chapter3 originofthesaturnsystem - |lasp|cu-boulder

Chapter 3Origin of the Saturn System

Torrence V. Johnson and Paul R. Estrada

Abstract Cassini mission results are providing new insightsinto the origin of the Saturn system and giant planet satel-lite systems generally. The chapter discusses current mod-els for the formation of giant planets and their satellites andreviews major Cassini findings which help advance our un-derstanding of the system’s formation and evolution to itscurrent state.

3.1 Introduction

The results of the Cassini/Huygens mission must be inter-preted in the context of what we know about the formationand history of the Saturn system. These results in turn pro-vide many constraints on and clues to the conditions and pro-cesses which shaped the system. The purpose of this chapteris not to provide a general tutorial or textbook on the the-ory of giant planet and satellite formation. This is currently avery active field of planetary and astrophysical research, withnew concepts being developed continually in response to theflood of new data from planetary missions and astrophysicalobservations of star- and planet-forming regions from groundand space-based telescopes. A full treatment of this topicis beyond the scope of the current book, let alone a singlechapter. The reader is referred to a number of recent pub-lications and reviews of the current state of research in thisfield (Canup and Ward 2009; Davis 2004; Estrada et al. 2009;Reipurth et al. 2007). What we hope to achieve in this chapteris to review the key concepts and current issues in planetaryformation and discuss Cassini/Huygens results that relate tothe problem of the origin of the system.

T.V. JohnsonJet Propulsion Laboratory, California Institute of Technology,MS 301-345E, 4800 Oak Grove Dr., Pasadena, CA 91109e-mail: [email protected]

P.R. EstradaCarl Sagan Center, SETI Institute, 515 N. Whisman Rd., MountainView, CA 94043e-mail: [email protected]

Historically, Saturn’s system was the next most readilyavailable after Jupiter’s for early telescopic observations.By the end of the seventeenth century, less than 100 yearsfollowing Galileo’s first astronomical observations with thetelescope, the general structure of Saturn’s ring and satellitesystem was known, encompassing its spectacular rings, gi-ant Titan and four of the seven medium sized ‘icy satellites’of the regular satellite system (satellites in nearly equato-rial, circular orbits). During the subsequent three centuriesof telescopic observations many details of the ring struc-ture were revealed and new satellites discovered, includingPhoebe, the first of the irregular outer satellites to be discov-ered. The Pioneer 11 and Voyager 1 and 2 flybys in 1979,and 1980/81 provided the first reconnaissance of the systemby spacecraft, discovering and characterizing its magneto-sphere and exploring the moons and rings from close rangefor the first time, setting the stage for the Cassini/Huygensmission.

In the past, the Jupiter and Saturn systems were frequentlyreferred to as archetypical examples of ‘miniature solar sys-tems’ and it was suggested that their ring and satellite sys-tems were possibly formed in an analogous manner to theplanets about the Sun. Current models of star and planetarysystem formation suggest a more complex picture, with arange of processes and timescales leading to the formation ofthe sun from interstellar material, an early gas and dust so-lar nebula, and finally planetary formation. Cassini/Huygensobservations have provided important new information andconstraints on many characteristics related to the system’sorigin, including the composition of Saturn, the compositionof the satellites and rings, the current dynamical state of thesatellites (spin, orbital eccentricity, resonances), and satel-lites’ internal structure and geological history.

In this chapter we will first review briefly current mod-els for the formation of gas giant planets such as Jupiter andSaturn and the formation of their satellite systems. The sec-ond portion of the chapter discusses Cassini results related tothe question of origins, emphasizing the chemical conditionsfor condensation of material in the outer solar system, andnew information about the densities and structures of Sat-urn’s satellites.

