icy satellites: geological evolution and surface processes

Chapter 20Icy Satellites: Geological Evolution and Surface Processes

Ralf Jaumann, Roger N. Clark, Francis Nimmo, Amanda R. Hendrix, Bonnie J. Buratti, Tilmann Denk, Jeffrey M. Moore,Paul M. Schenk, Steve J. Ostro�, and Ralf Srama

Abstract The sizes of the Saturnian icy satellites range from�1;500 km in diameter (Rhea) to �20 km (Calypso), andeven smaller ‘rocks’ of only a kilometer in diameter arecommon in the system. All these bodies exhibit remarkable,unique features and unexpected diversity. In this chapter, wewill mostly focus on the ‘medium-sized icy objects’ Mimas,Tethys, Dione, Rhea, Iapetus, Phoebe and Hyperion, andconsider small objects only where appropriate, whereas Titanand Enceladus will be described in separate chapters. Mimasand Tethys show impact craters caused by bodies that werealmost large enough to break them apart. Iapetus is unique inthe Saturnian system because of its extreme global brightnessdichotomy. Tectonic activity varies widely – from inactiveMimas through extensional terrains on Rhea and Dione tothe current cryovolcanic eruptions on Enceladus – and is notnecessarily correlated with predicted tidal stresses. Likelysources of stress include impacts, despinning, reorientation

R. JaumannDLR, Institute of Planetary Research, Rutherfordstrasse 2, 12489Berlin, GermanyandFreie Universität Berlin, Institute of Geological Sciences, Malteserstr.74-100, 12249 Berlin, Germany

T. DenkFreie Universität Berlin, Institute of Geological Sciences, Malteserstr.74–100, 12249 Berlin, Germany

R.N. ClarkU.S. Geological Survey, Denver Federal Center, Denver CO80225, USA

F. NimmoDepartment of Earth and Planetary Sciences, University of CaliforniaSanta Cruz, Santa Cruz, CA 95064, USA

A.R. Hendrix, B.J. Buratti, and S.J. Ostro�

Jet Propulsion Laboratory, California Institute of Technology,Pasadena, CA 91109, USA�Deceased

J.M. MooreNASA Ames Research Center, MS 245-3 Moffett Field, CA94035-1000, USA

P.M. SchenkLunar and Planetary Institute, Houston, TX 77058, USA

R. SramaMax Planck Institut für Kernphysik, 69117, Heidelberg, Germany

and volume changes. Accretion of dark material originatingfrom outside the Saturnian system may explain the surfacecontamination that prevails in the whole satellite system,while coating by Saturn’s E-ring particles brightens the innersatellites.

So far, among the surprising Cassini discoveries are thevolcanic activity on Enceladus, the sponge-like appearanceof Hyperion and the equatorial ridge on Iapetus – unique fea-tures in the solar system. The bright-ray system on Rhea wascaused by a relatively recent medium impact which formeda �40 km crater at 12ıS latitude, 112ıW longitude, whilethe wispy streaks on Dione and Rhea are of tectonic origin.Compositional mapping shows that the dark material on Ia-petus is composed of organics, CO2 mixed with H2O ice,and metallic iron, and also exhibits possible signatures ofammonia, bound water, H2 or OH-bearing minerals, and anumber of as-yet unidentified substances. The spatial pattern,Rayleigh scattering effect, and spectral properties argue thatthe dark material on Iapetus is only a thin coating on its sur-face. Radar data indicate that the thickness of the dark layerscan be no more than a few decimeters; this is also consistentwith the discovery of small bright-ray and bright-floor craterswithin the dark terrain. Moreover, several spectral features ofthe dark material match those seen on Phoebe, Iapetus, Hy-perion, Dione and Epimetheus as well as in the F-ring andthe Cassini Division, implying that throughout the Saturniansystem. All dark material appears to have a high content ofmetallic iron and a small content of nano-phase hematite.However, the complete composition of the dark material isstill unresolved, and additional laboratory work is required.As previously concluded for Phoebe, the dark material ap-pears to have originated external to the Saturnian system.

The icy satellites of Saturn offer an unrivalled natu-ral laboratory for understanding the geological diversity ofdifferent-sized icy satellites and their interactions within acomplex planetary system.

20.1 Introduction

Covering a wide range of diameters, the satellites of Sat-urn can be subdivided into three major classes: ‘icy rocks,’

M.K. Dougherty et al. (eds.), Saturn from Cassini-Huygens,DOI 10.1007/978-1-4020-9217-6_20, c� Springer Science+Business Media B.V. 2009

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with radii <10 km, ‘small satellites’ with radii <100 km,such as Prometheus, Pandora, Janus, Epimetheus, Telesto,Calypso, and Helene, and ‘medium-sized satellites’ withradii<800 km: Mimas, Enceladus, Tethys, Dione, Rhea, Hy-perion, Iapetus and Phoebe. The largest satellite, Titan, hasa radius of 2,575 km. We focus here on the ‘medium-sizedsatellites,’ considering the ‘small satellites’ where appropri-ate. Enceladus is discussed in a separate chapter (Spenceret al. 2009) and Titan in a dedicated book (Brown et al. 2009;Jaumann et al. 2009); therefore, these two satellites will notbe discussed in detail here. The main orbital and physicalcharacteristics of the medium and small satellites of Saturnare shown in Table 20.1.

Common properties of the Saturnian moons include a verylarge amount of water ice on the surface and in the inte-rior, resulting in low mean densities between �1:0 g=cm3

and �1:6 g=cm3, and numerous impact craters scattered overthe surface. Tectonic features and less heavily cratered plainsare restricted to only a few satellites and indicate endogenicgeological activity. The Voyager missions (e.g., Smith et al.1981, 1982) changed our perception of the medium-sizedSaturnian moons from small dots in large telescopes into realworlds with alien landscapes. Impact craters were recognised

as the dominant landforms, except for parts of Enceladus thatappeared to be crater-free. It was found that instead of crater-ing widespread tectonics seems to have shaped the surface ofEnceladus, and most craters showed evidence of viscous re-laxation. A causal connection with Saturn’s E-ring was alsosuspected because the E-ring is densest at Enceladus’ or-bit (Pang et al. 1984). However, similarly sized Mimas wasfound to look very different. Its trademark is a relativelylarge, 110 km crater, Herschel. Besides Herschel, there arenumerous other impact-related but few tectonic structures.On Tethys, a large valley (Ithaca Chasma), spanning about3/4 of Tethys’ globe, and a giant impact basin (Odysseus,�400 km in diameter) were found to dominate the surface.Dione and Rhea both showed wispy structures in early viewsof their trailing sides, but the resolution of the Voyager im-ages was insufficient to determine the origin of these sur-face features. Although it is the largest medium-sized moonof Saturn, Rhea shows numerous craters and few endo-genic geological structures. On Rhea’s trailing side, an ex-tended bright-ray pattern was detected in very low-resolutionVoyager images (Smith et al. 1981), possibly caused bya recent impact. On other outer solar system moons, onlyPwyll on Europa is known to have a similar structure, with

Table 20.1 Physical and orbital characteristics of the small (<100 km radius) and medium-sized moons (<800 km radius) of Saturn (from JPL athttp://ssd.jpl.nasa.gov/?satellites)Satellite a (103 m/ P (days) e i (deg) Mean radius (km) Mean density GM V0 or R P

Pan 134 0.575 0.0000 0.001 12.8 0.56 0:00033 ˙ 0:00005 19.4 0.5Atlas 138 0.602 0.0012 0.003 10. 0.5 0.00014 19.0 0.4Prometheus 139 0.613 0.0022 0.008 46:8 ˙ 5:6 0:435˙ 0:159 0:01246 ˙ 0:00083 15.8 0.6Pandora 142 0.629 0.0042 0.050 40:6 ˙ 4:5 0:530˙ 0:194 0:00995 ˙ 0:00155 16.4 0.5Janus 151 0.695 0.0068 0.163 90:4 ˙ 3:0 0:612˙ 0:062 0:1266 ˙ 0:0017 14.4 0.6Epimetheus 151 0.694 0.0098 0.351 58:3 ˙ 3:1 0:634˙ 0:102 0:0351 ˙ 0:0004 15.6 0.5Mimas 186 0.942 0.0202 0.001 198.2 1.150 2:530˙ 0:012 12.8 0.6Enceladus 238 1.370 0.0045 0.001 252.1 1.608 7:210 ˙ 0:011 11.8 1.0Tethys 295 1.888 0.0000 0.002 533 0.973 41:210 ˙ 0:007 10.2 0.8Telesto 295 1.888 0.0002 1.180 12:˙ 3:0 1.0 0.00048 18.5 1.0Calypso 295 1.888 0.0005 1.499 9:5˙ 1:5 1.0 0.00024 18.7 0.7Dione 377 2.737 0.0022 0.005 561.7 1.476 73:113 ˙ 0:003 10.4 0.6Helene 377 2.737 0.0071 0.213 16:˙ 4:0 1.5 0.0017 18.4 0.6Rhea 527 4.518 0.0010 0.029 764.3 1.233 154:07 ˙ 0:16 9.6 0.6Hyperion 1501 21.28 0.0274 0.461 133:0 ˙ 8:0 0:569˙ 0:108 0:37˙ 0:02 14.4 0.3Iapetus 3561 79.33 0.0283 14.968 735.6 1.083 120:50 ˙ 0:03 11 0.6Phoebe 12948 550.31 0.1635 26.891 106:6 ˙ 1:1 1:633˙ 0:049 0:5531 ˙ 0:0006 16.4 0:081 ˙ 0:002

Paaliaq 15200 686.95 0.3630 45.084 11.0 2.3 0.00055 21.3R 0.06Albiorix 16182 783.45 0.4770 34.208 16 2.3 0.0014 20.5R 0.06Siarnaq 17531 895.53 0.2960 46.002 20 2.3 0.0026 20.1R 0.06a – mean value of the semi-major axis.P – sidereal period P.e – mean eccentricity.i D mean inclination with respect to the reference plane ecliptic, ICRF, or local Laplace.V0 D mean opposition magnitude or R D red magnitude.p D geometric albedo (geometrical albedo is the ratio of a body’s brightness at zero phaseangle to the brightness of a perfectly diffusing disk of the same position and apparent sizeas the body). (After Thomas et al. 1983; Morrison et al. 1984; Thomas 1989; Showalter1991; Jacobson and French 2004; Jacobson et al. 2004; Spitale et al. 2006; Jacobson2006; Porco et al. 2005b; Jacobson et al. 2005; Jacobson 2007)

http://ssd.jpl.nasa.gov/?satellites

20 Icy Satellites: Geological Evolution and Surface Processes 639

rays extending over thousands of kilometers away from thecrater. During the Voyager era, the leading sides of Tethys,Dione and Rhea were found to be brighter than the trail-ing sides. Later, Mimas and Enceladus were found to haveslightly brighter trailing hemispheres; this general pattern ofalbedo dichotomies is consistent with alteration by Saturn’sE-ring (Hamilton and Burns 1994, Buratti et al. 1998). OnIapetus, the elliptical shape of the dark terrain on the lead-ing hemisphere was discovered along with giant mountainson the anti-Saturnian side. Hyperion was found by Voyagerto be very irregularly shaped and by ground-based tele-scopic observations to be without a regular rotational period(Klavetter 1989).

Most of our post-Voyager knowledge has been sum-marized by Gehrels and Matthews (1984) and Morrisonet al. (1986). However, important questions remained unan-swered before Cassini: the retrograde orbit and dark sur-face of Phoebe threw its Saturnian origin into doubt (e.g.,Cruikshank et al. 1983; Burns et al. 1979). Hyperion isdarker than the inner regular satellites (e.g., Cruikshank et al.1984), exhibits a surface dominated by cratering and spalla-tion (e.g., Thomas and Veverka 1985) and rotates chaotically(e.g., Wisdom et al. 1984), indicating a unique geologicalevolution. Almost all the small inner satellites move in re-markable coorbital, shepherd or Langrangian orbits, whichsuggests that they originated from the break-up of largerbodies (e.g., Morrison et al. 1986). The highly symmetricaldistribution of Iapetus’ dark material along the direction oforbital motion called for some external control, whereas thetopographical relationship of dark material on the floors ofbright-side craters indicated endogenic control, so that theorigin of the dark material was problematic (Morrison et al.1986). During the Voyager era the origin, nature and distribu-tion mechanism of dark material in the Saturnian system wasnot understood. The origin of bright features like the ‘wispystreaks’ on Rhea and Dione could not be determined by Voy-ager images due to their low resolution. Differences in craterdensity between some of the medium-sized satellites, suchas Dione, Rhea and Enceladus, as well as extended grabens,faults and valleys suggested endogenic activities. However,the nature of these activities was not well understood. For amore detailed discussion of our knowledge before Cassini,see Dougherty et al. (2009) and Orton et al. (2009). In termsof geology, the Cassini mission was designed to provide sig-nificantly better area coverage, spatial resolution, and spec-tral range, resulting in numerous new discoveries. A detaileddescription of the Cassini mission can be found in Doughertyet al. (2009) and Seal et al. (2009). Maps of the Saturniansatellites are included in Roatsch et al. (2009).

20.2 Cassini’s Exploration of Saturn’s IcySatellites

The Cassini spacecraft is equipped with instruments tailoredto investigate the surfaces, environments and interiors oficy satellites. The optical remote sensing (ORS) instrumentsuite includes cameras and spectrometers designed for highspatial and spectral resolution covering wavelengths between0:06 �m and 1; 000�m. With the imaging subsystem (ISS)(Porco et al. 2004), morphologic, stratigraphic and other ge-ological surface properties can be observed at spatial resolu-tions down to a few meters (locally) or a few hundred meters(globally), depending on flyby distances. Moreover, as manyas 33 color and polarization filter combinations permit map-ping geologically diverse terrain spatially at wavelengthsbetween �0:3 �m and �1:1 �m. The visible and infraredmapping spectrometer (VIMS; 0:4 �m to 5:1 �m) provideschemical and compositional spectral information with spatialresolutions from a few kilometers (globally) to better thanone hundred meters (locally) (Brown et al. 2004; Jaumannet al. 2006). The composite infrared spectrometer (CIRS) de-termines global and regional surface temperatures and ther-mal properties on a kilometer scale (Flasar et al. 2004). Theultraviolet imaging spectrograph (UVIS) provides informa-tion about thin atmospheres and volcanic plume structuresas well as about water ice and other minor constituents onthe surface (Esposito et al. 2004), operating in the 60 nm –190 nm wavelength range. A suite of magnetosphere andplasma science (MAPS) instruments characterizes the satel-lites’ environments by in-situ methods. Micron-sized dustgrains and neutral molecules released from the surface byactive or passive processes (volcanoes/sputtering) carry sur-face composition information to distances as far as hundredsof kilometers above the surface. Released material affects theplasma surrounding the satellite and is registered by varia-tions in the magnetic field, ion density, and neutral gas anddust density as well as gas and dusty composition.

The primary task of the radio science subsystem (RSS) isto determine the mass of the moons (Table 20.1) by track-ing deviations in Cassini’s trajectory. RADAR data can pro-vide unique information about the upper sub-surface (Elachiet al. 2004), though few close-up RADAR SAR observationswere performed during satellite flybys. The lack of a scanplatform for the remote sensing instruments prevents the ra-dio science, RADAR and remote sensing systems from op-erating simultaneously during a flyby because the antennaand remote sensing instruments are oriented 90ı apart on thespacecraft. Even joint surface scans by the ORS instrumentsrequire significant compromises between individual observa-tions due to data rate and integration time constraints. While


the ISS instrument needs short ‘dwell’ times for mosaicking,the VIMS depend on long ‘dwells’ for improved signal-to-noise ratios. On the other hand, the CIRS and UVIS instru-ments contain line-scanning devices, which depend on eitherslow or fast slews to scan the surfaces. Thus, reaching sat-isfactory compromises is a major challenge in the planningprocess.

Since the Cassini spacecraft orbits the planet rather thanindividual satellites, the moons can only be observed at var-ious distances and illumination conditions, ideally duringvery close-targeted flybys. Depending on the distance of amoon to Saturn, the flybys occur at very different veloci-ties. In March 2008, for instance, Enceladus was passed at�14 km=s, while the Iapetus flyby in September 2007 tookplace at a leisurely 2.4 km/s. The flyby geometry can alsovary significantly. The closest-ever approach was during theEnceladus flyby on October 9, 2008, when the spacecraftskimmed as low as 25 km over Enceladus’ surface. Other,more typical targeted flyby altitudes occur between 100 kmand 2,000 km. A flyby can be polar, equatorial, or in between,and at the closest approach, the sub-spacecraft point over amoon can be located either over its illuminated or over itsunlit side.

The MAPS instruments measure densities or field gradi-ents over time. Sputtering processes lead to high dust andneutral-gas densities above the surface, so that environmentalin-situ instruments rely on measurements as close to the sur-face as possible. Further plasma density variations caused bysatellites moving in the magnetosphere are empty flux tubes(wakes) and the drop of plasma along the related L-shell.Crossing those regions is of high value for in-situ plasmainvestigations that provide indirect information about thesatellite surface and the density of its neutral and plasmaenvironment.

During the nominal mission, Cassini performed nine tar-geted as well as numerous close flybys of icy satellites(Table 20.2).

Three of the targeted flybys were dedicated to Enceladus,two in 2005 and one in 2008. In addition, a very close butnon-targeted Enceladus flyby happened in February 2005.As the data provided by the MAPS instruments in March2005 had raised suspicions about potential internal activity(Spahn et al. 2006; Waite et al. 2006), the July 2005 flybywas lowered from 1,000 km to 173 km to allow for increasedsensitivity in MAPS measurements. This flyby supplied over-whelming evidence from multiple instruments of current ac-tivity on Enceladus, eventually confirmed by direct images ofplumes taken during a distant flyby in November 2005. Thelocations of the plumes were found to coincide with lineartectonic features near the south pole, dubbed ‘tiger stripes’.The CIRS instrument also identified these ‘tiger stripes’ asunusually warm linear features (Spencer et al. 2006); the

tiger stripes contain relatively large ice particles identifiedby VIMS (Brown et al. 2006; Jaumann et al. 2008).

All other moons had only one targeted flyby during thenominal mission, if any. The Tethys, Hyperion, and Dioneflybys in the late summer/early fall of 2005 were performedwithin 3 weeks of each other during the same orbit of Cassiniaround Saturn. Due to Hyperion’s chaotic rotation, it was im-possible to predict which part of its surface would be vis-ible and how the surface might look. By coincidence, theviewing perspective during the approach was almost iden-tical with the best Voyager 1 views. It showed a giant crater,dark material within numerous smaller craters and a verylow mean density of only 0:54 g=cm3. The ‘wispy streaks’ ofDione were observed at very oblique viewing angles duringthe October 2005 flyby, and were confirmed to be tectonic inorigin (Wagner et al. 2006). The images from this targetedDione flyby are among the most spectacular pictures of themission. The Rhea flyby in November 2005 was dedicated toradio science, and the results from that flyby are reported inMatson et al. (2009). However, while the spacecraft itself wasinertially pointed with its antenna towards Earth, the camerasslewed across the surface, obtaining the highest-resolved im-ages of Rhea to date.

