione and Epimetheus as well as in the F-ring and the Cassini Division,
1749 implying that the material has a common composition throughout the Saturnian system (Clark et al.,
1750 2008a, 2009). Dark material diminishes closer to Saturn, which might be due to surface
1751
contamination of the inner moons by E-ring particles and some chemical alteration processes (Clark
32

1752 et al., 2008a). Although the general composition of the dark material in the Saturnian system has not
1753
1754
yet been determined, it is known that it at least contains CO2, bound water, organics and, tentatively,
ammonia (Clark et al., 2008a). In addition, a blue scattering peak that dominates the dark spectra
1755 (Clark et al., 2008a, 2009) can be modelled by nano-phase hematite mixed with ice and dark
1756 particles. This led Clark et al. (2008b, 2009) to infer that sub-micron iron particles from meteorites,
1757 which oxidize (explaining spectral reddening) and become mixed with ice could best explain all the
1758 colors and spectral shapes observed in the system. Iron has also been detected in E-ring particles
1759 and in particles escaping the Saturnian system (Kempf et al., 2005). Previous explanations for the
1760 dark material were based on tholins (e.g. Cruikshank et al., 2005), but this is inconsistent with the
1761 lack of marked absorptions due to CH in the 3μm to 4μm region as well as other major absorptions at
1762 wavelengths longer than 2μm. Iron particles are consistent with previous Cassini studies that
1763 suggested an external origin for the dark material on Phoebe, Iapetus and Dione (Clark et al., 2005,
1764 2008a,b, 2009). Another remarkable issue relating to dark material is the albedo dichotomy on Iapetus,
1765 whose leading hemisphere is the darkest surface in the Saturnian system. In contrast to the outermost
1766 satellite, Phoebe, also covered by dark material, Iapetus' dark hemisphere is spectrally reddish, making
1767 it doubtful whether Phoebe is the source of the dark material. Clark et al. (2008a, 2009) explained the
1768 color discrepancy by red fine-grained dust (

1810
1811
1812
1813
1814
1815
1816
1817
1818
material has been observed throughout the system, and there is some evidence that this
contamination might also be a relatively recent process.
Acknowledgements: We gratefully acknowledge the long years of work by the entire Cassini team that allowed these
data of the Saturnian satellites to be obtained. We should also like to thank NASA, ESA, ASI, DLR, CNES and JPL for
providing support for the international Cassini team. We thank K. Stephan, R. Wagner and M. Langhans for data
processing. The discussions of reviewers C. Chapman and D. Domingue are highly appreciated. Part of this work was
performed at the DLR Institute of Planetary Research, with support provided by the Helmholz Alliance ’Planetary
Evolution and Life’ and the Jet Propulsion Laboratory under contract to the National Aeronautics and Space
1819 Administration.
We dedicate this work to Steve Ostro, who provided outstanding scientific contributions to the exploration of the
1820
1821
1822
Saturnian System.
1823 References
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2940
2941
Figures
2942 Chap.21_Icy Sat_Revised Version March 31, 2009
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2944
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Fig. 1
53

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Fig. 2
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Fig. 3
Fig. 4
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Fig. 8
Fig. 9a
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Fig. 9b
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Fig. 14 (left)
Fig. 14 (right)
62

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Fig. 15 (above)
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Fig. 15 (lower right)
Fig. 16 (left)
64

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Fig. 16 (right)
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Fig. 18
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Fig. 24
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69