M.K. Dougherty et al. (eds.), Saturn from Cassini-Huygens,DOI 10.1007/978-1-4020-9217-6_3, c� Springer Science+Business Media B.V. 2009

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56 T.V. Johnson and P.R. Estrada

3.2 Planet and Satellite Formation

3.2.1 Big Bang to the Solar Nebula

The events and conditions leading eventually to planetaryformation in our solar system can be traced back to the BigBang, about 13.7 Ga, based on current measurements of cos-mological constants (Spergel et al. 2007). Galactic formationbegan within a few Gyr or less. Current estimates for the ageof our Milky Way Galaxy are �10–13 Ga with some starsin associated globular clusters dating back even earlier. Thematerial out of which our solar system formed was the resultof billions of years of stellar evolution and processing, yield-ing the mix of elements in what is usually called ‘cosmic’ or‘solar’ abundance.

It is generally agreed that the Sun and its associatedcircum-stellar disk formed from the collapse of a largeinterstellar molecular cloud, very probably in associationwith a massive-star-forming region similar to that seen in theOrion nebula (Boss 2007). Two basic mechanisms for thecollapse have been proposed, gravitational instability andcollapse triggered by a shockwave from a supernova or othernearby stellar event. Both mechanisms appear to be viable,but have different timescales and implications for the earlysolar nebula, with gravitational collapse taking �10Myrwhile triggered collapse would be more rapid, �1Myr. Thepresence of short lived radioactive isotopes in the early solarsystem materials has favored the supernova shock-triggeredmodel (e.g., 26Al, and particularly 60Fe, which can only bemade in a supernova), with other shock sources, such asoutflow from massive AGB stars, regarded as less probablebased on stellar evolution models (Boss 2007; McKeeganand Davies 2007). Presolar grains and the daughter productsof the now extinct short lived isotopes found in primitivemeteorites are the only remaining physical evidence of thisearly phase of the solar system’s formation.

As the Sun formed from the collapsing molecular cloud,surrounded first by an envelope of gas and then an ac-cretion disk, the earliest condensable materials appeared.These materials, preserved in refractory fragments in primi-tive meteorites, are known as Calcium Aluminum Inclusions(CAIs), and provide our first direct ties to the timeline andconditions which led to planet formation. CAIs have beenradiometrically dated using lead isotopes, yielding an age of4:567:2˙ 0:6Ga (Amelin et al. 2002).

3.2.2 The Solar Nebula to Planets

3.2.2.1 The Inner Solar System

How rapidly the planets formed from the gas and dust of theearly solar nebula is the subject of much current research and

theoretical analysis. Isotopic dating of terrestrial materials,meteorites, lunar samples and Martian meteorites providesconstraints on the formation times for meteorite parent bod-ies and the terrestrial planets. These data can be correlatedusing both long-lived (e.g., the U-Pb system) and short-livedchronometers (e.g., 26Al=27Al and 53Mn=54Mn) and showthat the formation of chondrules and the parent bodies of ig-neous meteorites (Eucrites) began within a few million yearsfollowing CAI condensation (McKeegan and Davies 2007).The age of the Earth and Moon can be inferred from studiesof the oldest rocks on these objects and isotopic data con-straining the time of core formation on the Earth. There arestill significant uncertainties in these interpretations whichare beyond the scope of this chapter to explore. A recent re-view of relevant studies up to �2006 has the accretion of theEarth and core formation complete by �4:46Ga, with theMoon forming from a giant impact during the late stages ofEarth’s accretion �4:52Ga (Halliday 2007).

A dated event potentially related to the final stages ofplanet formation is the Late Heavy Bombardment. This isbased on a strong clustering of lunar impact breccia agesat �3:9Ga and the lack of breccias with significantly olderages. These data have been interpreted either as the ‘tail-off’of accretion or as a real ‘spike’ in the impact cratering rate inthe inner solar system at 3.9 Ga (see Warren 2007). Recentstudies of the dynamical evolution of the early solar systemhave provided support for the ‘spike’ or ‘lunar cataclysm’explanation (see discussion of migration and the ‘Nice’model below).