In 2006, no targeted flybys took place, primarily becausethe spacecraft spent most of its time in highly inclined orbits.The very close but non-targeted Rhea flyby in August 2007was dedicated to remote sensing and revealed small detailsof the prominent bright-ray crater and its ejecta environment.This might be the youngest large crater in the Saturnian sys-tem. The targeted Iapetus flyby in September 2007 over theanti-Saturnian and trailing hemispheres yielded very high-resolution observations of the highest parts of the equatorialridge, the transition zone between the dark and the bright ter-rain, and the bright trailing side. Combined with data from amore distant, non-targeted flyby on December 31, 2004, andother distant observations, these data permitted developing amodel of the formation of the global and unique brightnessdichotomy of this unusual moon (Denk et al. 2009; Denk andSpencer 2009). For more detailed descriptions of the Cassinimission, see Dougherty et al. (2009), Orton et al. (2009) andSeal et al. (2009).

20.3 Morphology, Geology and Topography

Voyager’s partial global mapping at 1 km resolution sug-gested that the morphology of the Saturnian satellites wasdominated by impact craters (Smith et al. 1981). Indeed,Voyager’s best images, 500 m resolution views of Rhea,showed a landscape saturated with craters. Voyager alsofound evidence of limited volcanic and tectonic activity in


Table 20.2 Missioncharacteristics of the Cassini icysatellites flybys; C/A: : : closestapproach Target Orbit Flyby date Altitude (km)

Solar PhaseAngle 2 h before,at, and 2 h afterC/A (deg)

Target (T)/non-targeted(nT)

Phoebe 0 June 11, 2004 1;997 85/25/90 TDione B December 15, 2004 72;100 108/85/62 nTIapetus B/C December 31, 2004 123;399 87/94/100 nTEnceladus 3 February 17, 2005 1;176 25/114/158 nTEnceladus 4 March 9, 2005 500 47/42/130 THyperion 9 June 10, 2005 168;000 25/35/47 nTEnceladus 11 July 14, 2005 170 47/63/132 TMimas 12 August 2, 2005 62;700 41/57/91 nTTethys 15 September 24, 2005 1;503 21/68/155 THyperion 15 September 26, 2005 505 51/50/127 TDione 16 October 11, 2005 500 23/66/156 TIapetus 17 November 12, 2005 415;400 89/91/93 nTRhea 18 November 26, 2005 1;264 19/87/161 TRhea 22 March 21, 2006 82;200 113/137/157 nTIapetus 23 April 11, 2006 602;400 123/124/125 nTEnceladus 27 September 9, 2006 39;900 105/118/105 nTDione 33 November 21, 2006 72;300 125/144/119 nTHyperion 39 February 16, 2007 175;000 66/64/64 nTTethys 47 June 27, 2007 18;400 148/116/49 nTTethys 49 August 29, 2007 55;500 135/102/72 nTRhea 49 August 30, 2007 5;800 125/46/41 nTIapetus 49 September 10, 2007 1;620 139/59/27 TDione 50 September 30, 2007 43;400 89/49/19 nTHyperion 51 October 21, 2007 122;000 115/116/115 nTMimas 53 December 3, 2007 84;200 161/139/82 nTEnceladus 61 March 12, 2008 52 115/135/65 TEnceladus 80 August 11, 2008 50 108/110/73 TEnceladus 88 October 9, 2008 25 108/112/73 TEnceladus 91 October 31, 2008 200 108/113/73 TTethys 93 November 16, 2008 57;100 79/41/58 nTTethys 94 November 24, 2008 24;200 123/159/84 nTTethys 115 July 26, 2009 68;400 105/89/79 nTRhea 119 October 13, 2009 40;400 140/82/25 nTMimas 119 October 14, 2009 44;200 152/102/44 nTHyperion 119 October 16, 2009 127;000 155/140/126 nTEnceladus 120 November 2, 2009 100 177/90/5 TEnceladus 121 November 21, 2009 1;600 146/87/36 TRhea 121 November 21, 2009 24;400 127/58/10 nTTethys 123 December 26, 2009 52;900 136/77/17 nTDione 125 January 27, 2010 45;100 158/106/53 nTMimas 126 February 13, 2010 9;500 173/99/21 nTRhea 127 March 2, 2010 100 176/87/3 TDione 129 April 7, 2010 500 167/79/11 TEnceladus 130 April 28, 2010 100 171/93/9 TEnceladus 131 May 18, 2010 200 153/108/29 TTethys 132 June 3, 2010 52;600 154/99/45 nT

the form of relatively smooth plains and linear features(Smith et al. 1981, 1982; Plescia and Boyce 1982, 1983;Moore et al. 1985; Moore and Ahern 1983), especially onEnceladus. The Enceladus discoveries aside, Cassini’s con-siderably improved resolution and global mapping coverageconfirmed the abundance of impact craters on all satellites

(Dones et al. 2009) and also revealed surprisingly varieddegrees of past endogenic activity on Tethys, Rhea and es-pecially Dione, indicating that these satellites were moredynamic than originally thought. Observed deformationsrange from linear grooves (Mimas) and nearly global-scalegrabens (Tethys, Rhea) to repeated episodes of tectonics and


resurfacing (Dione, Enceladus). To date, little or no directevidence of past or present endogenic activity has been iden-tified on Hyperion, Phoebe or Iapetus (other than its poten-tially endogenic equatorial ridge). This diversity is in itself apuzzle that has not yet been explained. In the following sec-tions, we examine the features observed, their morphology,and implications for internal evolution.

20.3.1 Craters

Comparison with craters on larger icy satellites (e.g.,Ganymede) is illuminating because surface gravities differby a factor of 5–10. Central-peak craters are common onSaturnian satellites. Central peaks in larger craters can be ashigh as 10 km, many of them rising above the local mean el-evation. Cassini has revealed little evidence of terrace forma-tion, and floor uplift dominates clearly over rim collapse inthe formation of complex craters (Schenk and Moore 2007).Pancake (formerly called pedestal) ejecta have been observedin smooth plains on Dione. This kind of ejecta deposit, onceonly known as a feature of the Galilean satellites, now ap-pears to be common on icy satellites in general. The fact thatit is not obvious on other satellites is most likely due to thedifficulty of identifying it on these rugged, heavily crateredsurfaces.

A transition to larger, more complex crater landforms(e.g., central pit craters, multi-ring basins) on these satellitesis predicted to occur at much larger diameters .>150 km/than those seen on Ganymede and Callisto (McKinnon andMelosh 1980; Schenk 1993), due to their weaker gravity.True central-pit craters like those on Ganymede or Cal-listo (Schenk 1993) have not been observed, although thecentral complex of Odysseus features a central depressionand closely resembles its counterparts on large icy moons.The dominant morphology at larger diameters .>150 km/is peak-ring or similar, with a prominent central peak sur-rounded by a ring or an elevated circular plateau. The num-ber of such basins discovered by Cassini on Iapetus and Rheawas greater than expected. Indeed, at least 10 basins largerthan 300 km across have been seen on Iapetus (e.g., Gieseet al. 2008).

The most ancient large basins on Rhea and Tethys havebeen modified by relaxation and long-term impact bombard-ment (e.g., Schenk and Moore 2007). The younger basins aregenerally deep and unrelaxed. However, the floor of Evan-der, is the largest basin identified on Dione, has been re-laxed and uplifted to the level of the original surface, and al-though its rim and peak remain prominent (e.g., Schenk andMoore 2007). Relaxation of this basin is consistent with ele-vated heat flows on Dione, as indicated by the relaxation ofsmaller (diameters of 25–50 km) craters within the extensive

smooth plains to the north of Evander. Relaxation in thisbasin, which has a low superposed crater density, is consis-tent with elevated heat flows, as indicated by the smaller re-laxed craters nearby and the extensive smooth plains to thenorth. The relaxation state of impact basins can be deter-mined by comparing their depth to that of unmodified cratersand examining their shape. The most ancient large basins onRhea (e.g., Tirawa) and Tethys are roughly 10–20% shal-lower than predicted and feature prominent domical or coni-cal central uplifts that approach or exceed the local elevationof the satellite (Schenk and Moore 2007), which suggeststhey have been modified to some degree by relaxation. Theyoungest basins appear to be deeper and possibly unrelaxed.The most prominent examples include Herschel on Mimas(diameter �120 km, depth �11 km) and Odysseus (diameter�450 km, depth >8 km). In neither case do the central up-lifts rise above the local topography. A prominent exceptionis the relatively young 400 km wide Evander basin on Dione,which may be either �3Gyr (Wagner et al. 2006) or less than2 Gyr old (Kirchoff and Schenk 2008; 2009). For a more de-tailed discussion of surface ages, see Dones et al. (2009).

Iapetus’ leading side shows more relief than the othermedium-sized Saturnian satellites (Thomas et al. 2007b;Giese et al. 2008). The lack of basin relaxation on Iapetusis consistent with the presence of a thick (50–100 km) litho-sphere in Iapetus’ early history. Such a lithosphere might alsosupport Iapetus’ prominent equatorial ridge. The ridge showsa complex morphology, probably the result of impact erosionin some places and of post-formation tectonic modificationsin others.

The abundance of large impact basins, even in Voyager’slimited view, suggested that such basins could have disrup-tive effects on these satellites (Smith et al. 1981). Severalattempts to discover evidence of these effects (e.g., seismicshaking, global fracturing) have met with limited success(e.g., Moore et al. 2004), but Cassini-based studies are nowin progress.

20.3.2 Tectonics

The Voyager 1 and 2 spacecraft revealed that the medium-sized icy Saturnian satellites have undergone varying degreesof tectonic deformation (Smith et al. 1981; 1982; Plesciaand Boyce 1982, 1983). Although such deformation can pro-vide constraints on the interior structure and evolution of asatellite, the imaging resolution was often insufficient to per-mit sustainable conclusions. Post-Voyager advances includea vastly better understanding of tectonic deformation andthermal evolution in the Jovian system thanks to Galileo, andgreatly improved experimental measurements of relevant iceproperties (Schmitt et al. 1998). In this section, we combine


this improved understanding with the results of the ongoingCassini mission to discuss the current state of knowledge oficy-satellite tectonics in the Saturnian system. We also exam-ine the extent to which cryovolcanic processes may be rele-vant. Most of the topics covered in this section are treated ingreater detail in Collins et al. (2009).

Tectonics on icy satellites comprises three major aspects:(1) the various mechanisms by which ice deforms when sub-jected to tectonic stress; (2) the mechanisms by which tec-tonic stress may arise; and (3) the observational constraintson these deformation mechanisms.

From a tectonics point of view, ice is rather similar to sil-icate materials: under surface conditions its behavior is rigidor brittle, while at greater depths it is ductile. The rheologyof ice is complicated and cannot be discussed in depth here;useful discourses on its elastic, brittle and viscous propertiesmay be found in Petrenko and Whitworth (1999), Beemanet al. (1988) and Durham and Stern (2001), respectively. Al-though a real ice shell will exhibit all three kinds of behav-ior, it can be modeled as a simple elastic layer (e.g., McNutt1984). This effective elastic thickness depends on the temper-ature gradient within the ice shell during deformation and canbe inferred from topographical measurements (e.g., Gieseet al. 2007). Thus, measuring the elastic thickness places auseful constraint on satellite thermal evolution (e.g., Nimmoet al. 2002). Pressures within Saturnian satellites are gener-ally low (see Table 20.3).

Thus, although ice transmutes into high-pressure forms atpressures greater than about 0.2 GPa (e.g., Sotin et al. 1998),this effect is relevant mainly on Titan.

Tectonic features may be used to infer the direction and(sometimes) the magnitude of the driving stresses. Observa-tions can thus place constraints on the mechanisms whichgave rise to these tectonic features. We briefly discuss sev-eral mechanisms which are potential candidates for produc-ing tectonic features.

Satellites experience both rotational and tidal deforma-tion. If the rotation rate or the magnitude of the tide changes,

the shape of the satellite will change as well, generatingglobal-scale patterns of stress. A recent compilation of thepresent-day shapes of the Saturnian satellites is given byThomas et al. (2007b); a good general description of tidaland rotational satellite deformation is given in Murray andDermott (1999).

If a synchronously rotating satellite’s orbital eccentricityis zero, the tidal bulge will be at a fixed geographic point, andof constant amplitude. The maximum amplitude H of thisstatic (or permanent) tidal bulge depends on the mass andradius of the satellite, the mass of the primary, the distancebetween them, and the rigidity of the satellite. The latter isdescribed by the h2 Love number, where h2 has a value of 5/2for a fluid body of uniform density which, however, declinesas the rigidity (or shear modulus) � of the satellite increases(e.g., Murray and Dermott 1999).

In a synchronous satellite with non-zero orbital eccen-tricity, the amplitude and direction of the static tidal bulgewill change slightly, resulting in a time-varying diurnal tide.Causing diurnal stresses, this tide has an amplitude ex-pressed by 3eH, where e is the eccentricity and H is theequilibrium tidal bulge (e.g., Greenberg and Geissler 2002).Obliquity librations are potentially an additional source oftidal stress (Bills and Ray 2000); they can be large if theirperiod is a small multiple of the orbital period (Wisdom2004).

Table 20.3 shows the diurnal tidal stresses of the icy Sat-urnian satellites at zero rigidity, demonstrating the generaldecrease in stress as the semi-major axis lengthens. For com-parison, the diurnal stresses on Europa thought to be respon-sible for cycloidal features are a few tens of kPa (Hoppaet al. 1999a). Because the principal stresses rotate duringeach orbit, contraction, extension and mixed-mode horizon-tal principal stresses which promote strike-slip motions areall possible as a satellite orbits (Hoppa et al. 1999b). Thismotion can lead to both shear heating (Nimmo and Gaidos2002) and the monotonic accumulation of strike-slip offsets(Hoppa et al. 1999b).

Table 20.3 Data from Yoder (1995) and Thomas et al. (2007b).R; Ms , a and e are the radius, mass, semi-major axis, period and eccen-tricity, respectively. H is the permanent tidal bulge assuming h2 D 2:5.3 eH is the approximate magnitude of the diurnal tidal bulge – note thatthis is likely an overestimate for small satellites because their rigiditywill reduce h2. “Tidal” is the tidal heat production assuming a homoge-neous body with k2 D 1:5 (again, a likely overestimate for small satel-

lites) and Q D 100. “Rad” is the radiogenic heat production assuminga chondritic rate of 3:5�10�12 W=kg. “Stress” is the approximate diur-nal tidal stress given byEeH=RwhereE is Young’s modulus (assumed9 GPa). Pcen is the central pressure, assuming uniform density. “MG”is the “metamorphic grade”, modified after Johnson (1998), which is aqualitative assessment of the degree of deformation: I D unmodified,II D intermediate, III D heavily modified

R

(km)Ms

.1020 kg/A

.103 km/ e H (m)3 eH(m)

Tidal1010 .W/

Rad..1010 W/

Stress(kPa)

Pcen(MPa) MG

Mimas 198 0.37 186 0:0202 9;180 556 78 0:013 8;400 7:1 IEnceladus 252 1.08 238 0:0045 3;940 53 2 0:038 630 23 IIITethys 533 6.15 295 0: 7;270 0 0 0:22 0 37 IIDione 562 11 377 0:0022 2;410 16 0.86 0:39 85 97 IIRhea 764 23 527 0:001 1;440 4:3 0.067 0:81 17 124 IITitan 2;575 1,346 1;222 0:0292 254 22 45 47 26 3;280 -Iapetus 736 18 3;561 0:0283 5 0:4 2:6� 10�5 0:63 1:7 88 I/II


A satellite surface may move longitudinally with re-spect to the static tidal bulge due to tidal (Greenberg andWeidenschilling 1984) or atmospheric torques (Lorenz et al.2008) or, in the case of floating ice shells, if the thermallyequilibrated ice shell thickness variations are not compat-ible with rotational equilibrium (Ojakangas and Stevenson1989a). This non-synchronous rotation leads to surfacestresses (Helfenstein and Parmentier 1985). Stresses fromnon-synchronous rotation increase with the angular speed ofrotation but will dissipate with time owing to viscous re-laxation (Greenberg and Weidenschilling 1984; Wahr et al.2009). Maximum tensile stress depends on the amount of re-orientation and the degree of flattening, f , which depends inturn on the satellite rotation rate (Leith and McKinnon 1996).In non-rigid satellites, the stresses due to non-synchronousrotation will exceed the diurnal tidal stresses for a reorien-tation typically in excess of roughly one degree. In certaincircumstances, the surface of a satellite may reorient rel-ative to its axis of rotation (Willemann 1984; Matsuyamaand Nimmo 2008) and give rise to global fracture pat-terns (e.g., Schenk et al. 2008). This process of true po-lar wander is conceptually very similar to non-synchronousrotation (Melosh 1980; Leith and McKinnon 1996), and ‘re-orientation’ will in the following be understood as imply-ing either non-synchronous rotation or true polar wander.In general, rotation axis motion occurs roughly perpendic-ular to the tidal axis because this path is energetically fa-vored (Matsuyama and Nimmo 2007). True polar wander canoccur because of ice shell thickness variations (Ojakangasand Stevenson 1989b), long-wavelength density anomalies(Janes and Melosh 1988; Nimmo and Pappalardo 2006),volatile redistribution (Rubincam 2003) or impacts, either di-rectly (Chapman and McKinnon 1986) or due to the creationof a new impact basin (Melosh 1975; Murchie and Head1986; Nimmo and Matsuyama 2007). Like non-synchronousrotation, the maximum stress developed by true polar wan-der depends on the amount of reorientation and the effectiveflattening f , with the value of f depending on whether anyreorientation of the tidal axis has occurred (Matsuyama andNimmo 2008).

Satellites that are initially rotating at a rate faster than syn-chronous will despin (i.e., reduce their rotation rate) to syn-chronous rotation in timescales that are generally very shortcompared to the age of the solar system (e.g., Murray andDermott 1999). Spin rate reductions result in shape changesand stresses – compressive at the equator as the equatorialbulge collapses, and tensile at the poles as the poles elon-gate. The maximum differential stress caused by despinningin a hydrostatic body depends mainly on the satellite radiusand the difference between the initial and the final angular ro-tation velocities (Melosh 1977). Spin-up may occur in somecircumstances (e.g., if the satellite undergoes differentiation),and in this case the signs of all the stresses are reversed.

Stresses along lines of longitude are always larger than thestresses along lines of latitude. Thus, irrespective of whethera satellite is spinning down or spinning up, equatorial tec-tonic features are expected to be oriented in a north-south di-rection. This result is important when considering the originof the global equatorial ridge on Iapetus.

Synchronous satellites, which are evolving outwards dueto tidal torques from the primary, will undergo both areduction in their spin rate and a parallel reduction in the tidalbulge amplitude. This combination of despinning and tidalbulge reduction causesastresspattern in whicha regionaroundthe sub-planet point experiences compressive stress, middlelatitudes experience horizontal shear stress, and the polesundergo tension (Melosh 1980; Helfenstein and Parmentier1983). The maximum principal stress difference is similar tothat produced by despinning alone (Melosh 1980).