3.2.2.2 The Outer Solar System

For the outer solar system, we do not yet have sample-derived dates for the giant planets or their satellites, and sothe formation times must be constrained by astrophysicalobservations of other star systems and theoretical studies ofgiant planet formation. Two major advances in astrophysicalresearch have recently led to a greatly improved understand-ing of planetary formation around stars similar to our Sun:the explosive rate of discovery of extrasolar planets in thelast decade and the study of protoplanetary disks, dust disksand debris disks around other stars with powerful groundand space-based telescopes (such as the Keck telescope,Hubble Space Telescope and the Spitzer mission). Twoimportant results from this rapidly growing area of researchcan be summarized as: 1. Planetary formation is frequentlyassociated with star formation and 2. Giant, gas-rich planetsform very quickly – in less than 10 million years for mostsystems and in many cases within a few million years. Inparticular, evidence for gas loss in the disks around youngstars suggest that giant planets (which must form beforemost of the gas is removed from a planetary disk) form inless than 107 year (Meyer et al. 2007) and gap clearing in

3 Origin of the Saturn System 57

dust disks, interpreted as evidence for planet formation, isseen in many young systems, the youngest a system only�106 year old (Espaillat et al. 2007).

3.2.2.3 Giant Planet Formation

Discussions of giant planet formation have been dominatedby two competing concepts for initiating the formation andcapture of essentially solar composition gas from the sur-rounding solar nebula. These models are generally referred toas the gas instability model and the core nucleated accretionmodel. There are proponents of both approaches and debatecontinues on the pro’s and con’s of each. Currently, the core-accretion approach is more heavily favored by researchers inthe field, although there is continuing development of boththeories in response to new observations and constraints onthe origin of both our solar system and exoplanetary systems.

Gas Instability Models

The gas instability (GI) approach to giant planet formationinvolves the formation of dense clumps in the solar neb-ular gas disk due to non-linear instabilities. In this modelthese clumps are the self gravitating precursors of the giantplanets (Cameron 1978; Durisen et al. 2007; Kuiper 1951).Recent work on GI models for both solar system and exo-planet formation is associated with Alan Boss and his col-laborators. The theoretical development and numerical mod-eling of GIs in pre-stellar gas disks is complex and thereis still considerable debate about whether GIs could lead togas giant formation in realistic models of the solar nebula(Bodenheimer et al. 2000; Durisen et al. 2007; Lissauer andStevenson 2007). An advantage of the GI models is that theynaturally lead to formation of giant planets very early in asystem’s lifetime, in agreement with mounting evidence thatat least some stellar systems have gas giants forming withina few million years. The GI models do not do as well atexplaining the enrichment of heavy elements over solar abun-dances in Jupiter and Saturn, since the pure instability mod-els should produce planets reflecting the composition of thelocal nebula. GI models also do not address directly the for-mation of the terrestrial planets, although they are not intrin-sically incompatible with solid planet formation by planetes-imal accretion.

Core Nucleated Accretion Models

Core nucleated accretion models derive from the theories de-veloped by Safronov (Safronov 1967, 1969, 1991; Safronovand Ruskol 1994) and others to explain the accretion of theterrestrial planets from the collision of planetesimals in a

planetary accretion disk. The extension of this approach togiant planets envisions the growth of cores of rock and icewhich then accrete gas from the local nebula when they be-come massive enough to hold substantial atmospheres (typ-ically a few M˚). Continued gas accumulation dominatesthe final stage of planet growth until the gas is depleted inthe planet’s vicinity, either through tidally opening a gapin the gas disk, or the nebula gas dissipating, upon whichplanetary growth is halted – e.g. (Estrada et al. 2009; Lissauerand Stevenson 2007).

Core nucleated models with cores enriched in rock andice are qualitatively better able to explain the observed non-solar giant planet compositions than GI models (Estradaet al. 2009). The differences between Jupiter and Saturnand the observed enrichments in noble gases in Jupiter’s at-mosphere, however, are not yet fully explained by any onetheory.

A difficulty with the traditional treatment of core accre-tion models has been the time scale for accretion of the req-uisite mass in the core. Extensions of Safronov’s theory to theproblem suggested time scales of 108 to 109 year or longer(Safronov 1969; Wetherill 1980) would be required for late-stage, high velocity growth in the outer solar system. Giventhe requirement that, by definition, the gas giants must formbefore the solar nebula gas dissipates, and the growing evi-dence noted earlier for formation of extrasolar giant planetsin<107 years, mechanisms for much more rapid core growthare needed.

Lissauer has suggested that this implies that run-away/oligarchic growth (Kokubo and Ida 1998) must becapable of growing large cores at the distance of Jupiter,which may in turn require higher surface mass densitiesthan given by calculations of the ‘minimum mass solar neb-ula’ (Estrada et al. 2009; Lissauer 1987, 2001; Lissauer andStevenson 2007). In models incorporating these concepts,the time scale of giant planet formation is strongly depen-dent on the mass of the solid core, the rate of energy inputfrom continued accretion of solids, and the opacity of thegaseous envelope. There are uncertainties in the appropriatevalues for all of these quantities, but models for reasonableranges are capable of forming giant planets in �1–12Myr(Hubickyj et al. 2005; Lissauer et al. 2009). Figure 3.1 illus-trates the evolution of a giant protoplanet from models ofJupiter’s growth (Lissauer et al. 2009).

3.2.2.4 Planetary Migration and the Nice Model

Earlier generations of planetary formation models startedwith the assumption that the major planets formed at or veryclose to their present distance from the sun, although theeffects of orbital resonances and orbit changes due to closeencounters with the giant planets were recognized for theirroles in the current structure of the asteroid belt (e.g., the


Fig. 3.1 Evolution of a proto-Jupiter: The mass of solids in the planet(solid line), gas in the planet (dotted line) and the total mass of theplanet (dot-dashed line) are shown as functions of time. Note the slow,

gradually increasing, buildup of gas, leading to a rapid growth spurt, fol-lowed by a slow tail off in accretion. Figure from Lissauer et al. (2009)

Kirkwood gaps) and comet orbits (e.g., the Oort cloud).However, planets will interact gravitationally with the disk ofmaterial from which they are forming and the gravitationaltorque between a planet and the disk should lead to migra-tion of the planet’s orbit (Goldreich and Tremaine 1980;Ward 1986, 1997). Many of the recently discoveredexoplanet systems have giant planets orbiting extremelyclose to their primary star (inside 1 AU), strongly suggestingthat they must have migrated inward from their formationregions in the more distant parts of the circumstellar disk.Current modeling of the early solar system addresses thepossibility that significant rearrangement of planetary orbitsmay have occurred in our solar system as well.

How a planet migrates during growth in a massive gas diskdepends on its location and whether it is massive enough toclear a gap in the disk. For lower mass objects, migrationis dominated by what is known as Type I migration, due tothe difference in torques exerted on material outside of andinward of their orbits. Planets which clear a gap then are sub-ject to Type II migration, whereby they are in effect draggedalong with the viscously evolving disk. Type I migration poseproblems for the growth of large cores in the core nucleatedaccretion scenario and Type II migration creates potentialdifficulties for both core nucleated and gas instability models.

Once the giant planets have formed, a further compli-cation is that in a system with multiple giant planets, theplanets will create instabilities in the planetesimal diskthrough orbital resonances, which increase the eccentricitiesof planetesimals and can also affect the orbits of the other

major planets, creating planet-induced migration. As theplanets migrate, these resonances will ‘sweep’ though thesystem, creating the possibility of periods of rapid rearrange-ment of planetary orbits. Investigation of these effects withlarge scale computer numerical simulations is currently arapidly growing and evolving discipline. For recent reviews,see (Levison et al. 2007; Lissauer and Stevenson 2007;Papaloizou et al. 2007).