A large variety of mechanisms can lead to volume changeswithin a satellite, and thus to extensional or contractionalfeatures on the surface (Squyres and Croft 1986; Kirk andStevenson 1987; Mueller and McKinnon 1988). Volumechanges generate isotropic stress fields on the surface. Apotentially important source of expansion or contraction isthe large density contrast between ice I and water. As waterfreezes to ice I, large surface extensional stresses may re-sult (Cassen et al. 1979; Smith et al. 1981; Nimmo 2004a).The reason for this effect is that, on a spherical body, the in-creased volume of ice relative to water drives the ice shelloutwards; and this increase in radius results in extension. Afractional change in radius of 0.1% gives rise to stresses ofabout 10 MPa. This effect is much reduced if there are high-pressure ice phases freezing simultaneously (Squyres 1980;Showman et al. 1997).

Similar but smaller effects are caused by thermal expan-sion or contraction (Ellsworth and Schubert 1983; Hillierand Squyres 1991; Showman et al. 1997; Nimmo 2004a). Aglobal temperature change of 100 K will cause stresses onthe order of 100 MPa for a thermal expansivity of 10�4=K.Warming may also lead to silicate dehydration (e.g., Squyresand Croft 1986), which in turn causes expansion.

Satellites are quite likely to have non-axisymmetric struc-tures, in which case some of the above-mentioned globalmechanisms may lead to local deformation. For instance, lo-cal zones of weakness or pre-existing structures can resultin enhanced tidal dissipation and deformation (e.g., Sotinet al. 2002; Nimmo 2004b) and/or the alteration of localstress trajectories. Convection (thermal or compositional)and buoyancy forces due to lateral shell thickness variationscan generate local deformation. They are discussed at greaterlength in Collins et al. (2009). We will not discuss these pro-cesses further here, as they do not tend to produce globallyorganized tectonic structures.

Large impact events can also potentially create globalfracture networks. This mechanism is implicit in the model


presented by E. Shoemaker in Smith et al. (1981) for thecatastrophic breakup and re-accretion of the medium-sizedicy satellites of Saturn. Presumably, impacts not quite largeenough to completely disrupt a satellite could do consid-erable damage in the form of global fracturing, and post-Voyager studies of these satellites focused in part on attemptsto find evidence of an incipient break-up of this kind. For in-stance, the surface of Mimas is crossed by a number of linear-to-arcuate, sub-parallel troughs (Fig. 20.1) that are thoughtto be the consequence of global-scale fractures during theHerschel impact event (McKinnon 1985; Schenk 1989a).

Voyager and Cassini mapping indicates that these fea-tures, if related to Herschel, do not form a simple radial orconcentric pattern. Nor has it been demonstrated that anyof these features form on a plane intersecting Herschel, orwhether they should do so. It thus remains possible that theseare random fractures formed during freeze expansion of theMimas interior. Detailed mapping remains to be done to testthese hypotheses. Mimas does not otherwise exhibit endo-genic landforms and will, therefore, not be discussed further.

Enceladus shows the greatest tectonic deformation of anyof the Saturnian satellites, and very significant spatial vari-ations in surface age (Smith et al. 1982; Kargel and Pozio1996; Porco et al. 2006; Jaumann et al. 2008). Spencer et al.(2009) discusses the tectonic behavior of Enceladus in muchgreater detail; only a very brief summary is given here.

Fig. 20.1 Linear troughs on Mimas, extending east to west. Thesetroughs are interpreted as fractures, possibly related to the Herschelcrater. Orthographic map projection at 400 m/pixel is centered on theantipode to Herschel, the largest impact basin on Mimas

Three different terrain types exist on Enceladus. Ancientcratered terrains cover a broad band from 0 to 180ı longitude.Centred on 90ı and 270ı longitude there are younger, de-formed terrains roughly 90ı wide at the equator. The south-ern polar region south of 55 ıS is heavily deformed and al-most uncratered. The most prominent tectonic features ofthis region are the linear depressions called ‘tiger stripes’(Porco et al. 2006). These tiger stripes are typically �500mdeep, �2 km wide, up to �130 km long and spaced �35 kmapart. Crosscutting fractures and ridges with almost no su-perimposed impact craters characterize the area between thestripes. The tiger stripes are apparently the source of the gey-sers observed to emanate from the south polar region (Spitaleand Porco 2007). The energy source of these geysers is un-known, but it may be shear heating generated by strike-slipmotion at the tiger stripes (Nimmo et al. 2007). The coher-ent (but latitudinally asymmetrical) global pattern of defor-mation on Enceladus is currently unexplained, but it is mostprobably due to shape changes, perhaps related to a hypothet-ical episode of true polar wander (Nimmo and Pappalardo2006).

Nearly encircling the globe, Ithaca Chasma is the mostprominent feature on Tethys (Fig. 20.2) (Chen and Nimmo2008).

It extends approximately 270ı around Tethys and is not acomplete circle. It is narrowly confined to a zone which liesalong a large circle whose pole is only �20ı from the centerof Odysseus, the relatively fresh largest basin on the satellitewhich is 450 km wide (Smith et al. 1982; Moore and Ahern1983). Smith et al. (1981) suggested that Ithaca Chasma wasformed by freeze-expansion of Tethys’s interior. They notedthat if Tethys was once a sphere of liquid water covered by athin solid ice crust, freezing in the interior would have pro-duced an expansion of the surface comparable to the area ofthe chasm (�10% of the total satellite surface area). This hy-pothesis fails to explain why the chasm occurs only withina narrow zone; besides, it is difficult to account for the geo-physical energy needed for the presence of a molten interior.Expansion of the satellite’s interior should have caused frac-turing over the entire surface in order to effectively relievestresses in a rigid crust (Moore and Ahern 1983).

The nearly concentric geometrical relationship betweenOdysseus and Ithaca Chasma prompted Moore and Ahern(1983) to suggest that the trough system was an immedi-ate manifestation of the impact event, perhaps caused by adamped, whole-body oscillation of the satellite. Based onthe deep topography of Odysseus, Schenk (1989b) also sug-gested that it had not undergone substantial relaxation, andthat Ithaca Chasma was the equivalent of a ring grabenformed during the impact event by a prompt collapse of thefloor involving a large portion of the interior. In a Cassini-era study, Giese et al. (2007) doubted the Odysseus-related


N630

–3–5

10

Ele

vatio

nm, k

m

0

0

0

0 100

200 km

200 300

vertic. exagg. = 3

km

Te = 4.9 km

Ithaca chasma Te = 7.2 km

p3p2

p1

p3

p2

p1

Elevation, km

Fig. 20.2 Tethys and Ithaca Chasma. Topography was derived fromCassini stereo images and has a horizontal resolution of 2–5 km anda vertical accuracy of 150–800 m. Profiles have superimposed model

flexural profiles (in red). Te is elastic thickness (reproduced from Gieseet al. (2007))

hypothesis (Moore and Ahern 1983; Schenk 1989b) on thebasis of crater densities within Ithaca Chasma relative toOdysseus, from which they concluded that Ithaca Chasmais older than Odysseus and therefore could not have influ-enced its formation. They went on to measure photogram-metrically the raised flanks of Ithaca Chasma, which theyattributed to flexural uplift. Assuming that these raised flankswere indeed due to flexural uplift, they derived a mechanicallithospheric thickness of �20 km and surface heat fluxes of18–30mW=m2 at assumed strain rates of 10�17 to 10/s. Theyconcluded that Ithaca Chasma is an endogenic tectonic fea-ture but could not explain why it is confined to a narrow zoneapproximating a large circle.

Several Voyager-era studies of Dione (Plescia 1983;Moore 1984) noted a near-global network of tectonic troughsand the presence of a smooth plain crossed by both ridgesand troughs located near the 90 ıW longitude and extend-ing eastward. Cassini coverage revealed the westward exten-sion of this plain (Wagner et al. 2006) and demonstrated thatthe fractured and the cratered dark plains are spectrally dis-tinct (Stephan et al. 2008). Like the eastern portion observedby Voyager, the western plain also exhibits both ridges andtroughs (Fig. 20.3).

However, the ridge system is much more pronounced. Thewestern ridges are planimetrically narrower and morpholog-ically better expressed (e.g., their relative lack of degrada-tion). These western ridges also conform much more closelyto the boundary between the plain and the older crateredrolling terrain to the west. The linear to arcuate troughs in-

terpreted to have been created by tectonic extension forma complex network that divides Dione’s surface into largepolygons (Smith et al. 1981, 1982; Plescia 1983; Moore1984). The troughs often form parallel to sub-parallel sets,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 witha decline in oblateness due to either despinning or orbital re-cession (Melosh 1977, 1980).

The topography derived from Cassini images of Dioneshows one clear case of tectonic orientation regionally influ-enced by a preexisting but very degraded large impact basincentered on the smaller crater Amata (Schenk and Moore2007 (Figs. 20.3 and 20.4)).

Moore et al. (2004) identified another example of tectoniclineament orientation associated with an unambiguouslyidentified impact feature, 90 km Turnus. Such crater-focusedtectonics has been identified on Ganymede from Galileo data(e.g., Pappalardo and Collins 2005). This crater-lineament re-lationship on Dione is the most pronounced among any of themiddle-sized icy satellites except Enceladus (Spencer et al.2009). Moore (1984) proposed a tectonic history of extensionfollowed by (at least regional) compression, the latter beingresponsible for the ridges. Cassini-era images of the tectonictroughs (mostly grabens) show them to be morphologicallyfresh, implying that extension continued into geologically re-cent times (Wagner et al. 2006).

After Voyager, Rhea was held to be the least endogeni-cally evolved satellite larger than 1,000 km in diameter.


Fig. 20.3 (a) Image mosaic of Dione, centered on 180ı longitude. Redand white arrows point to Amata and Turnus impact structures, respec-tively. (b) Topography of Dione, from Schenk and Moore (2007). An

ancient degraded impact basin centered just west of Amata can be iden-tified by the two concentric rings in the DEM (white arrows). DEM hasa dynamic range of 8 km (red is highest, magenta is lowest)

Moore et al. (1985) pointed out a few parallel lineamentsand noted several large ridges or scarps (which theytermed ‘megascarps’) and speculated that these features wereformed by a period of extension followed by compression.Cassini images, and digital terrain models derived fromthem, of the trailing hemisphere of Rhea seen from both mid-dle northern and far southern latitudes clearly show that thewispy arcuate albedo markings seen in Voyager images areassociated with a graben/extensional fault system dominantlytrending north-south (Schenk and Moore 2007; Wagner et al.2007) (Fig. 20.5a).

This is the first identification of unambiguous regionalendogenic activity on Rhea. The digital terrain models ofRhea derived from Cassini images also show a set of roughlynorth-south trending ridges (Figs. 20.5b and 20.6) that corre-spond to the ‘megascarps’ of Moore et al. (1985).

If these ridges are in fact due to compressional tecton-ics, their geometrical relation to the graben/scarp system mayalso indicate a genetic relationship. However, the poor mor-phological presentation of these ridges implies that they areold, while the ‘wispy terrain’ grabens are fresh and presum-ably much more recently formed.


Fig. 20.4 Fracture zones on Dione. Global base map of Cassini imag-ing has a resolution of 400 m. The circle represents the outer ridge corre-sponding to the rim of the degraded basin near Amata shown in Fig. 20.3

While the troughs and ridges suggest mainly extensionaland (minor) compressional tectonics (Thomas 1988), theen-echelon pattern of the scarps and troughs on the trailinghemisphere suggests shear stress (Wagner et al. 2007). Thispattern may have been influenced by the possible presence ofa large, degraded impact basin (basin C of Moore et al. 1985).

The equatorial ridge of Iapetus (Fig. 20.7) is the mostenigmatic feature on any of the medium-sized icy satellites(Porco et al. 2005a; Giese et al. 2008); it extends across morethan half of Iapetus’ circumference (Giese et al. 2008; Denket al. 2008). Different parts show very different morpholo-gies, from trapezoidal to triangular cross-sections with steepwalls and a sharp rim to isolated mountains with moderateslopes and rounded tops (Giese et al. 2008). Although theridge is variable in height, it rises up to 10 km above the lo-cal mean but also has a gently sloping flank. The ridge isabundantly cratered, indicating that it is comparable in ageto other terrains on Iapetus (Schmedemann et al. 2008). Al-though the equatorial location of the ridge might be due todespinning (Castillo-Rogez et al. 2007), it represents such alarge mass that, if it had formed in another location, it would

Fig. 20.5 (a) Cassini image ofthe northern trailing hemisphereof Rhea showing the brightscarps of the northern extensionof the graben system and theirrelationship to ‘wispy’ terrain.(b) Mosaic of Cassini imagesshowing curvilinear grabens(black arrows) and an orthogonalset of ridges and troughs (whitearrow). Base mosaic resolution is1 km/pixel

Fig. 20.6 Preliminary stereotopography map, from Schenkand Moore (2007). White arrowsindicate graben systems. TheDEM has a dynamic range of10 km (red is highest, magenta islowest)


(730) Height, km

p_glob

ridge

14

30˚

60˚

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200 km

Hei

ght,

km

302010

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0 200 400 600 800 1000 1200 1400 1600 1800 km

vert, exag,=6p_glob

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(748 x 748 x 713) Height, km

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300˚

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0˚

(748 x 748 x 713) Height, km1470

–7–12

N31.520.010.0

–12.50.0

Fig. 20.7 Topography of Iapetus as derived from stereo Cassini imaging. Note the sharply defined equatorial ridge. The DEM at left is referencedto a sphere, the DEM at right to a biaxial ellipsoid (adapted from Giese et al. 2008)

almost certainly have reoriented to the equator anyway. How-ever, as already mentioned above, irrespective of whether asatellite is spinning down or spinning up, equatorial tectonicfeatures can be expected to be oriented north-south, makingit difficult to explain the ridge with despinning. One possi-bility is that despinning stresses combined with lithosphericthickness variations due to convection may have resulted inequatorial extension and diking (Roberts and Nimmo 2009;Melosh and Nimmo 2009). Other proposals suggest thatthe morphology of the ridge indicates that the surface waswarped up by tectonic faulting (Giese et al. 2008), that it is aresult of extensional forces acting above an ascending currentof solid-state convection from a two-cell convection pattern(Czechowski and Leliwa-Kopystynski 2008), or that it wasformed by the impact of an ancient debris disk on Iapetus (Ip2006). The great apparent elastic thickness suggested by thesurvival of the ridge may be used to place constraints on thethermal evolution of Iapetus (Castillo-Rogez et al. 2007).

Although the global shape of Iapetus does suggest an ear-lier, more rapid spin rate (Castillo-Rogez et al. 2007), despin-ning cannot explain the equatorial ridge (see above): stressesin the north-south direction are always smaller than those inthe east-west direction, and will result in N-S oriented fea-tures (Melosh 1977). Thus, the origin of the equatorial ridgeis currently a mystery.

The deformation of Saturnian satellites varies wildly,from currently active and heavily deformed Enceladus to

nearby Mimas, which is tectonically uninteresting. Nor doesthe degree of deformation correlate in any straightforwardway with predicted tidal stresses. However, there are at leasttwo global yet somewhat contradictory aspects, which maybe identified. On the one hand, coherent global tectonicpatterns are evident on Dione, Tethys and Enceladus andto a lesser extent on Rhea and Iapetus. On the other, sev-eral satellites, especially Tethys and Enceladus, show verypronounced spatial variations in the extent of deformation.These two aspects suggest that global sources of stress arecommon, but also that lateral variations in the mechani-cal properties of the satellites’ ice shells play an impor-tant role.

Regarding this second point, localization of deformation(e.g., at the south pole of Enceladus or at Ithaca Chasma) mayarise naturally in ice shells (e.g., Sotin et al. 2002; Nimmo2004b) but is currently poorly understood and makes mod-eling tectonic deformation much harder. As far as globaldeformation patterns are concerned, several satellites (Ence-ladus, Iapetus, perhaps Dione) show patterns with simplesymmetries about the rotational and tidal axes. This symme-try may be due either to feature reorientation or the operationof mechanisms (tides, despinning etc.), which have a natu-ral symmetry. On Dione, Rhea and Iapetus at least, diurnaltidal stresses are rather small (Table 20.3), and most of thetectonic features are apparently ancient, which suggests thatdespinning, volume changes or reorientation are more likely


to be responsible for the observed deformation than present-day tidal stresses.

As in the Galilean satellites, extensional features are gen-erally more common than compressional features on theother Saturnian satellites. This observation is readily ex-plained if we assume that the satellites once contained sub-surface oceans which progressively froze (e.g., Smith et al.1981; Nimmo 2004a); alternative explanations are that com-pressional features such as folds (Prockter and Pappalardo2000) are less easy to recognize than extensional features, orthat larger stresses are required to form compressional fea-tures (see Pappalardo and Davis 2007).

20.3.3 Cryovolcanism

Cryovolcanism can be defined as ‘the eruption of liquid orvapor phases (with or without entrained solids) of wateror other volatiles that would be frozen solid at the normaltemperature of the icy satellite’s surface’ (Geissler 2000).Cryovolcanism therefore includes effusive eruptions of icyslurries and explosive eruptions consisting primarily of vapor(Fagents 2003). Voyager-era data suggested that such activ-ity might be common in the outer solar system, resulting in aflurry of theoretical papers (e.g., Stevenson 1982; Allison andClifford 1987; Crawford and Stevenson 1988; Kargel 1991;Kargel et al. 1991). Higher-resolution images from Galileogenerally revealed tectonic rather than cryovolcanic resurfac-ing processes, although some features may be due to cryovol-canic activity (e.g., Schenk et al. 2001; Fagents et al. 2000;Miyamoto et al. 2005). In the following, we discuss the cur-rent state of knowledge about cryovolcanism in the Saturniansatellites. Enceladus, which is definitely cryovolcanically ac-tive (Porco et al. 2006; Dougherty et al. 2006), and Titan,which may be (Lopes et al. 2007; Nelson et al. 2009a,b), arediscussed by Spencer et al. (2009), Jaumann et al. (2009) andSotin et al. (2009).

The principal difference between cryovolcanism and sili-cate volcanism is that in the former, the melt is denser thanthe solid (by �10%). This means that special circumstancesmust exist for an eruption of non-vapor cryovolcanic ma-terials to occur (though of course on Earth it is not un-common for denser basaltic lavas to erupt through lightergranitic material). Such circumstances include one or moreof the following: (1) the melt includes dissolved volatiles,which exsolve and reduce the bulk melt density (Crawfordand Stevenson 1988); (2) the melt is driven upwards by pres-sure differences caused by tidal effects (Greenberg et al.1998), surface topography (Showman et al. 2004) or freez-ing (Fagents 2003; Manga and Wang 2007); (3) composi-tional differences render the ice shell locally dense or themelt less dense (e.g., by the addition of ammonia) (Kargel

1992, 1995). The latter mechanism can be nothing more thana local effect, since if the bulk ice shell is denser than theunderlying material; it will sink if the ice shell is disrupted.

Erupted material consisting primarily of water plus solidice will simultaneously freeze and boil, the latter processceasing once a sufficiently thick ice cover is formed (Allisonand Clifford 1987). The resulting flow velocity will dependon the viscosity of the ice-water slurry (Kargel et al. 1991) aswell as on flow thickness and local slope (Wilson et al. 1997;Miyamoto et al. 2005). If the erupted material is dominatedby the vapor phase, the vapor plume will expand outwards(Tian et al. 2007) and escape entirely if the molecular thermalvelocities exceed the escape velocity. Any associated solidmaterial will follow ballistic trajectories, ultimately formingdeposits centered on the eruption vents (Fagents et al. 2000).

The plume properties measured on Enceladus by Cassiniimply considerably smaller velocities for icy dust grains thanfor vapor. On hypothesis is that the dust grains condense andgrow in vertical vents of variable width. Repeated wall col-lisions and re-acceleration by the gas in the vents lead toeffective friction, which explains the observed plume prop-erties (Schmidt et al. 2008).). Most dust grains are growingby direct condensation of the ascending water vapor withinthe ice channels and contain almost pure water ice; they arefree of atomic sodium. In 2008 the CDA instrument whileanalyzing the composition of E-ring particles discovered asecond type of icy dust particles. These grains are rich insodium salts (0.5–2% in mass) and there origin is explainedby the existence of an ocean of liquid water below the surfaceof Enceladus and which is in direct contact with its rock core(Postberg et al. 2009). This shows that cryovolcanism canprovide valuable information about geological conditions ofthe interior.

The plains of Dione are another candidate for cryo-volcanic resurfacing (Fig. 20.8). Several Voyager-era stud-ies concluded that these plains were 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 60ıW may have beenthe source vents. Surprisingly, the plains unit may be to-pographically higher than the cratered terrain (Schenk andMoore 2007). This observation seems a little difficult to rec-oncile with the plains being emplaced as a low-viscosity cry-ovolcanic flow. One possibility is that the plains materialwas emplaced as a deformable plastic unit (like terrestrialglaciers), but this, or any other, explanation needs both fur-ther evidence and thought. The ridges discussed above bound(or at least occur near the edges) the plains unit in the westand southeast and may be related to their formation.

A plains unit on Tethys is located antipodal to the verylarge Odysseus impact feature. Moore and Ahern (1983) pro-posed that the plains were formed by cryovolcanism. Alter-natively, Moore et al. (2004) considered the possibility that it


Fig. 20.8 Cassini view of smooth plains on Dione. Older heav-ily cratered terrains lie to the west. Note the prominent north-southridge and the subtle linear features to the east. Mosaic resolution is400 m/pixel

was extreme seismic shaking caused by energy focused thereby the Odysseus impact that formed the antipodal plains. Ifseismic shaking was responsible, there would probably bea significant transition zone between that unit and the adja-cent hilly and cratered terrain. Digital terrain models derivedfrom Cassini images indicate that the transition between thecratered terrain and the plains is abrupt and the plains them-selves are level, supporting the hypothesis that they are en-dogenic (Schenk and Moore 2007).

Both Tethys and Dione show evidence for weak plasmatori (Burch et al. 2007) somewhat analogous to the muchmore marked feature at Enceladus, which is caused by gey-sers (Porco et al. 2006). It is not yet clear whether surfaceactivity or some other process is generating these tori. Simi-larly, Clark et al. (2008a) report a possible, but currently un-confirmed, detection of material around Dione, which mayalso indicate some surface activity. The Cassini magnetome-ter has also measured small amounts of mass loading nearDione (Burch et al. 2007).

So far, none of the digital terrain models of Rhea devel-oped in the Cassini era shows any regions that would qual-ify as plains, supporting the perception that Rhea shows norecord of cryovolcanic resurfacing. However, unlike Dione,the diffuse nature of the wisps (the ‘wispy terrain’) on Rhea

has yet to be explained as numerous smaller bright scarpsassociated with fracture and faulting, and may yet be shownto be an expression of cryo-pyroclastic mantling (Stevenson1982). Furthermore, Rhea, surprisingly, exhibits evidence forfaint particle ring (Jones et al. 2008), the origin of whichis unclear but may be related to endogenic or impact pro-cesses. On the other hand, Cassini VIMS observations foundno evidence of a water vapor column around Rhea (Pitmanet al. 2008). On Iapetus, no evidence of cryovolcanism hasbeen detected thus far. If volcanism ever operated on thesesatellites, it occurred in ancient times and its evidence is nowunseen.

20.4 Composition and Alteration of SurfaceMaterials

The composition of the satellites of Saturn is determined byremote spectroscopy, by sampling materials that were sput-tered from the satellite surfaces or, in the case of Enceladus,injected into space where they can be directly sampled byspacecraft instruments.

It has long been known that the surfaces of Saturn’s majorsatellites, Mimas, Enceladus, Tethys, Dione, Rhea, Hyper-ion, Iapetus and Phoebe, are predominantly icy objects (e.g.,Fink and Larson 1975; Clark et al. 1984, 1986; Roush et al.1995; Cruikshank et al. 1998a; Owen et al. 2001; Cruikshanket al. 2005; Clark et al. 2005, 2008a; Filacchione et al. 2007,2008). While the reflectance spectra of these objects in thevisible range indicate that a coloring agent is present on allsurfaces except Mimas, Enceladus and Tethys (Buratti 1984).Only Phoebe and the dark hemisphere of Iapetus displaynear-IR spectra markedly different from very pure water ice(e.g., Cruikshank et al. 2005; Clark et al. 2005).

The major absorptions in icy satellite spectra shown inFigs. 20.9 a–e are due to water ice and can be observed at1:5 �m:; 2:0 �m:; 3:0 �m., and 4:5 �m. Weaker ice absorp-tions appear at 1:04 �m and 1:25 �m in the spectra of brightregions where the ice is purer. Trapped CO2 was first detectedin Galileo near-infrared mapping spectrometer (NIMS) spec-tra of the Galilean satellites (Carlson et al. 1996) and in theSaturnian system on Iapetus (Buratti et al. 2005), Phoebe(Clark et al. 2005), Hyperion (Cruikshank et al. 2007), andEnceladus (Brown et al. 2006). Weak CO2 absorptions werealso detected in VIMS data on Dione, Tethys, Mimas andRhea (Clark et al. 2008a).

Organic compounds were detected directly in the Ence-ladus plume by the INMS (Waite et al. 2006) and in traceamounts by the VIMS on Phoebe (Clark et al. 2005), Hype-rion (Cruikshank et al. 2007) and Iapetus (Cruikshank et al.2008), and possibly Enceladus (Brown et al. 2006). While thespectral signatures indicate aromatic hydrocarbon molecules,


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Fig. 20.9 VIMS spectra of the Saturnian icy satellites: (a) The majorabsorptions in icy satellite spectra are due to water ice and appear at 1.5,2.0, 3.0, and 4:5�m; (b) Representative (A) bright and (B) dark regionDione spectra (the gaps in the bright region spectra (A) near 1�m aredue to sensor saturation); the inset shows an area of high concentrationof CO2 (after Clark et al. 2008a); (c) Spectra of Hyperion showing anincrease in blue reflectivity and absorptions of water ice and CO2 (after

Clark et al. 2008a); (d) Spectra of Iapetus showing a range of blue peakintensities, absorptions of different intensities by water ice and CO2 ab-sorption bands (after Clark et al. 2008a); (e) Spectra of Phoebe showinga range of blue peak intensities and absorptions of different intensitiesby water ice and CO2 absorption bands (after Clark et al. 2008a). How-ever, no CO2 was found on the small satellites Atlas, Pandora, Janus,Epimetheus, Telesto, and Calypso (Buratti et al. 2009a)


exact identifications have remained elusive, partly becausethe detections have low signal-to-noise ratios. Another spec-tral feature observed in VIMS data of dark material onPhoebe, Iapetus and Dione appears to be an absorption cen-tered near 2:97�m (Fig. 20.9b,d,e). The detection of thisfeature is complicated by the fact that VIMS has an order-sorting filter change at this wavelength. Clark et al. (2005)presented evidence that this feature is real and maps to ge-ological patterns on Dione and Iapetus (Clark et al. 2008a,2009). However, a better understanding of the instrument re-sponse is needed to increase confidence in the identification,including confirmation by a different instrument. If correct,the feature strength indicates approximately 1% ammonia inthe dark material.

The position shift of amorphous ice absorptions from theircrystalline counterparts (Mastrapa et al. 2008 and referencestherein) can be used to determine whether amorphous orcrystalline ice is present on a surface (Newman et al. 2008).Traditionally, amorphous ice displays almost no 1:65 �mfeature and the 3:1�m Fresnel peak shifts, broadens andweakens. Newman et al. (2008) used variations in the 3:1�mpeak and the 1:65 �m absorption to argue for amorphousice around the ‘tiger stripes’ on Enceladus. However, Clarket al. (2009) showed that scattering and sub-micron grainsizes also influence the intensities of the 3:1-�m peak andthe 1:65-�m absorption, and can be confused with tempera-ture and crystallinity effects. Further examination of the icefeature positions in spectra of Enceladus and other satellitesin the Saturnian system showed that positions and strengthsare consistent with those of fine-grained crystalline ice (Clarket al. 2009). Newman et al. (2009) showed that water ice onDione exists exclusively in crystalline form.

The spectra of the satellites in the visible wavelengthrange are smooth with no sharp spectral features, but all showan ultraviolet absorber affecting the blue end of the visible-wavelength spectral slope. Visible spectra can be classifiedby their spectral slope: from bluish Enceladus and Phoebeto redder Iapetus, Hyperion and Epimetheus. In the 1�m to1:3�m range, the spectra of Enceladus, Tethys, Mimas andRhea are characterized by negative slope, consistent with asurface mainly dominated by water ice, while the spectra ofIapetus, Hyperion and Phoebe show considerable reddening,pointing out the relevant role played by the darkening mate-rials present on the surface (Filacchione et al. 2007, 2008).In between these two classes are Dione and Epimetheus,which have a relatively flat global spectrum in this range. Atwavelengths less than about 0:5 �m, all satellite spectra dis-play a UV absorber, which might be due to charge transferin oxides such as nano-phase hematite (Clark et al. 2008a),with the apparent strength of the absorber generally indicat-ing the amount of contaminant. Trends in the UV absorberare complicated by ice dust contributions to satellite surfacestransmitted from Enceladus via the E-ring. Satellites close

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Fig. 20.10 Cassini UVIS disk-integrated spectra of the icy moons, allat phase angles between 10ı and 13.6

to the rings, such as Janus, Atlas and Prometheus, show anintermediate red tint between that of the rings and Enceladus.Mimas and Tethys show a red tint between that of Janus,Atlas and Prometheus and that of Enceladus on the other(e.g., Cuzzi et al. 2009; Buratti et al. 2009a).

Disk-integrated ultraviolet observations by the CassiniUVIS (Esposito et al. 2004; Hendrix and Buratti 2009)(Fig. 20.10) show that Tethys and Dione exhibit markedUV leading-trailing differences whereas Mimas, Enceladusand Rhea do not, suggesting compositional variations amongthe satellites as well as spectrally active components inthe ultraviolet. UVIS spectra of all the icy moons of Sat-urn are dominated by a pronounced water-ice absorptionedge at �0:165�m. Its strength varies with the abundanceof water ice, and its wavelength depends on grain size.Nearly all UVIS spectra display a small absorption featureat �0:185�m (Hendrix and Hansen 2008b) whose originis not clear at present. It is possible that this feature is dueto water ice as it is seen in published reflectance spectra ofwater ice (Pipes et al. 1974; Hapke 1981). Published opticalconstants (Warren 1982) do not exhibit the feature, but theconstants have not been well measured in the �0:17 �m to0:4 �m range.

Cassini spectral coverage of the icy moons is complete inthe 0:4 �m to 5:2 �m and in the 0:1 �m to 0:19 �m region,but there is a gap in the 0:19 �m to 0:4 �m region (except forthe broad band spectral measurements of ISS). This is an im-portant wavelength region for non-water ice absorptions, andthe Cassini UVIS data in the far UV provide a critical tie-point between the far UV and the visible band that is helpfulin understanding the nature of the non-ice species present inthe surfaces of the moons. The icy moons are dark in the farUV compared to their spectra in the Vis-NIR. For instance,the disk-integrated albedo of Enceladus at 0:19 �m is �0:4,while at 0:439�m (Verbiscer et al. 2005) it is �1:2. Water ice


is thought to have a relatively flat spectral signature in thiswavelength range, which suggests that some other speciesdarkens Enceladus in the FUV. Ammonia, which exhibits anabrupt absorption edge at 0:2 �m (Pipes et al. 1974), is a can-didate (and ammonia hydrate has been tentatively detectedat Enceladus and Tethys in Earth-based spectra (Verbisceret al. 2005; Verbiscer et al. 2008)), though at this time otherspecies cannot be ruled out other materials.

20.4.1 Dark Material

Cassini (1672) was the first to infer the presence of darkmaterial in the Saturnian system, which was verified byMurphy et al. (1972) and Zellner (1972). Dark material ap-pears throughout the system, most predominantly on Dione,Rhea, Hyperion, Iapetus and Phoebe. The nature of the darkmaterial especially on Iapetus has been studied by numer-ous authors, often with conflicting conclusions, includingCruikshank et al. (1983), Wilson and Sagan (1995, 1996)Vilas et al. (1996), 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 satel-lites’ spectra, particularly in the visible. In dark-region spec-tra, additional absorptions are seen at 2:42 �m and 4:26 �m,with the latter due to CO2, as well as stronger water ice ab-sorption near 3�m (Fig. 21.5.1d). The 2:42 �m feature wasfirst tentatively identified as a CN overtone band by Clarket al. (2005) in spectra of Phoebe, but evidence of CN funda-mentals in the 4�m to 5�m region has not been confirmed.At present, the 2:42 �m feature is believed to be due to OHor hydrogen (Clark et al. 2008a, b, 2009), but complicationswith the VIMS calibration have limited any detailed defini-tion of the absorption.

Diverse spectral features of the dark material on Dionematch those seen on Phoebe, Iapetus, Hyperion andEpimetheus as well as in the F-ring and the Cassini Division,implying that its composition is the same throughout the Sat-urnian system (Clark et al. 2008a, 2009). Models of the darkmaterial on Iapetus have used tholins to explain the reddishvisible spectral slope (e.g., Cruikshank et al. 2005). However,VIMS data, which better isolate the spectral signature of thedark material from that of icy regions thanks to the high spa-tial resolution afforded by close fly-bys, show that the darkmaterial has a remarkably linear spectral increase in the nearinfrared (Clark et al. 2008a, 2009) (Fig. 20.12) not explainedby tholins which have a non-linear spectral response in thiswavelength region. Tholins also have strong CH absorbersin the 3–4-�m region as well as other marked absorptionsat wavelengths longer than 2�m that are not seen in, forexample, spectra of the dark regions of Iapetus. The prob-lem of matching one spectrum might be accomplished with

a specific mixture of many compounds, but with spatially re-solved spectra on the satellites we now see a large range ofmixtures, so models must show consistency across that rangeof observed mixtures, and many models do not.

The source of the dark material in the Saturnian system isunclear. Several researchers (e.g., Clark et al. 2005) suggest asource outside the system, perhaps the Kuiper belt or comets.Owen et al. (2001) suggested that an impact on Titan cre-ated a spray of dark, organic-rich material, part of which wasswept up by Iapetus’ leading hemisphere, another part beingtransported to the inner system to accrete onto the surfaces ofthe inner icy moons. The Dione VIMS spectra provide a keyto solving some of the mysteries of the dark material. Clarket al. (2008a) showed that a pattern of bombardment by fineparticles below 0:5 �m impacted Dione from the trailing-sidedirection. Several lines of evidence point to an external originof the dark material on Dione, including its global spatial pat-tern, local patterns including crater and cliff walls shieldingimplantations on slopes facing away from the trailing side,exposing clean ice, and slopes facing the trailing directionthat show higher abundances of dark material.

A blue scattering peaks is seen in the spectrum of thedark material and was initially observed in the spectra ofTethys, Dione and Rhea by Noland et al. (1974). Clark et al.(2008a) suggest that the blue peak is caused by particles lessthan 0:5 �m in diameter embedded in the ice. The VIMS re-solved areas on Phoebe, Iapetus and Hyperion also show bluepeaks (Fig. 20.9 c–e), indicating that this is a common prop-erty of satellites in the Saturnian system (Clark et al. 2008a,2009). The blue scattering peak with a strong UV/visible ab-sorption was modelled with ice plus 0:2 �m diameter car-bon grains (to provide the Rayleigh scattering effect) plusnano-phase hematite .Fe2O3/ to produce the UV absorber(Fig. 20.11). But carbon is inconsistent with the linear-redslope seen in the purest dark material spectra from 0:4 �mto 2:5 �m. Metallic iron is the only known compound todisplay this spectral property and nano-phase metallic ironplus nano-phase hemtatite mixed with ice show consistencywith the mixture spectra observed on the icy satellites and inSaturn’s rings (Clark et al. 2009).

Previous models of the dark material on Iapetus typicallyused tholins (e.g., Cruikshank et al. 2005). However, VIMSdata, which better isolate the spectral signature of the darkmaterial from that of icy regions thanks to the high spatialresolution afforded by close fly-bys, show that the dark ma-terial has a remarkably linear spectrum (Clark et al. 2008a,2009) (Fig. 20.12).

Tholins have strong CH absorbers in the 3–4-�m regionas well as other marked absorptions at wavelengths longerthan 2�m that are not seen in, for example, spectra of thedark regions of Iapetus.

Using VIMS data, Clark et al. (2008a) found the gen-eral composition of the dark material in the Saturnian system


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likely contains bound water, CO2 and, tentatively, ammonia.More recently, Clark et al. (2008b, 2009) argued that thesignature of iron identifies the major coloring agent on theicy surfaces in the Saturnian system. Metallic iron does notexhibit sharp diagnostic spectral features, but it does haveunique spectral properties in that its absorption increases lin-early with decreasing wavelength (Clark et al. 2009).