The most comprehensive model to date for the dynamicalevolution of the outer solar system through the interactionsof the giant planets and the planetesimal disk has been devel-oped by researchers at the Nice Observatory in France andtheir collaborators. Known as the ‘Nice model’ and presentedin a series of papers (Gomes et al. 2005; Levison et al. 2008;Morbidelli et al. 2005; Tsiganis et al. 2005), the model at-tempts to explain the current location and orbital characteris-tics of the outer planets as well as the structure of the KuiperBelt though the combined effects of orbital migration and or-bital resonances. Numerous possible initial configurations ofthe outer planets were considered. The version most closelymatching the observed characteristics of the current outer so-lar system starts with a more compact configuration of plan-ets and the relative positions of Neptune and Uranus reversedfrom their final locations. In this scenario the planetary or-bits evolve relatively slowly initially until Jupiter and Saturnreach a 1:2 mean motion resonance condition, at which pointthe orbits of all the outer planets as well as the planetesimaldisk of the proto-Kuiper Belt are strongly perturbed, mostof the planetesimal disk is dispersed and the outer planets


Fig. 3.2 The planetary orbits and the positions of the disk particles,projected on the initial mean orbital plane: The four panels correspondto four different snapshots taken from our reference simulation. In thisrun, the four giant planets were initially on nearly circular, co-planar or-bits with semimajor axes of 5.45, 8.18, 11.5 and 14.2 AU. The dynam-ically cold planetesimal disk was 35ME with an inner edge at 15.5 AU

and an outer edge at 34 AU. Each panel represents the state of the plan-etary system at four different epochs: a, the beginning of planetary mi-gration (100 Myr); b, just before the beginning of LHB (879 Myr); c,just after the LHB has started (882 Myr); and d, 200 Myr later, whenonly 3% of the initial mass of the disk is left and the planets haveachieved their final orbits. From Gomes et al. (2005)

settle into their current more extended orbital configuration(Gomes et al. 2005). Figure 3.2 illustrates the history of theplanetary orbits for this case.

In addition to rearranging the giant planets’ orbital char-acteristics, the Nice scenario also may explain the existenceand nature of the extensive population of irregular satellitesat Jupiter and Saturn. Several dynamical studies have sug-gested that capture of scattered planetesimals can produce theobserved characteristics of these irregular systems (Nesvornyet al. 2007; Turrini et al. 2008, 2009).

An interesting consequence of this model for the earlyhistory of the Saturn system is the timing of the resonance-driven rearrangement (Gomes et al. 2005). The conditionsfor this version of the Nice model were chosen to require theJupiter/Saturn 1:2 resonance to occur at the time of the so-called Lunar Late Heavy Bombardment (LHB) at �3:9Ga.The Nice model provides a plausible scenario for creatingsuch a spike, and also implies that the LHB should have been

a solar system-wide event. While the impacting objects onthe Moon in this scenario may originate both from the grav-itational scattering of bodies from the asteroid belt into theinner solar system and outer solar system planetesimals, thedisruption of the planetesimal disk in the outer solar systemwould have created a spike in large impacts throughout theouter solar system at the same time. The large impact basinson the oldest surfaces on Saturn’s icy satellites, such as Iape-tus, may be evidence for this event (Johnson et al. 2007).

3.2.3 Circumplanetary Disks to Satellites

The satellites of the giant planets are characterized by pro-grade, low inclination orbits that are relatively compactin their spacing compared to the radial extent of theirplanet’s Hill sphere. This indicates that they formed within


a circumplanetary disk around their parent planet. Given thatall the giant planets of our solar system possess regular satel-lite systems, it suggests that satellite formation may be a nat-ural consequence of giant planet formation. Thus in order toaccount for the formation of their satellite systems, modelsof giant planet accretion must include the formation of theirassociated circumplanetary disk.

3.2.3.1 Formation of the Subnebula

The circumplanetary gas disk, or subnebula, is a by-productof the giant planet’s later stages of accretion and forms overa period in which the giant planet transitions from runawaygas accretion to its eventual isolation from the circumstellardisk. This isolation, which effectively signals the end of thegiant planet’s formation, occurs because the giant planet suc-cessfully opens a deep, well-formed gas gap in the nebula, orthe nebula gas dissipates. The process of gap-opening by thegiant planet is relevant to subnebula formation because theradial extent of the circumplanetary gas disk is determinedby the specific angular momentum of gas that enters the gi-ant planet’s gravitational sphere of influence.