The theory of Clark et al. (2008b, 2009) envisages small-particle meteoroid dust contaminating the system as well asmeteoroids with a high metallic-iron content. As one movescloser to the rings, the dark material diminishes (e.g., Mimasand Enceladus have little dark material). However, the com-peting process of accretion by Saturn’s E-ring, which origi-nates on Enceladus, also contributes to the greater abundanceof high-albedo material on Mimas and Enceladus (Burattiet al. 1990, Verbiscer et al. 2007), The main rings do notshow any strong signature of metallic iron particles, but theCassini Division displays spectra that closely match those ofIapetus in Iapetus’ transition zone from dark to light material,implying mixtures of the same materials. The ring spectra aremuch redder in the 0:3 �m to 0:6 �m region than the spectraof dark material on Phoebe, Iapetus and Hyperion. Cosmicrays impacting the huge surface area of the rings break upwater molecules, thus creating an atmosphere of hydrogenand oxygen. The hydrogen escapes, leaving an atmosphere ofoxygen around the rings (Johnson et al. 2006). Oxygen ions

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Fig. 20.12 Spectrum of the dark material of Iapetus compared to lab-oratory spectra of mixtures of iron powder, H2O; CO2 and NH3. Thesmall Rayleigh scattering peak and UV absorber can be explained bytraces of iron oxide nano-powder. The 3-micron absorption is best ex-

plained by a combination of bound water, ice and trace ammonia. Finer-grained iron would require higher abundances of water, ammonia andcarbon dioxide to match the spectrum of Iapetus


are very reactive, and when they encounter iron particles,they oxidize them and turn them into nano-phase hematite.Fe2O3/. The spectra of nano-phase hematite mixed with iceand dark particles closely match the UV red slope seen inthe spectra of Dione (Clark et al. 2008a) (Fig. 20.11). To ex-plain the dark material in the Saturnian system, Clark et al.(2008b, 2009) suggest that oxidized sub-micron iron parti-cles from meteorites, mixed with ice, might explain all thecolors and spectral shapes observed in the system. A smallamount of carbon might also be identified in the mixture,which would change the slope of iron and further lower thealbedo. Sub-micron particles also explain the origin of theobserved Rayleigh scattering. Trace amounts of CO2, ammo-nia and organics would not significantly influence the visibleand UV spectra down to 0:3 �m. Finally, the inner satellitesmay be exposed to as much iron as the outer satellites, butE-ring ice particles mix with and cover up the iron particles.Species found by INMS, CAPS and MIMI in the magneto-sphere are candidates for surface components as the mag-netospheric material originates either directly from satellitesurfaces or from ring particles. Compositional analyses ofEnceladus’ plume are described in Spencer et al. (2009).

The idea of iron in the Saturn system is supported bythe Cassini CDA instrument, which has discovered small.<20 nm/ particles escaping the Saturnian system, which arecomposed predominantly of oxygen, silicon and iron, withsome evidence of H2O ice, ammonium, and possibly carbon(Kempf et al. 2005). The escaping iron is not in a silicate and

is not metallic; iron oxide is most consistent with the data(Kempf et al. 2008, personal communication). Iron has alsobeen detected in the E-ring by CDA (Kempf et al. 2008). Theorigin of the iron may be meteoric dust falling into the Sat-urnian system, and may indicate that space weathering pro-cesses operate there as in other parts of the solar system (e.g.,Chapman 2004).

20.4.2 Iapetus’ Hemispheric Dichotomy

Iapetus has intrigued planetary scientists for centuries, pri-marily due to its striking hemispheric albedo dichotomy.The leading hemisphere (centered exactly on the apex ofmotion at 90 ıW) is very dark, reflecting just �5% of thevisible light that hits it (at 3ı phase angle), while the trail-ing hemisphere (centered on 270 ıW), is relatively brightand has a visible albedo of �50% to �60% (Fig. 20.13). Along-standing question has been how the global albedo di-chotomy of

Iapetus originated. Has the leading hemisphere’s dark ter-rain been created by exogenic processes, as most scientistshypothesized (e.g., Cook and Franklin 1970; Soter 1974;Cruikshank et al. 1983; Squyres and Sagan 1983; Bell et al.1985; Tabak and Young 1989; Matthews 1992; Buratti andMosher 1995; Wilson and Sagan 1996; Denk and Neukum2000; Owen et al. 2001; Buratti et al. 2002, 2005a), or

Fig. 20.13 Patterns of the global dichotomies on Iapetus (Denk et al.2009). 1 : : : Uniformly dark, reddish terrain ‘deep within’ Cassini Re-gio; 2 : : : Dark, reddish terrain similar to (1) but showing bright slopes(esp. crater walls) facing poleward; 3 : : : Bright, reddish terrain on theleading side, with dark spots located on equatorward-facing slopes; 4: : : Dark terrain, but less reddish (at visible to IR wavelengths) than (1);5 : : : Bright equatorial trailing-side terrain containing numerous darkspots colored similarly to (4); 6 : : : Low- to mid-latitude bright terrainwith multiple dark equatorward-facing slopes; 7 : : : Bright ‘polar caps’on the trailing side, no dark spots, flattest spectra (near UV to near IR)

of all Iapetus terrains. These might be the sinks for the water ice in thethermal segregation process. The global color dichotomy separates ar-eas (1)–(3) (leading side; redder color in visible to IR wavelengths) from(4) to (7) (trailing side; less reddish). The global brightness dichotomyseparates areas (1), (2) and (4) (mainly dark terrain; named Cassini Re-gio) from (3), (5), (6) and (7) (mainly bright terrain, named RoncevauxTerra and Saragossa Terra). Latitudinal inspection: the differences be-tween areas (1), (2) and (3) result from the thermal segregation of waterice, as do the differences between areas (5), (6) and (7)


did geological activity emplace dark material from withinIapetus (Smith et al. 1981, 1982)? Voyager images of dark-floored craters within the bright terrain point to an endogenicsource; they also suggest that the bright-dark boundary istoo irregular to be consistent with infalling dust (Smith et al.1981, 1982; Denk and Spencer 2008). However, albedo pat-terns observed by Cassini cameras in late 2004 suggest exter-nal material emplacement (e.g., dark material on ram-facingcrater walls at high latitudes) (Porco et al. 2005a), but thiswas later found to be due to photometric effects. The ini-tial theory of an exogenically-created dark pattern (Cook andFranklin 1970) suggested that pre-existing dark material wasuncovered by meteoritic bombardment; this idea was extrap-olated by Wilson and Sagan (1996). Researchers theorized(Soter 1974) that the dark material is exogenically emplacedon Iapetus’ leading hemisphere as material is lost from themoon Phoebe (Burns et al. 1979). Retrograde Phoebe dustfrom a distance of 215 Saturn radii would travel inwardfrom the effects of Poynting-Robertson drag and impact theleading hemisphere of Iapetus orbiting at 59 Saturn radii.However, Phoebe is spectrally neutral at visible wavelengths,while the Iapetus dark material is reddish (Cruikshank et al.1983; Squyres et al. 1984). If the material does come fromPhoebe, then some sort of chemistry or impact volatiliza-tion must occur to change the color and darken the mate-rial (Cruikshank et al. 1983; Buratti and Mosher 1995). An-other possibility is that the exogenic source of the dark mate-rial is Hyperion (Matthews 1992; Marchi et al. 2002), Titan(Wilson and Sagan 1995; Owen et al. 2001), Iapetus itself(Tabak and Young 1989) or the debris remnants of a collisionbetween a former outer moon and a planetocentric object(Denk and Neukum 2000). Matthews (1992) suggested thatthe hypothetical impact that disrupted Hyperion created a de-bris cloud that subsequently hit Iapetus. Both Hyperion andTitan dark materials are spectrally reddish (Thomas and Vev-erka 1985; McDonald et al. 1994; Owen et al. 2001), thoughnot as dark as those of Iapetus. Buratti et al. (2002, 2005a)suggested that both Hyperion’s and Iapetus’ leading hemi-spheres are impacted by dark, reddish dust from recently dis-covered 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) whenmixed in small amounts with water ice produces a Rayleighscattering effect, enhancing the blue response. This blue en-hancement appears just enough on Phoebe to create an ap-proximately neutral spectral response. An examination of thespectral signatures of Dione, Hyperion, Iapetus and Phoebeby Clark et al. (2008a) showed a continuous mixing space onall these satellites, demonstrating variable mixtures of fine-grained dark material in the ice. Thus, the colors observedare simply variable amounts of dark material, with the red-der slopes indicating greater abundance.

Ground-based radar observations at 13 cm (Black et al.2004) and Cassini RADAR data at 2.2 cm (Ostro et al.2006; 2009) indicate that the dark terrain on Iapetus mustbe quite thin (one to several decimeters); an ammonia-waterice mixture may be present at a depth of several decime-ters below the surface on both the leading and trailing hemi-spheres of Iapetus. The RADAR results appear to rule outany theories of a thick dark material layer (Matthews 1992;Wilson and Sagan 1996) and are consistent with the spectralcoating indicated by VIMS data as described by Clark et al.(2008a). Cassini radio science results (Rappaport et al. 2005)indicate a bulk density for Iapetus of 1:1 g=cm3, from whichit can be inferred that the moon is composed primarily ofwater ice.

Because the large craters within the dark terrain appearevenly colored by the dark material in Voyager data, nocraters significantly break up the dark material to exposebright underlying terrain (Denk et al. 2000, 2009). This sug-gests that the emplacement of dark material by whatevermechanism is relatively new or ongoing.

With the arrival of Cassini at Saturn and the first Iapetusflyby (Dec. 2004), the idea of thermal segregation developed(Spencer et al. 2005; Hendrix and Hansen 2008a; Spencerand Denk 2009). This idea was originally mentioned as oneoption by Mendis and Axford (1974). The September 2007close flyby of Iapetus (Fig. 20.14) permitted testing of thishypothesis, resulting in the following theory. A small amountof exogenic dust from an unknown source created a ‘globalcolor dichotomy’: the leading side became slightly redderand darker than the trailing with fuzzy boundaries close tothe sub-Saturn and anti-Saturn meridians (Denk et al. 2009).A runaway thermal segregation process might have beeninitiated primarily at the low and middle latitudes of theleading side (Spencer and Denk 2009) as well as in somelocal areas at low latitudes on the trailing side, especiallyon equatorward-facing slopes (Denk et al. 2009). Indeed,poleward-facing slopes even within the dark terrain (atmiddle latitudes) are bright. On these slopes the temperaturenever raises enough to allow the thermal process to act fasterthan the destruction process by micrometeorite gardening,which transports bright sub-surface material to the surface.In those locations that are now dark, the originally brightsurface material becomes warm enough for volatiles (mainlywater ice) on the top part of the surface to grow unstableand migrate towards cold traps, leaving behind residual darkmaterial that was originally intimately mixed with the waterice as a minor constituent. The slow 79.3 day rotation is acrucial boundary condition for this process to work in theSaturnian system. ISS images suggest that significant coldtraps might be located at high latitudes on the trailing side,where lower-resolution images show bright white polar caps(Denk et al. 2009).


Compositional data from the UVIS and the VIMS(Hendrix and Hansen, 2009a, b; Clark et al. 2009) indicatevariable amounts of ice and volatiles in the dark material. Thewater-ice absorption edge in the UV at 0:165�m was shownto increase with latitude, consistent with decreasing tem-peratures and increasing numbers of visibly bright regions(Hendrix and Hansen 2009a). This is consistent with ther-mal segregation as a possible process explaining the overallappearance of Iapetus and, more particularly, the observedcorrelations of the dark/bright boundary with geographicallocation and topography (Fig. 20.15) and the small-scale dra-matic brightness variations (Fig. 20.15).

The 3�m absorption observed in VIMS data (Clark et al.2009) might be due to ice, ammonia or bound water. As the3�m absorption is not been well matched by those due to ice,bound water and ammonia, a combination seems to be indi-cated, but a model showing a good match has not yet beendeveloped (Clark et al. 2009). The amount of ice contained

in the dark materials from VIMS data could be on the orderof 10%, but is estimated to be <5% from UVIS data. Thepresence of H2O- or NH3-ice in the dark material is incon-sistent with sublimation rates of �1�m=10 years. Therefore,further spectral observations of Iapetus and laboratory inves-tigations are needed to finally solve the problem.

20.4.3 Surface Alterations and Photometry

Modifications to the surfaces of the icy Saturnian satel-lites include those due to charged-particle bombardment andsputtering, E-ring grain bombardment or coating, thermalprocessing, UV photolysis, and micrometeoroid bombard-ment, many of which can lead to leading-trailing asymme-tries. Photometry is a particularly useful observation methodfor detecting surface alterations, and it was used in the first

Fig. 20.14 Cassini ISS images from the September 2007 flyby, demon-strating the large- and small-scale albedo effects of thermal segregation.(Left) The central longitude of the trailing hemisphere is 24ı to the leftof the mosaic’s center. (Right) The mosaic consists of two image foot-

prints across the surface of Iapetus. The view is centered on terrain near42ı southern latitude and 209:3ı western longitude, on the anti-Saturnfacing hemisphere. Image scale is approximately 32 m/pixel; the craterin the center is about 8 km across


5

4

3

2

10 1 2

Solar phase angle

Brig

htne

ss

3 4

0.08

0.06

0.04

0.02

00.04 0.06

Albedo

Pha

se c

oeffi

cien

t 1–8

deg

(m

ag/d

eq)

0.08 0.1

0.3

0.2

0.1

0.02 0.04 0.06 0.08Albedo

Am

plitu

de

0.1 0.12

Fig. 20.15 (Above) Cassini visual infrared mapping spectrometer(VIMS) disk-integrated observations of the opposition surge on Iape-tus between 0.35 and 3:6�m. (Lower left) The slope of the solar phasecurve (the ‘phase coefficient’) between 1ı and 8ı decreases as thealbedo increases. This effect is expected for shadow hiding, since bright

surfaces have partly illuminated primary shadows. (Lower right) On theother hand, the amplitude of the ‘spike’ under 1ı (shown as a fractionalincrease above the value of the phase curve at 1ı) increases with thealbedo, suggesting that it is a multiple scattering effect, such as coher-ent backscatter

detection of large-scale surface alterations that are the resultof exogenic processes.

All the major Saturnian satellites (with the exception ofHyperion) are synchronously locked, so that one hemispherefaces Saturn at all times. The Saturn-facing hemisphere isapproximately centered on 0 ıW longitude, with the lead-ing hemisphere (facing the direction of orbital motion) cen-tered on 90 ıW and the trailing hemisphere on 270ıW. It isexpected that gardening of the regolith by micrometeoritesis especially important on the leading hemisphere, which‘scoops up’ incoming dust and exogenic material through-out an orbit. In contrast, charged-particle bombardment af-fects primarily (but not exclusively) the trailing hemisphere.Because Saturn’s magnetosphere co-rotates at a rate fasterthan the orbital speed of these moons, the satellites’ trailinghemispheres are preferentially affected by magnetospheric

particle bombardment (cold ions and keV electrons). Me-teoritic bombardment of icy bodies acts in two major waysto alter the optical characteristics of the surface. First, theimpacts excavate and expose fresh material. Second, impactvolatilization and the subsequent escape of volatiles resultsin a lag deposit enriched in opaque, dark materials. The rel-ative importance of the two processes depends on the flux,size distribution and composition of the impacting particlesas well as on the composition, surface temperature and massof the satellite (Buratti et al. 1990). Impact gardening andexcavation tend to brighten an icy surface at visible wave-length (Smith et al. 1981), while impact volatilization tendsto darken it. Buratti (1988) suggested that E-ring coatingdominates any micrometeorite gardening.

Coating by E-ring grains affects the reflectance spec-tra of the satellites and could mask the signatures of other


species. The E-ring is a region of fine-grained .�1�m/ par-ticles primarily consisting of H2O ice that extends from theorbit of Mimas (�3 Saturn radii) outwards past the orbit ofRhea (�15 Saturn radii). The ultimate source of the E-ringis Enceladus and its geyser-like plumes (e.g., Porco et al.2006; Spahn et al. 2006). Compositional measurements ofdust particles in the ring directly reveal the composition oftheir sources. E-ring particles are of two types: (1) H2O iceand (2) organics and/or silicates in icy particles (Postberget al. 2008). E-ring particles coat Rhea and other satellitescloser to Saturn like Mimas, Enceladus, Tethys and Dioneas well as contributing to the G-, F-, and A-rings at least(Buratti et al. 1990; Hamilton and Burns 1994; Buratti et al.1998; Verbiscer et al. 2007). The geometric albedos of theseembedded icy moons were measured in the visible wave-length by the Hubble space telescope (Verbiscer et al. 2007).Their brightness changes with their distance from the E-ring,which is consistent with the assumption that their albedo re-sults from E-ring coating. Hamilton and Burns (1994) pre-dicted that the relative velocities between the E-ring particlesand the icy satellites might explain the large-scale longitu-dinal albedo patterns on the icy satellites. Because of theirhigher relative velocity, the satellites exterior to the dens-est point in the E-ring have brighter leading hemispheres,while those interior to it have brighter trailing hemispheresbecause of the higher velocity of the particles. Indeed, Mimasand Enceladus are both slightly darker on their leading thanon their trailing hemispheres, while the leading hemispheresof Tethys, Dione and Rhea are brighter than their trailinghemispheres (Buratti et al. 1998). The visible orbital phasecurve of Enceladus shows that the leading hemisphere is�1:2 times darker than its trailing hemisphere. At visiblewavelengths, the leading hemisphere of Tethys is �1:1 timesbrighter than its trailing hemisphere, Dione’s leading hemi-sphere is up to �1:8 times brighter and Rhea’s leading hemi-sphere is �1:2 times brighter than the trailing hemisphere(Buratti and Veverka 1984).

The bombardment of icy surfaces by charged particlescauses both structural and chemical changes in the ice(Johnson et al. 2004). Effects include the implantation ofmagnetospheric ions, sputtering and grain size alteration.Low-energy electrons can cause chemical reactions on thesurface, while high-energy electrons and ions tend to causeradiation effects at depth. The latter can result in bulk chem-ical changes in composition as well as structural damage. Aprimary effect of energetic particle bombardment on ice is tocause defects in it (Johnson 1997; Johnson and Quickenden1997; Johnson et al. 1998). Extended defects, in the formof voids and bubbles, affect the light-scattering properties ofthe surface. They also act as reaction and trapping sites forgases, which can produce spectral absorption features. Inci-dent ions deposit energy over a small region around their paththrough the ice, causing local damage. Depending on excita-

tion density and cooling rate, this can result in either localcrystallization or amorphization (Strazzulla 1998).

Energetic ions and electrons have been seen to brightenlaboratory surfaces in the visible by producing voids andbubbles (Sack et al. 1992). Depending on the size and den-sity of these features, a clear surface can become brighter dueto enhanced scattering in the visible. The plasma-inducedbrightening competes with annealing, which operates effi-ciently in regions where the ice temperature exceeds �100Kand can make equatorial regions less light-scattering despitethe plasma bombardment.

A significant effect of the ion bombardment of ice isthe creation of new species through radiolysis. Bombard-ing particles can decompose H2O ice, creating molecular hy-drogen and oxygen. Mini-atmospheres of trapped volatileshave been suggested to form in voids in the ice (Johnsonand Jesser 1997). As a result, species such as O3; H2O2; O2

and oxygen-rich trace species have been detected as gasestrapped in the regoliths of the icy Galilean satellites (Nelsonet al. 1987; Noll et al. 1996; Hendrix et al. 1998, 1999;Spencer et al. 1995; Carlson et al. 1996, 1999; Hibbitts et al.2000, 2003). On these satellites, the hydrogen and oxygenproduced by radiation decomposition form tenuous atmo-spheres; in contrast, atmospheres have not been detectedon satellites in the Saturnian system (with the exception ofthe gaseous plume of Enceladus). There are several possi-ble reasons for this. Lower gravities might support the lossof volatiles; temperatures are lower, which might slow downthe breakup of H2O and the loss of H; the particle environ-ment is different (discussed below); the electron density andtemperature is not effective in exciting and ‘lighting up’ anygases around the moons; and sputtering is not as intense asin the Jupiter system due to differences in the net charge andenergy of the particle environment. On Saturn, the primarysurface alterations found thus far are due to the E-ring, dust(e.g., on Dione) and thermal processing (especially apparenton Iapetus, as previously discussed).

Neutrals, rather than ions as on Jupiter, dominate the Sat-urnian system. The energetic heavy ions (�10 keV to MeV)that are critical at Europa and Ganymede are lost by chargeexchange to the neutral cloud of the Saturnian system as theydiffuse inwards from beyond �10 Saturn radii (Paranicaset al. 2008). However, although the energy fluxes due totrapped plasma and the solar UV are relatively low on theicy Saturnian satellites (Paranicas et al. 2008), radiation pro-cessing appears to be occurring. Radiation fluxes are knownto cause decomposition in ice, possibly producing ozone inthe icy surfaces of Dione and Rhea (Noll et al. 1997) andO2

C in Saturn’s plasma (Martens et al. 2007). In addition,the role of low energy ions has been recently re-examinedusing Cassini CAPS data (Sittler et al. 2008), indicatingthat sputtering by this component is not negligible (John-son et al. 2008). Spencer (1998) searched for but did not


find O2 inclusions on Rhea and Dione. Such inclusions havebeen seen on Ganymede in the low latitudes of the trailinghemisphere, and it was suggested that they might be relatedto charged-particle bombardment or UV photolysis. The lackof O2 on Rhea and Dione suggests that low temperatures pre-clude the formation of O2 bubbles in the ice. Hydrogen per-oxide .H2O2/ has been tentatively VIMS data (Newman et al.2007) of the south pole tiger-stripes region of Enceladus, al-though it is not clear whether exogenic processing is neededto create peroxide.

It is important to distinguish between hemispheric di-chotomies in color, albedo, texture and macroscopic structurebecause each exogenic alteration mechanism leaves a differ-ent signature on the physical and optical properties of thesurface. Photometric methods have been employed over thepast two decades to characterize the roughness, compactionstate, particle size and phase function, and single scatteringalbedo of planetary and satellite surfaces, including the Sat-urnian satellites (Buratti 1985; Verbiscer and Veverka 1989,1992, 1994; Domingue et al. 1995). An indication of the scat-tering properties of the satellites’ surfaces can be obtainedby extracting photometric scans across their disks at smallsolar phase angles. Dark surfaces such as that of the Moon,in which single scattering is dominant, are not significantlylimb-darkened, whereas bright surfaces in which multiplescattering is important could be expected to show limb dark-ening. Physical properties such as roughness, compactionstate and particle size give clues as to the past evolution ofthe satellites’ surfaces and reveal the relative importance ofexogenic alteration processes. Accurate photometric model-ing is also necessary to define the intrinsic changes in albedoand color on a planetary surface (Noland et al. 1974); much,and in many cases most of the change in intensity is due tovariations in the radiance viewing geometry (incident, emis-sion, and solar phase angle) and is not intrinsic to the surface.

Planetary surface scattering models have been developed(e.g., Horak 1950; Goguen 1981; Lumme and Bowell 1981a,b; Hapke 1981, 1984, 1986, 1990; Shkuratov et al. 1999,2005) which express the radiation reflected from the surfacein terms of the following physical parameters: single scat-tering albedo, single particle phase function, the compactionproperties of the optically active portion of the regolith,and the scale and extent of macroscopically rough surfacefeatures. Observations at small solar phase angles .0–10ı/are particularly important for understanding the compactionstate of the upper regolith: this is the region of the well-known opposition effect, in which the rapid disappearance ofshadowing among surficial particles causes a nonlinear surgein brightness as the object becomes fully illuminated. At thesmallest phase angles .<1ı/, coherent backscatter dominatesthe surge (Hapke 1990). Larger phase angles .>40ı/ are nec-essary for investigating the macroscopically rough nature ofsurfaces, ranging in size from clumps of several particles

to mountains, craters and ridges. Such features cast shad-ows and alter the local incidence and emission angles. Twoformalisms have been developed to describe roughness forplanetary surfaces: a mean-slope model in which the surfaceis covered by a Gaussian distribution of slope angles witha mean angle of ™ (Hapke 1984; Helfenstein and Shepard1998), and a crater model in which the surface is coveredby craters with a defined depth-to-radius ratio (Buratti andVeverka 1985). The single particle phase function, which ex-presses the directional scattering properties of a planetarysurface, is an indicator of the physical character of individ-ual particles in the upper regolith, including their size, sizedistribution, shape, and optical constants. Small or transpar-ent particles tend to be more isotropically scattering becausephotons survive to be multiply scattered; forward scatter-ing can exist if photons exit in the direction away from theobserver. The single particle phase function is usually de-scribed by a 1 or 2-term Henyey-Greenstein phase functiondefined by the asymmetry parameter g, where g D 1 is purelyforward-scattering, g D �1 is purely backscattering, andg D 0 is isotropic. The opposition surge of airless surfaceshas been described by parameters that define the compactionstate of the surface and the degree of coherent backscatter ofmultiply scattered photons (Irvine 1966; Hapke 1986, 1990).One final physical photometric parameter is the single scat-tering albedo . K /, which is the probability that a photon willbe scattered into 4 steradians after one scattering.

Three other important physical parameters are the ge-ometrical albedo (p), the phase integral (q) and the Bondalbedo .AB/. The geometrical albedo is the flux receivedfrom a planet at a solar phase angle of 0ı compared withthat of a perfectly diffusing disk of the same size, also at 0ı.The phase integral (q) is the flux, normalized to unity at asolar phase angle of 0ı, integrated over all solar phase an-gles. The Bond albedo is p q, and when integrated over allwavelengths it yields the bolometric Bond albedo, which is ameasure of the energy balance on the planet or satellite. Ex-cept for the bolometric Bond albedo, all these parameters arewavelength-dependent.

Buratti (1984) and Buratti and Veverka (1984) performedphotometric analyses of visible-wavelength Voyager data ofthe icy Saturnian moons based on Voyager disk-resolved datasets, and disk-integrated Voyager observations at large phaseangles to supplement existing ground-based observations atsmaller phase angles. These studies investigated limb darken-ing with implications for multiple scattering (significant limbdarkening suggests that multiple scattering is important).Physical photometric modeling of Saturnian moons based onVoyager and ground-based data was done by Buratti (1985),Verbiscer and Veverka (1989, 1992, 1994), Domingue et al.(1995) and Simonelli et al. (1999). In general, the surfaces ofthe Saturnian satellites are backscattering and exhibit rough-ness comparable to that of the Moon, indicating that impact


processes produce similar morphological effects on rockyand icy bodies at low temperatures. The photometric andphysical properties of the satellites determined so far fromboth Voyager and Cassini data are summarized in Table 20.4.

Opposition surge data from 0.2 to 5:1 �m were obtainedfor Enceladus, Tethys, Dione, Rhea and Iapetus. Figure 20.16shows the opposition curve of Iapetus in between 0:35 �mand 3:6 �m (Buratti et al. 2009b; Hendrix and Buratti 2009).The curve is divided into two components: at phase anglesgreater than 1ı, brightness increases linearly with decreas-ing phase angle, while at phase angles smaller than onedegree the increase is exponential. Moreover, the albedo-dependence of the two types of surge is different. For the

first type, the magnitude of the surge decreases as the albedoincreases, while for the second type of surge the reverse istrue. This result suggests that different phenomena are re-sponsible for the two types of surge. At angles greater than1ı, the surge is caused by mutual shadowing, in which shad-ows cast among particles of the regolith rapidly disappearas the face of the satellite becomes fully illuminated to theobserver. Since multiply scattered photons, which partly il-luminate primary shadows, become more numerous as thealbedo increases, this effect becomes less pronounced forbright surfaces. On the other hand, the huge increase atvery small solar phase angles can be explained by coher-ent backscatter, a phenomenon in which photons following

Table 20.4 Photometric and physical parameters of the icy Satur-nian satellites (references: (1) Buratti and Veverka 1984; (2) Buratti1985; (3) Verbiscer and Veverka 1994; (4) Verbiscer and Veverka 1989;

(5) Verbiscer and Veverka 1992; (6) Simonelli et al. 1999; (7) Burattiet al. 2008; (8) Verbiscer et al. 2005)

Parameter Mimas Enceladus Tethys Dione Rhea Phoebe

Mean slopeangle � .ı/

30 (2) 0˙ 1 (5) 6˙ 1 (3) 13˙ 5 (4) 31˙ 4 (6) 3˙ 3 (7)

Asymmetryparameter g

�0:3˙ 0:05 (2)0:213 ˙ 0:005

(5)

�0:35 ˙ 0:03 (2)0:399 ˙ 0:005

(3)

�0:287 ˙ 0:007

(4)�0:24.6/ � 0:3˙ 0:1

(plus a forwardcomponent) (7)

Single scatteringalbedo K�

0:93˙ 0:03 (2)0:951 ˙ 0:002

(5)

0:99 ˙ 0:02 (2)0:998 ˙ 0:001

(3)

0:861˙ 0:008 (4) 0.068 (6) 0:07˙ 0:01

(7)

Phase integral q 0:82˙ 0:05 (1)0:78˙ 0:05 (5)

0:86 ˙ 0:1 (1)0:92˙ 0:05 (3)

0:7˙ 0:1 (1) 0:8˙ 0:1 (1) 0:70˙ 0:06 (1)0:78˙ 0:05 (4)

0:24 ˙ 0:05 (6)0:29˙ 0:03 (7)

Visual geometricalbedo pv

0:77˙ 0:15 (1) 1:04 ˙ 0:15 (1)1:41˙ 0:03 (8);

0:80˙ 0:15 (1) 0:55˙ 0:15 (1) 0.65 (1) 0:081 ˙ 0:002 (6);0:082 ˙ 0:002 (7);

Bond albedo AB 0:6˙ 0:1 (1)0:5˙ 0:1 (5)

0:9˙ 0:1 (1)0:91˙ 0:1 (3)

0:6˙ 0:1 (1) 0:45˙ 0:1 (1) 0:45˙ 0:1 (1)0:49˙ 0:06 (4)

0:020 ˙ 0:004 (6)0:023 ˙ 0:007 (7)

0.02

0.015

0.01

0.005

00 20 40

Emission Angle [deg]

Ref

lect

ivity

q=0.140f(a)=0.023

I/F OF LOW ALBEDO SIDEI/F OF MODEL

8060

Fig. 20.16 A macroscopic roughness model for the low-albedohemisphere of Iapetus, based on the crater roughness model by Burattiand Veverka (1985). The best-fit roughness function corresponds to adepth-to-radius of 0.14, or a mean slope angle of 11ı, which is much

lower than the �30ı typical of other icy satellites (see Table 20.4). Forthese fits, a surface phase function f .˛/ of 0.023 was derived. FromLee et al. 2009


identical but reversed paths in a surface interfere construc-tively in exactly the backscattering direction, causing bright-ness to increase by up to a factor of two (Hapke 1990).

The Cassini flyby in June 2004 prior to Saturn orbit in-sertion afforded views at large phase angles – importantfor modeling roughness - and over a full excursion in geo-graphical longitudes. Phoebe exhibits a substantial forward-scattering component to its single particle phase function.This result is consistent with a substantial amount of fine-grained dust on the surface of the satellite generated by par-ticle infall, as suggested by Clark et al. (2005, 2008a). Asobserved in Cassini images (Porco et al. 2005a), Phoebe alsoshows extensive macroscopic roughness, hinting at the vio-lent collisional history predicted by Nesvorny et al. (2003).

ISS images of the low-albedo hemisphere of Iapetus werefitted to the crater roughness model (Buratti and Veverka1985), yielding a markedly smooth value for the degree ofmacroscopic roughness on this side with a mean slope angleof only a few degrees (Lee et al. 2009). This result suggeststhat small-scale rough features on the dark side have beenfilled in (features which probably dominate the photometriceffects; see Helfenstein and Shepard 1998), which is consis-tent with the idea of Spencer and Denk (2009) that thermalmigration of water ice leaves behind a fluffy residue of darkmaterial. The low-albedo side of Iapetus shows a roughnessmuch smaller than that of the other satellits and compara-ble within a factor of 2 with Enceladus, which is coated withmicron-sized particles from its plume and the E-ring (Verbis-cer and Veverka 1994; see Table 20.4). Model-fits for the lowalbedo hemisphere are shown in Fig. 20.16.

Finally, observations at very large solar phase angles.155ı–165ı/ have been used to search for cryovolcanic ac-tivity on satellites. By comparing the brightness of Rhea toEnceladus in this phase angle range at 2:02 �m, Pitman et al.(2008) were able to place an upper limit on water vapor col-umn density based on a possible plume of 1:52 � 1014 to1:91 � 1015=cm2, two orders of magnitude below the ob-served plume density of Enceladus (Fig. 20.17).

The microphysical structure of the E-ring grain coatingsof the inner icy moons embedded in the E-ring can be probedby investigating the opposition surge. Verbiscer et al. (2007)found that the amplitude and angular width of the oppositioneffect on the inner icy moons are correlated with the moons’position relative to the E-ring.

20.5 Constraints on the Top-Meter Structureand Composition by Radar

The wavelength of Cassini’s 13.8 GHz RADAR instrument,2.2 cm, is about six times shorter than the only ground-based radar wavelength available to study the satellites

Fig. 20.17 Cassini VIMS disk-integrated brightness as a function ofsolar phase angle at 2:23�. Symbols: Normalized Rhea and EnceladusVIMS brightness (every fifth data point plotted). Solid line: Third-orderpolynomial fit to Rhea data. Dashed lines: (polynomial fit to Rhea data).Enceladus’s plume can be seen as a peak at a solar phase angle of 159ı.Adapted from Pitman et al. (2008)

(13 cm, at Arecibo) and 22 times longer than the millimeterwavelengths at the limit of the CIRS.

The echoes result from volume scattering, and theirstrength is sensitive to ice purity. Therefore, they provideunique information about near-surface structural complexityas well as about contamination with non-ice material. Thissection summarizes the most basic aspects of Cassini- andground-based radar observations of the satellites, the theo-retical context for interpreting the echoes, and the inferencesdrawn about subsurface structure and composition.

Cassini measures echoes in the same linear polarizationas transmitted, whereas most ground-based radar astronomyconsists of Arecibo 13 cm or 70 cm observations (Black et al.2001a) and Goldstone 3.5 cm observations that almost al-ways use transmission of a circularly polarized signal andsimultaneous reception of echoes in same-circular (SC) andopposite-circular (OC) polarizations. A radar target’s radaralbedo is defined as its radar cross section divided by its pro-jected area. The radar cross section is the projected area ofa perfectly reflective isotropic scatterer which, if observed atthe target’s distance from the radar and using the same trans-mitted and received polarizations, would return the observedecho power. The total-power radar albedo is the sum of thealbedos in two orthogonal polarizations, TP D SL C OL DSC C OC. Here, we will use ‘SL-2’ to denote the 2 cm radaralbedo in the SL polarization, ‘TP-13’ to denote the 13 cmtotal-power albedo, etc.

For a smooth, homogeneous target, single back reflectionswould be entirely OC when the transmission is circularlypolarized, or entirely SL when the transmission is linearlypolarized, so the polarization ratios SC/OC and OL/SLwould equal zero. Saturn’s icy satellites and the icy Galilean


satellites not only have radar albedos an order of magnitudegreater than those of the Moon and most other solar systemtargets but also unusually large polarization ratios.

These strange radar signatures were discovered in obser-vations of Europa, Ganymede and Callisto in the mid-1970s.It took a decade and a half to realize that the signatures arethe outcome of ‘coherent backscattering,’ which is phase-coherent, multiple scattering within a dielectric medium thatis both heterogeneous and non-absorbing (e.g., MacKintoshand John 1989; Hapke 1990).

During the next decade, Peters (1992) produced a vectorformulation of coherent backscattering that led to predictionsof radar albedos and polarization ratios, and Black (1997)and Black et al. (2001b) built on Peters’ work with modelsof scatterer properties.

Extremely clean ice is insufficient for anomalouslylarge radar albedos and circular polarization ratios, whicharise from multiple scattering within a random, disorderedmedium whose intrinsic microwave absorption is very low.Thus, perfectly clean ice that is structurally homogeneous(constant density at millimeter scales) would not give anoma-lous echoes. The regoliths on solar system objects are struc-turally very heterogeneous, with complex density varia-tions from the distribution of particle sizes and shapes pro-duced by impacts of projectiles of different sizes and impactenergies.

It is to be expected that the uppermost few meters of Sat-urn’s icy satellite regolith are basically similar to those ofthe icy Galilean satellites or the Moon’s in terms of par-ticle size distribution, porosity and density as functions ofdepth, and lateral/vertical heterogeneity. It is the combina-tion of naturally complex density variations and the extremetransparency of pure water ice that gives icy regolith its ex-otic radar properties.

The icy Galilean satellites were observed in dual-circular-polarization experiments from the ground at 3.5, 13, and70 cm; experiments measuring SL-13 and OL-13 were doneonly for Europa, Ganymede and Callisto (Ostro et al. 1992).Arecibo obtained estimates of SC-13 and OC-13 for Ence-ladus, Tethys, Dione, Rhea and Iapetus. The Cassini radarmeasured SL-2 for Mimas, Hyperion and Phoebe. The satel-lites’ albedos show different styles and are location- andwavelength-dependent, as discussed below. The results de-scribed in this chapter make use of 73 Cassini SL-2 measure-ments, more than twice the number reported by Ostro et al.(2006).

If a radar target is illuminated by a uniform antenna beam(that is, if the antenna gain is constant across the disk), thenit is a simple matter to use the appropriate ‘radar equation’to convert the measured echo power spectrum (power as afunction of Doppler frequency) into a radar cross section(e.g., Ostro 1993). For Cassini measurements, except forthe SAR imaging of Iapetus discussed below, the antenna

beam width is within a factor of several times the target’sangular width, so the gain varies across the target’s disk,usually - significantly. Data reduction yields an estimate ofecho power, which is the sum of contributions from differentparts of the surface, with each contribution weighted by thetwo-way gain. However, the distribution of echo power as afunction of location on the target disk is not known a pri-ori. Thus, we do not know how much the echo power fromany given surface element has been ‘amplified’ by the an-tenna gain. To estimate a radar albedo from our data we mustmake an assumption about the homogeneity and angular de-pendence of the surface’s radar scattering. For the purposeof obtaining a radar albedo estimate from an echo powerspectrum, one assumes uniform, azimuthally isotropic, co-sine scattering laws, and uses least squares to estimate thevalue of a single parameter that defines the echo’s incidence-angle dependence. In general, only modest departures fromLambert limb darkening are found.

20.5.1 Radar-Optical Correlations

Figure 20.18 shows all measurements of the satellites’ in-dividual SL-2 albedos along with their average optical ge-ometrical albedos, plotted vs. distance from Saturn. A verypronounced radar-optical correlation means that variations inoptical and radar albedos involve the same or related con-taminant(s) and/or surface processes. Absolutely pure waterice at satellite surface temperatures is extremely transparentto radar waves. Contamination of water ice with no morethan trace concentrations of virtually any other substance (in-cluding earth rocks, chondritic meteorites, lunar soil, ammo-nia, metallic iron, iron oxides and tholins) can decrease itsradar transparency (and hence the regolith’s radar albedo)by one to two orders of magnitude. Phenomena that dependon the distance from Saturn include primordial composi-tion, meteoroid flux and the influx of dark material (per-haps from several sources), the balance between the influx ofammonia-containing particles and the removal of ammoniaby a combination of ion erosion and micrometeoroid garden-ing, variation in E-ring fluxes, and the systematics of thermalvolatiles redistribution.

Modeling of icy Galilean satellite echoes (Black et al.2001b) suggests that penetration to about half a meter shouldbe adequate to produce the range of measured values of SL-2.

The explanation by Clark et al. (2008b, 2009) for darkSaturnian system material, which suggests that it consistsof sub-micron iron particles from meteorites which oxidizedand became mixed with ice, would also suffice to explain,at least qualitatively, the overall radar-optical albedo corre-lation in Fig. 20.18 if the contamination of the ice extendssome one to several decimeters below the optically visible


Radar SL-2 Albedos andOptical Geometric Albedos

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Fig. 20.18 Individual Cassini measurements of SL-2 (small symbols)and average optical geometrical albedos, plotted for the satellites vs.their distances from Saturn. Optical albedos of Mimas, Enceladus,Tethys, Dione and Rhea are from Verbiscer et al. (2007); those of Ia-petus and Phoebe are from Morrison et al. (1986) and involve a V fil-ter; Hyperion’s is from Cruikshank and Brown (1982). Note the similardistance dependences. Hyperion’s radar albedo is large and appears tobreak the pattern, perhaps (Ostro et al. 2006) because Hyperion’s opti-cal coloring may be due to a thin layer of exogenously derived materialthat, in contrast to the other satellites, has never been worked into theicy regolith, which remains very clean and radar-bright

surface, as expected by Clark et al. (2008b, 2009). The ten-tative detection of ammonia in VIMS spectra by Clark et al.(2008a) opens up the possibility that contamination by thismaterial may play an important role in inter- and intra-satellite radar albedo variations. If so, this contaminationwould affect the radar but not the optical albedo.

Almost all Cassini radar data on the icy satellites consistof real-aperture scatterometry, which means that the beamsize determines the resolution. Cassini has no real-apertureresolution of Mimas, Enceladus, Hyperion or Phoebe, butdoes have abundant measurements for Tethys, Dione, Rheaand Iapetus with beams smaller than, and in many cases lessthan half, the size of the disk (Ostro et al. 2006).

The SL-2 distributions of the individual satellites can besummarized briefly as follows. As regards Mimas, Herschelmay or may not contribute to the highest value of thedisc-integrated albedo. Enceladus shows no dramatic lead-ing/trailing asymmetry and no noteworthy geographical

variations except for a low albedo in an observation of theSaturn-facing side. Tethys’ leading side is brighter both op-tically and in SL-2. Dione’s highest-absolute-latitude viewsgive the highest radar albedos. Rhea’s leading side has alower SL-2 in the north, where there is much cratering, andhigher in the south, where there are optically bright rays. Itstrailing and Saturn-facing sides, which look younger, show arelatively high SL-2. Iapetus displays an extremely markedcorrelation between optical brightness and SL-2.

In general, a first-order understanding of the SL-2 distri-bution is that higher SL-2 means cleaner near-surface ice, butthe nature of the contaminant(s) and the depth of the contam-ination are open questions. The latter can be constrained bymeasuring radar albedos at more than one wavelength. For-tunately, there are Arecibo 13-cm observations of Enceladus,Tethys, Dione, Rhea (Black et al. 2004) and Iapetus (Blacket al. 2007).

Also fortunately, the wavelength dependence observed forEuropa, Ganymede and Callisto (EGC) gives us a context fordiscussing the wavelength dependence observed for Saturn’ssatellites.

20.5.2 Wavelength Dependence

EGC total-power radar albedos and circular polarization ra-tios at 3.5 cm and 13 cm are indistinguishable. This meansthat their regolith structure (distribution of density variations)and ice cleanliness look identical at the two wavelengths, foreach object. That is, near-surface structures are ‘self-similar’at 3.5 cm and 13 cm scales; there are no depth-dependentphenomena, like contamination levels or structural character,perceptible in EGC’s 3.5 cm and 13 cm radar signatures.

At 70 cm, moreover, EGC SC/OC ratios show little wave-length dependence, so subsurface multiple scattering remainsthe primary source of the echo even on this longer scale.However, the radar albedos (TP-70) plummet for all threeobjects. The implication is that at 70 cm, scattering hetero-geneities are relatively sparse, and the self-similarity of theregolith structure disappears on the larger scale. On Europaat least, with the most striking albedo drop from TP-13 toTP-70, the regolith may simply be too young for it to be het-erogeneous at 70 cm.

Within the framework of the Black et al. (2001b) model,this implies that at 70 cm ‘the scattering layer is disappear-ing in the sense that there are too few scatterer at this scaleand they are confined to a layer too shallow to permit nu-merous multiple scatterings. At even longer wavelengths, ofthe order of several meters, the reflectivity from the sur-faces of these moons, and Europa in particular, may drop


precipitously or even vanish as the layers responsible for thecentimeter-wavelength radar properties become essentiallytransparent.’

Now, given that the 3.5 cm and 13 cm polarization proper-ties of EGC show no significant wavelength dependence, andthat the 13 cm OL/SL ratios of those objects are very similar,let us assume that the linear polarization properties of theSaturnian satellites are wavelength-invariant. Let us then useEGC’s mean OL/SL value to estimate total-power 2 cm albe-dos: TP-2 D .1:52 ˙ 0:13/ SL-2 and compare them to theArecibo 13 cm total-power albedos (TP-13).

Figure. 20.19 shows relevant Cassini and Arecibo infor-mation.

Moving outward from Saturn, the magnitude of 2 cm to13 cm wavelength dependence is unremarkable for Ence-ladus’ leading side, extreme for Enceladus’ trailing side, verylarge for Tethys, large for Dione, unremarkable for Rhea,large for Iapetus’ leading side, and extreme for Iapetus’ trail-ing side.

Whereas compositional variation (ice cleanliness) almostcertainly does not contribute to the wavelength dependenceseen in EGC, it may well play a key role in the wavelengthdependence seen in Saturn’s satellites. In this complicatedsystem, the tapestry of satellite radar properties may involvemany factors, including inter- and intra-satellite variationsin E-ring flux, variations in the meteoroid or ionizing flux,variations in the concentration of radar absorbing materials

3

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ENCTETDIORHEIAPENC_L_13cmENC_T_13cmTET_13cmDIO_13cmRHE_13cmIAP_L_13cmIAP_T_13cmEuropaGanymedeCallisto

Fig. 20.19 Wavelength dependence of total-power radar albedos. In-dividual symbols are estimates of TP-2 obtained from individual mea-surements of SL-2 using TP-2 D 1.52 SL-2 ˙ 0.13 SL. Arecibo 13-cmalbedos are horizontal lines. Dashed lines represent EGC

(including optically dark material and possibly ammonia) asa function of depth, the systematics of thermal volatile re-distribution (especially for Iapetus), and/or constituents withwavelength- or temperature-dependent electrical properties.Here, we will present what we consider reasonable hypothe-ses about what the most important factors are.

The unremarkable 2- to-13-cm wavelength dependenceseen on Enceladus’ leading side and also on Rhea is reminis-cent of Europa’s 3.5- to-13-cm wavelength invariance. Thesimplest interpretation is that at least the top few decime-ters are uniformly clean and old enough for meteoroid bom-bardment to have created density heterogeneities that areable to produce coherent backscattering and extreme radarbrightness.

Enceladus’ leading-side 2 cm and 13 cm TP albedos re-semble Europa’s 3.5 cm and 13 cm TP albedos, whereasRhea’s 2 cm and 13 cm TP albedos resemble Ganymede’s3.5 cm and 13 cm TP albedos, reflecting much greater con-tamination of Rhea’s upper regolith.

Dione’s and Rhea’s TP-2 albedos are similar, but Dioneis dimmer at 13 cm. Dione VIMS data show evidence of abombardment of apparently dark particles on the trailing side(Clark et al. 2008a). If the dark material contains ammonia,and the ejecta systematics of meteoroid bombardment effi-ciently distribute the dark material over the surface of theobject, multiple scattering will be cut off at a depth that maybe too shallow for TP-13 to be enhanced.

The striking wavelength dependence of Enceladus’ trail-ing side .TP-2 >> TP-13/ mirrors Europa’s .TP-13 >>

TP–70/. It is easily understood if we assume that the topmeter or so of ice was deposited so recently that there hasnot been time for the regolith to be matured by meteoroidbombardment. That is, there are small-scale heterogeneitiesthat produce coherent backscattering at 2 cm, but none ofthe larger-scale heterogeneities needed to produce coherentbackscattering at 13 cm.

This idea is consistent with the fact that Enceladus’ trail-ing side is optically brighter than its leading side, apparentlybecause of the greater E-ring flux it receives (Buratti 1988;Showalter et al. 1991; Buratti et al. 1998). So all of Ence-ladus is very clean, but the uppermost layer of one or severalmeters on the trailing side is also extremely young, resultingin a striking hemispheric dichotomy at 13 cm that is the mostintense seen at that wavelength in the icy-satellite systems ofeither Jupiter or Saturn.

At least some of the decrease in average SL-2 (Fig. 20.18)and the decrease in average wavelength dependence outwardfrom Enceladus (Fig. 20.19) is probably due to the outwarddecrease in E-ring flux (Verbiscer et al. 2007).

Finally, we get to Iapetus, where TP-2 is strongly cor-related with the optical albedo: low for the optically darkleading-side material and high for the optically brighttrailing-side material. However, Iapetus’ TP-13 values show


much less of a dichotomy and are several times lower thanTP-2, being indistinguishable from the weighted mean of TP-13 for main-belt asteroids, 0:15˙ 0:10 (Table 20.5).

Therefore, whereas Iapetus is anomalously bright at 2 cm(similar to Callisto) and mimics the optical hemispheric di-chotomy, at 13 cm it looks like a typical main-belt asteroid,with minimal evidence of mimicking the optical hemisphericdichotomy.

The Iapetus results are understandable if (1) the leadingside’s optically dark contaminant is present to depths of atleast one to several decimeters, so that SL-2 can see the op-tical dichotomy, and (2) ammonia and/or some other radar-absorbing contaminant is globally much less abundant withinthe upper one to several decimeters than at greater depths,which would explain the wavelength dependence, that is,the virtual disappearance of the object’s 2-cm signature at13 cm. The first conclusion is consistent with the dark ma-terial being a shallow phenomenon, as predicted from thethermal migration hypothesis discussed elsewhere in thischapter.

Moreover, ammonia contamination may be uniformwithin the top few meters, so that the wavelength depen-dence arises because ammonia’s electrical loss is muchgreater at 13 cm (2.38 GHz) than at 2 cm (13.8 GHz). Un-fortunately, the likelihood of this possibility cannot yet bejudged because measuring the complex dielectric constant ofammonia-contaminated water ice as a function of concentra-tion, frequency and temperature is so difficult that, althoughneeded, it has not been done yet. This is a pity, becausethe lack of this information undermines the interpretation

Table 20.5 Total-power radar albedosDark/Leading Bright/Trailing

2:2-cm 0:32˙ 0:09 0:58˙ 0:18

13-cm 0:13˙ 0:04 0:17˙ 0:04

of the 2 cm to 13 cm wavelength dependence seen in otherSaturnian satellites as well as of the possible influence of am-monia and its ‘modulation’ by magnetospheric bombardment(which varies with distance from Saturn) on the dependenceof SL-2 on the distance from Saturn.

Verbiscer et al. (2006) suggested that, while their spec-tral data do not contain an unambiguous detection of ammo-nia hydrate on Enceladus, their spectral models do not ruleout the presence of a modest amount of it on both hemi-spheres. However, they note that the trailing hemisphere isalso exposed to increased magnetospheric particle bombard-ment, which would preferentially destroy the ammonia. Inany case, ammonia does not appear to be a factor in Ence-ladus’ radar albedos.

20.5.3 Iapetus Radar Image

During the Iapetus 49 flyby, RADAR obtained an SAR im-age at 2 km to 12 km surface resolution that covers much ofthe object’s leading (optically darker) side. From a compar-ison with optical imaging (Fig. 20.20), we see that, as ex-pected, the SAR reliably reveals the surface features seen inISS images.

The pattern of radar brightness contrast is basically sim-ilar to that seen optically, except for the large basin at30 ıN; 75 ıW, which has more SL-2-bright than opticallybright highlights. This suggests that the subsurface of thesehighlights at decimeter scales may be cleaner than in otherparts of the optically dark terrain, which is consistent with theidea that the contamination by dark material might be shal-low, at least in this basin. Conversely, many optically brightcrater rims are SL-2-dark, so the exposure of clean ice cannotbe very deep.

Fig. 20.20 RADAR SAR image of the leading side of Iapetus (Sept. 2007, left) and the ISS mosaic PIA 08406 (Jan. 2008), using similarprojections and similar scales (the large crater on the right in the ISS image is about 550 km in diameter)


20.6 Geological Evolution

The densities of the Saturnian satellites constrain theircomposition to be about 2/3 ice and 1/3 rock. Accretionalenergy, tidal despinning and radioactive decay might haveheated and melted the moons. Due to the very low melt-ing temperatures of various water clathrates, it is likely thatthe larger bodies differentiated early in their evolution whilesmaller satellites may still be composed and structured intheir primitive state. For a more detailed discussion of theaspects of origin and thermal evolution, see Johnson et al.(2009) and Matson et al. (2009).

The heavily cratered solid surfaces in the Saturnian sys-tem indicate the role of accretional and post-accretional bom-bardement in the geological evolution of the satellites. Ob-served crater densities are summarized in Dones et al. (2009).However, the source as well as the flux of bombarding bodiesis still under discussion, and Dones et al. (2009) came to theconclusion that, since no radiometric sample data exist fromsurfaces in the outer solar system and the size-frequency dis-tribution of comets is less well understood than that of as-teroids, the chronology of the moons of the giant planetsremains in a primitive state. Therefore, we refer the reader toDones et al. (2009) for the crater chronology discussion. Ininterpreting the geological evolution of the Saturnian satel-lites, we will rely on relative stratigraphic relations. Never-theless, absolute ages are given in Dones et al. (2009).

In a stratigraphic sense, three major surface unit cate-gories can be distinguished on all icy satellites:

(1) Cratered plains, which are characterized by dense craterpopulations that are superimposed by all other geologicalfeatures and, thus, are the relatively oldest units.

(2) Tectonized zones of narrow bands and mostly sub-parallel fractures, faults, troughs and ridges of global ex-tension indicating impact-induced or endogenic crustalstress. These units dissect cratered plains and partly dis-sect each other, indicating different relative ages.

(3) Resurfaced units consisting of terrain that has undergonesecondary exogenic and endogenic processes such asmass wasting, surface alterations (due to micrometeoritegardening, sputtering, thermal segregation and crater rayformation by recent impacts), contamination by dark ma-terial and cryovolcanic deposition mostly resulting insmoothed terrain. Superimposed on cratered plains andtectonized regions, these are the youngest surface units.

Heavily cratered surfaces are found on all the icy rocksand small satellites (Fig. 20.21) as well as on the medium-sized moons Mimas, Enceladus, Tethys, Dione, Rhea, Titan,Hyperion, Iapetus and Phoebe, indicating an intense post-accretional collision process relatively early in the geologi-cal history of the satellites. Cassini observations added to thenumber of large basins and ring structures mainly on Rhea,

Fig. 20.21 Phoebe, 213 km in average diameter, was the first one of thenine classic moons of Saturn to be captured by the ISS cameras aboardCassini in a close flyby about three weeks prior to Cassini’s Saturn OrbitInsertion (SOI) in July 2004. The mosaic shows a densely cratered sur-face. Higher-resolution data revealed that the high frequency of cratersextends with a steep slope down to craters only tens of meters in diame-ter (Porco et al. 2005a). Image source: http://photojournal.jpl.nasa.gov,image identification PIA06064

Dione and Iapetus. Differences in the topographical expres-sion of these impact structures suggest post-impact processessuch as relaxation, which may reflect differences in the ther-mal history and/or crustal thickness of individual satellites(e.g., Schenk and Moore 2007; Giese et al. 2008). Althoughage models are based on different impactor populations, thecratered surfaces on Iapetus seem to be the oldest in the Sat-urnian system, whereas cratered plains on the other satellitesseem to be younger (see Dones et al. 2009). One large (48 kmin diameter) stratigraphically young ray crater was found onRhea (Wagner et al. 2007) (Fig. 20.22).

While the small satellites lack tectonic features, themedium-sized bodies experienced rotation-, tide- andimpact-induced deformation and volume changes, which cre-ated extensional or contractional tectonic landforms. Parallellineaments on Mimas seem to be related to the Herschel im-pact event (McKinnon 1985; Schenk 1989a), although freez-ing expansion cannot be ruled out as the cause (Fig. 20.23).

http://photojournal.jpl.nasa.gov


Fig. 20.22 A prominent bright ray crater with a diameter of 48 km islocated approximately at 12:5 ıS; 112 ıW on Saturn’s second-largestsatellite Rhea. The crater is stratigraphically young, indicated by thelow superimposed crater frequency on its floor and continuous ejecta

(Wagner et al. 2008). The high-resolution mosaic (right) of NAC im-ages, shown in context of a WAC frame, was obtained by Cassini ISSduring orbit 049 on Aug. 30, 2007. The global view to the left showingthe location of the detailed view is a single NAC image from orbit 056

Fig. 20.23 The anti-Saturnian hemisphere of Mimas, centered approx-imately at lat. 22 ıS, long. 185 ıW, is shown in a single image (NACframe N1501638509) at a resolution of 620 m/pixel, taken during a non-targeted encounter in August 2005. The densely cratered landscape im-plies a high surface age. Grooves and lineaments infer weak tectonic onthis geologically unevolved body

Enceladus exhibits the greatest geological diversity in acomplex stress and strain history that is predominantly tec-tonic and cryovolcanic in origin, with both processes strati-graphically correlated in space and time, although the energysource is so far not completely understood (e.g., Smith et al.1982; Kargel and Pozio 1996; Porco et al. 2006; Nimmo andPappalardo 2006; Nimmo et al. 2007; Jaumann et al. 2008).Enceladus is discussed in detail in Spencer et al. (2009).

Located in densely cratered terrain, the nearly globe-encircling Ithaca Chasma rift system on Tethys (Fig. 20.2) is100 km wide, consists of at least two narrower branches to-wards the south and was caused either by expansion when theliquid-water interior froze (Smith et al. 1981) or by deforma-tion by the large impact which formed Odysseus (Moore andAhern 1983). However, crater densities derived from Cassiniimages suggest Ithaca Chasma to be older than Odysseus(Giese et al. 2007), so that the rift system may be an endo-genic feature.

Dione exhibits a nearly global network of tectonic fea-tures such as troughs, scarps, ridges and lineaments partlysuperimposed on each other (Fig. 20.24), which correlatewith smooth plains (e.g., Plescia 1983; Moore 1984; Stephanet al. 2008). The global orientation of Dione’s tectonic pat-tern is non-random, suggesting a complex endogenic his-tory of extension followed by compression. Cassini imagesindicate that some of the grabens are fresh, which impliesthat tectonic activity may have persisted into geologicallyrecent times (Wagner et al. 2006). Cassini also revealed acrater-lineament relationship on Dione (Moore et al. 2004;Schenk and Moore 2007), indicating exogenic-induced tec-tonic activity in the early geological history of the moon.Thus, Dione’s crust not only experienced stress caused bydifferent exogenic and endogenic processes, such as ancientimpacts, despinning and volume change over a long period oftime, but also more recent tidal stress. As on Dione, wispy ac-curate albedo markings on Rhea are associated with an exten-sional fault system (Schenk and Moore 2007; Wagner et al.2007), although Rhea seems less endogenically evolved thanthe other moons (Fig. 20.25).

Parallel north-south trending lineaments are thought tobe the result of extensional stress followed by compression


Fig. 20.24 This image of Tethys, taken by the Cassini ISS NA camerain May 2008 in visible light, shows the old, densely cratered surfaceof Tethys and a part of the extensive graben system of Ithaca Chasma.The graben is up to 100 km wide and is divided into two branches in theregion shown in this image. Using digital elevation models infers flex-ural uplift of the graben flanks and a total topography ranging severalkilometers from the flanks to the graben floor (Giese et al. 2007). Imagedetails: ISS frame number N1589080288, scale 1 km/pixel, centered atlat. 42 ıS, long. 50 ıW in orthographic projection

(Moore et al. 1985) and appear in Cassini observations aspoorly developed ridges. This suggests that they are olderthan the grabens identified in the wispy terrain that appear tohave formed more recently (Wagner et al. 2006).

Iapetus’ equatorial ridge (Porco et al. 2005a; Giese et al.2008) (Fig. 20.7) is densely cratered, making it comparablein age to any other terrains on this moon (Schmedemann et al.2008). Models explain the origin of the ridge either by en-dogenic shape and/or spin state changes (e.g., Porco et al.2005a; Giese et al. 2008) or exogenic causes, with the ridgeaccumulating from ring material left over from the formationof proto-Iapetus (Ip 2006). On Iapetus, diurnal tidal stressis rather low, indicating that the ridge is old, and its preser-vation suggests the formation of a thick crust early in themoon’s history (Castillo-Rogez et al. 2007).

Enceladus, where cryovolcanism is definitely active(Porco et al. 2006), is discussed in Spencer et al. (2009).Titan, where some evidence of cryovolcanism exists (Lopeset al. 2007; Nelson et al. 2009a, b) is discussed in Jaumannet al. (2009) and Sotin et al. (2009).

There is no direct evidence of cryovolcanic processes(Geissler 2000) on the other Saturnian satellites. Although ithas been suggested in pre-Cassini discussions that the plainson Dione (Plescia and Boyce 1982; Plescia 1983; Moore1984) and Tethys (Moore and Ahern 1983) are endogenic inorigin, Cassini has not directly observed any cryovolcanicactivity on these satellites. The plains correlate with tectonicfeatures and exhibit distinct boundaries with more denselycratered units (Schenk and Moore 2007), which indicate

Fig. 20.25 Map of Dione’s trailing hemisphere showing an age se-quence of troughs, graben and lineaments, inferred from various cross-cutting relationships (Wagner et al. 2009). Three age groups of troughsand graben can be distinguished: Carthage and Clusium Fossae areoldest, cut by the younger systems of Eurotas and Palatine Chasmata,

which in turn are truncated by Padua Chasmata which are the youngest.Several systems of lineaments with various trends, either parallel orradial (the latter previously assumed to be a bright ray crater namedCassandra) can be mapped which appear to be younger than the troughsand graben


that the plains are a stratigraphically younger formation. Inaddition, there is evidence of material around Dione, Tethysand even Rhea (Burch et al. 2007; Jones et al. 2008; Clarket al. 2008a), but its origin is unclear. There is no evidenceof cryovolcanism on Iapetus. Although the plains on Dioneand Tethys are younger than the surrounding terrain, cryovol-canic activity must have occurred in ancient times, assumingthat such processes caused these surface units.

All medium-sized icy satellites appear to have a debrislayer, as may be inferred from (1) the non-pristine mor-phology of most craters; (2) the presence of groups of largeboulders and debris lobes seen in the highest-resolution ISSimages; and (3) derived thermal-physical properties consis-tent with loose particulate materials. With the exception ofEnceladus, which is addressed in Spencer et al. (2009), thesedebris layers were generated by impacts and their ejecta onthe other satellites. The upper limiting case for a given satel-lite is crustal disruption from extreme global impact-inducedseismic events. A way to evaluate the depth of fracture insuch events is to assume that these events were energeticenough to open fractures in the brittle ice (brittle on thetimescale of these events) composing the interiors of the tar-get satellite, and that, during the impact event, the lithostaticpressure everywhere was less than the brittle tensile strengthof cold ice. Brittle fracture caused by global seismicity, allelse being equal, extends deeper into smaller satellites, giventhat larger satellites are sheltered from dynamic rupture bytheir interior pressure. These being transient events, a simpleestimate of the realm where brittle fracture may be expectedto be induced, and visibly sustained, can be obtained fromglobal seismicity by applying an expression (e.g., Turcottand Schubert 2002) of the lithostatic pressure at radius Rp,and comparing this to the brittle tensile strength of waterice at timescales in which extreme global impact-inducedseismic events produce energy sufficient to break ice me-chanically. The determination of the fracture depth is basedon laboratory-derived dynamic values of �1MPa (Langeand Ahrens 1987; Hawkes and Mellor 1972). With the(bulk) density of the body (approximated as 1; 000 kg=m3),G the gravitational constant, and R the radius of the satel-lite, the depth of the fracture can be constrained. Based onthis, the value of R � Rp for Mimas is �20 km (�10% ofits radius), �10 km (�2% of radius) for Tethys, and �5 km(�0:5% of radius) for Rhea. This is a measure of the depthto which it is reasonable to expect the surface layer (i.e.,the megaregolith) to be intensely fractured. Of course, abody may be fractured to greater depths by large impacts,and the resulting pore space may not easily be eliminated.Deep porosity is most easily maintained on smaller and pre-sumably colder bodies such as Mimas (Eluszkiewicz 1990).Thus, for Mimas we expect that relatively recent major

impact events, such as Herschel, might exploit pre-existingfracture patterns in the ‘megaregolith’ which might penetratedeep into the body rather than focus energy, as has been sug-gested for the formation of the plains of Tethys (e.g., Brueschand Asphaug 2004; Moore et al. 2004). Consequently, thelack of antipodal focusing or axial symmetry about a radiusthrough Herschel in the lineament pattern on Mimas may beunsurprising, if that pattern is indeed an expression of tec-tonic movements initiated by the Herschel impact.

The impact-gardened regolith of most middle-sized icysatellites can be assumed to be many meters thick if the re-golith in the lunar highlands is any guide. Calculations ofimpact-generated regolith, based on crater densities observedin Voyager images, yielded mean depths of 500 m for Mimas,740 m for Dione, 1,600 m for Tethys and 1,900 m for Rhea(Veverka et al. 1986).

This regolith will be further distributed by gravity-drivenmass wasting (non-linear creep, i.e., disturbance-driven sed-iment transport that increases nonlinearly with the slope dueto granular creep (e.g., Roering et al. 2001; Forsberg-Tayloret al. 2004)) which reduces high-standing relief on craters(e.g., lowering of crater rims and central peaks) and tec-tonic features (e.g., the brinks of scarps) as well as fillingdepressions (e.g., Howard 2007). The process of non-linearcreep, when not competing with other processes, producesa landscape, which has a repetitive consistent appearance.Impact-induced shaking and perhaps diurnal and seasonalthermal expansion facilitate such mass wasting. Note thatthis common expression of mass wasting on regolith-coveredslopes assumes no topography-maintaining precipitation offrosts on local topographical highs, as discussed below.

Bedrock erosion might be caused by the sublimation ofmechanically cohesive interstitial volatiles. The icy bedrockwould crumble, move down slope through nonlinear creepand form mantling slopes, perhaps leaving only topograph-ical highs (crests) composed of bedrock exposed. Thissituation would occur if the initial bedrock landscape dis-aggregates to form debris at a finite rate, so that exposure ofbedrock gradually diminishes, lasting longest at local crests.Such a situation might explain the landscape of Hyperion. Ifthere is re-precipitation of volatiles in the form of frosts onlocal topographical highs, these frosts would armor the to-pographical highs against further desegregation. It has beenproposed that landscapes evolved in this fashion on Callisto(Moore et al. 1999; Howard and Moore 2008). Thermallydriven global and local frost migration and redistribution mayoccur on Iapetus (Spencer et al. 2005; Giese et al. 2008;Spencer and Denk 2009), as previously discussed. This sur-ficial frost redistribution may occur regardless of whethervolatiles in the bedrock of Iapetus sublimate and contributeto bedrock erosion.


20.7 Conclusions

The icy satellites of Saturn show great and unexpecteddiversity. In terms of size, they cover a range from�1;500 km (Rhea) to �270 km (Hyperion) and even smaller‘icy rocks’ of less than a kilometer in diameter. The icysatellites of Saturn offer an unrivalled natural laboratory forunderstanding the geology of diverse satellites and their in-teraction with a complex planetary system.

Cassini has executed several targeted flybys over icy satel-lites, and more are planned for the future. In addition, thereare numerous other opportunities for observation during non-targeted encounters, usually at greater distances. Cassiniused these opportunities to perform multi-instrument obser-vations of the satellites so as to obtain the physical, chemicaland structural information needed to understand the geolog-ical processes that formed these bodies and governed theirevolution.

All icy satellites exhibit densely cratered surfaces. Al-though there have been some attempts to correlate craters andpotential impactor populations, no consistent interpretationof the crater chronology in the Saturnian system has been de-veloped so far (Dones et al. 2009). However, most of the sur-faces are old in a stratigraphical sense. The number of largebasins discovered by Cassini was greater than expected (e.g.,Giese et al. 2008). Differences in their relaxation state pro-vide some information about crustal thickness and internalheat flows (Moore and Schenk 2007). The absence of basinrelaxation on Iapetus is consistent with a thick lithosphere(Giese et al. 2008), whereas the relaxed Evander basin onDione exhibits floor uplift to the level of the original surface(Moore and Schenk 2007), suggesting higher heat flow.

The wide variety in the extent and timing of tectonicactivity on the icy satellites defies any simple explanation,and our understanding of Saturnian satellite tectonics andcryovolcanism still is at an early stage. The total extent ofdeformation varies wildly, from Iapetus and Mimas (barelydeformed) to Enceladus (heavily deformed and currently ac-tive), and there is no straightforward relationship to predictedtidal stresses (e.g., Nimmo and Pappalardo 2006; Matsuyamaand Nimmo 2007, 2008; Schenk et al. 2008; Spencer et al.2009). Extensional deformation is common and often formsglobally coherent patterns, while compressional deformationappears rare (e.g., Nimmo and Pappalardo 2006; Matsuyamaand Nimmo 2007, 2008; Schenk et al. 2008; Spencer et al.2009). These global patterns are suggestive of global mech-anisms, such as despinning or reorientation, which proba-bly occurred early in the satellites’ histories (e.g., Thomaset al. 2007b); present-day diurnal tidal stresses are unlikelyto be important, except on Enceladus (e.g., Spencer et al.2009) and (perhaps) Mimas and Dione (e.g., Schenk andMoore 2007). Ocean freezing gives rise to isotropic exten-sional stresses; modulation of these stresses by pre-existing

weaknesses (e.g., impact basins) may explain some of the ob-served long-wavelength tectonic patterns (e.g., Schenk andMoore 2007). Except for Enceladus, there is as yet no ir-refutable evidence of cryovolcanic activity in the Saturniansystem, either from surface images and topography or fromremote sensing of putative satellite atmospheres. The nextfew years will hopefully see a continuation of the flood ofCassini data, and will certainly see a concerted effort to char-acterize and map in detail the tectonic structures discussedhere. Understanding the origins of these structures and thehistories of the satellites will require both geological andgeophysical investigations and probably provide a challengefor many years to come.

Although the surfaces of the Saturnian satellites are pre-dominately composed of water ice, reflectance spectra in-dicate coloring agents (dark material) on all surfaces (e.g.,Fink and Larson 1975; Clark et al. 1984, 1986; Roush et al.1995; Cruikshank et al. 1998a; Owen et al. 2001; Cruikshanket al. 2005; Clark et al. 2005, 2008a, 2009; Filacchione et al.2007, 2008). Recent Cassini data provide a greater range inreflected solar radiation at greater precision, show new ab-sorption features not previously seen in these bodies, and al-low new insights into the nature of the icy-satellite surfaces(Buratti et al. 2005b; Clark et al. 2005; Brown et al. 2006;Jaumann et al. 2006; Cruikshank et al. 2005, 2007, 2008;Clark et al. 2008a, b, 2009). Besides H2O ice, CO2 absorp-tions are found on all medium-sized satellites (Buratti et al.2005b; Clark et al. 2005; Cruikshank et al. 2007; Brownet al. 2006; Clark et al. 2008a), which make CO2 a commonmolecule on icy surfaces as well. Various spectral features ofthe dark material match those seen on Phoebe, Iapetus, Hy-perion, Dione and Epimetheus as well as in the F-ring andthe Cassini Division, implying that the material has a com-mon composition throughout the Saturnian system (Clarket al. 2008a, 2009). Dark material diminishes closer to Sat-urn, which might be due to surface contamination of the in-ner moons by E-ring particles and some chemical alterationprocesses (Clark et al. 2008a). Although the general compo-sition of the dark material in the Saturnian system has notyet 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 thedark spectra (Clark et al. 2008a, 2009) can be modelled bynano-phase hematite mixed with ice and dark particles. Thisled Clark et al. (2008b, 2009) to infer that sub-micron ironparticles from meteorites, which oxidize (explaining spectralreddening) and become mixed with ice could best explain allthe colors and spectral shapes observed in the system. Ironhas also been detected in E-ring particles and in particlesescaping the Saturnian system (Kempf et al. 2005). Previ-ous explanations for the dark material were based on tholins(e.g., Cruikshank et al. 2005), but this is inconsistent with thelack of marked absorptions due to CH in the 3�m to 4�m


region as well as other major absorptions at wavelengthslonger than 2�m. Iron particles are consistent with previousCassini studies that suggested an external origin for the darkmaterial on Phoebe, Iapetus and Dione (Clark et al. 2005,2008a, b, 2009). Another remarkable issue relating to darkmaterial is the albedo dichotomy on Iapetus, whose leadinghemisphere is the darkest surface in the Saturnian system.In contrast to the outermost satellite, Phoebe, also coveredby dark material, Iapetus’ dark hemisphere is spectrally red-dish, making it doubtful whether Phoebe is the source of thedark material. Clark et al. (2008a, 2009) explained the colordiscrepancy by red fine-grained dust .<0:5 �m/ that, mixedin small amounts into water ice, produces a Rayleigh scat-tering effect enhancing the blue response, so that the colorsobserved simply depend on the amount of dark material.

Radar observations constrain the dark material on Iapetusto a thickness of only a few decimeters (Ostro et al. 2006),which is consistent with the thin dark coating derived fromspectral data (Clark et al. 2008a, 2009) as well as with thelow density of 1:1 g=cm3 of Iapetus (Rappaport et al. 2005).In addition, the coated surface in the dark terrain is nowherebroken up by large impacts, suggesting that the emplacementof dark material occurred relatively recently (Denk et al.2000). In addition to its brightness dichotomy, Iapetus alsoshows a hemispheric color with the leading side dark materialbeing redder than the trailing side bright material (Bell et al.1985; Buratti and Mosher 1995). This also seems to correlatewith the low degree of macroscopic roughness on the darkside (Lee et al. 2009), indicating the infilling of rough surfacestructures. Moreover, there are also latitude-dependent spec-tral variations in the UV (Hendrix and Hansen 2009b) thatindicate temperature decreases and increases in the bright re-gions, particularly at the dark/bright boundary. These obser-vations are consistent with a model process of thermal seg-regation in which water ice leaves behind a residue of darkmaterial (Denk and Spencer 2009).

Alteration mechanisms cause not only compositional vari-ations on the satellites but also structural differences. Thesatellites serve as sources and sinks of dust grains within theSaturnian system. Plume activity on Enceladus, meteoriticbombardment and subsequent escape of particles, capture ofinterplanetary dust by Saturn’s gravitational field, and ringparticles all serve to provide material that is subsequentlyaccreted by the satellites. The fate of these dust grains inthe system depends on a complex set of interactions deter-mined by competing processes that move particles inward oroutward. Forces at work include gravitational perturbations,plasma drag, radiation pressure and Poynting-Robertson drag(Burns 1979). These mechanisms are sensitive to particlecomposition and size; furthermore, the latter diminishes overtime as particles are subjected to collisions.

Radar echos result from volume scattering, and as theirstrength is sensitive to ice purity they provide unique

information about near-surface structural complexity as wellas contamination with non-ice material. The observed radar-optical albedo correlation (Ostro et al. 2006) is consistentwith the explanation developed for the dark material in theSaturnian system in which Clark et al. (2008b, 2009) proposethat sub-micron iron particles from meteorites might oxidizeand become mixed with ice, assuming that the contaminationof the ice extends between one and several decimeters belowthe optically visible surface. Inter- and intra-satellite radaralbedo variations might be due to contamination by ammo-nia, which was tentatively detected in VIMS spectra (Clarket al. (2008a) and would affect the radar but not the opticalalbedo.

The Saturnian icy satellites exhibit stratigraphically oldsurfaces showing craters and early-formed tectonic structuresinduced by despinning, reorientation and volume changes.However, they also show evidence of younger geologicalprocesses such as tide-induced crustal stress and plains for-mation. Recent and current cryovolcanism exists on Ence-ladus, and ancient cryovolcanism cannot be ruled out forDione, Tethys and Rhea. Younger events, such as the forma-tion of ray craters, can be observed on Rhea. Surface alter-ation by charged-particle bombardment, sputtering, microm-eteorite bombardment and thermal processing are ongoinggeological activities. Surface coating by E-ring particles isalso a recent process. Coating of the satellites’ surfaces bydark material has been observed throughout the system, andthere is some evidence that this contamination might also bea relatively recent process.

Acknowledgements We gratefully acknowledge the long years ofwork by the entire Cassini team that allowed these data of the Saturniansatellites to be obtained. We also acknowledge NASA, ESA, ASI, DLR,CNES and JPL that provide 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 highlyappreciated. Part of this work was performed at the DLR Institute ofPlanetary Research, with support provided by the Helmholtz Alliance‘Planetary Evolution and Life’ and the Jet Propulsion Laboratory undercontract to the National Aeronautics and Space Administration.

We dedicate this work to Steve Ostro, who provided outstandingscientific contributions to the exploration of the Saturnian System.

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