The contraction of the giant planet envelope down to afew planetary radii occurs relatively quickly compared to therunaway gas accretion phase (see Lissauer et al. 2009), mean-ing that the subnebula likely begins to form relatively earlyon in the planetary formation process when the protoplanetis not sufficiently massive enough to open a substantial gasgap in the surrounding nebula. Prior to opening a deep, wellformed gap, the planet accretes low specific angular momen-tum gas from its vicinity, which results in the formation ofa rotationally supported compact disk component. This disksize is consistent with the observed radial extent of the gi-ant planet satellite systems (Mosqueira and Estrada 2003a).Furthermore, since the protoplanet may still need to accretethe bulk of its gas mass after envelope collapse, this suggeststhat the gas mass deposited over time in the relatively com-pact subnebula could be substantial compared to the mass ofthe satellites, a picture that may be consistent with the viewthat the planet and disk may receive similar amounts of angu-lar momentum (Mosqueira and Estrada 2003a, b; Stevensonet al. 1986). One potential caveat regarding Saturn is that itslower mass may mean that the time between envelope col-lapse and its final mass could be considerably shorter thanthe Jupiter-mass case.

As the giant protoplanet grows more massive and the gapbecomes deeper, the continued gas inflow through this gapcan significantly alter the properties of the subnebula. In par-ticular, as gas in the protoplanet’s vicinity is depleted, theinflow begins to be dominated by gas with specific angularmomentum which is much higher than what previously ac-creted. This is because the gas must now come from increas-

ingly farther away in heliocentric distance. A more extended,less massive disk component forms as a result. Recent sim-ulations of gas accretion onto a giant planet embedded in acircumstellar disk that are able to resolve structure on thescale of the regular satellite systems indicate that the sizeof the disk formed by the inflow through the gap likely ex-tended as much as �5 times the size of the observed regularsatellite systems of Jupiter and Saturn (D’Angelo et al. 2003)(see also Ayliffe and Bate 2009, who specifically treat bothJupiter- and Saturn-sized planets). The low and high specificangular momentum gas accretion phases along with the ob-served mass distribution of the regular satellites of Saturn(and Jupiter) thus argue in favor of a two-component circum-planetary disk: a compact, relatively massive disk that formsover the period prior to the opening of a deep, well formedgap; and, a much more extended, less massive outer disk thatforms from gas flowing through the gap and at a lower in-flow rate (Bryden et al. 1999; D’Angelo et al. 2003). Currentmodels for disk and satellite formation for the Jupiter sys-tem are discussed more extensively in Estrada et al. (Estradaet al. 2009; See also Canup and Ward 2009).

3.2.3.2 Conditions for Satellite Accretion

Satellite formation is expected to begin at the tail end ofplanetary formation. Therefore, the character of the gasinflow through the gap during the waning stage of the giantplanet’s accretion is a key issue in determining the conditionsunder which the regular satellites form. A consequence ofpersistent gas inflow through the gap is that the subnebulawill continue to evolve due to the turbulent viscosity gen-erated as a result of gas accretion onto the circumplanetarydisk. If the subnebula is turbulent, it can affect satelliteformation in two major ways. The first is that (even weak)turbulence can pose a problem for satellitesimal formation(the subnebula equivalent of planetesimals). Dependingon its strength, nebula turbulence presents a problem forparticle growth because the relative velocities between grow-ing particles can quickly become large enough to lead tofragmentation rather than growth (e.g., Dominik et al. 2007).This problem is exacerbated in the circumplanetary diskenvironment where gas densities, temperatures, pressures,and dynamical times are considerably different than theircircumsolar analog (Estrada et al. 2009). Second, even weak,ongoing inflow through the gap can generate a substantialamount of viscous heating which would generally result in acircumplanetary disk that is too hot for ice to condense andsatellites to form and survive (Coradini et al. 1989; Klahrand Kley 2006; Makalkin et al. 1999).

Despite these difficulties several models in the litera-ture attempt to explain satellite formation under turbulentconditions, where satellite survival hinges on the turbulent


dissipation of the gas disk (e.g., Alibert et al. 2005;Canup and Ward 2002; Makalkin et al. 1999; Mousis andGautier 2004). However, given the complications associatedwith satellite accretion in even weak turbulence generatedby the inflow, Mosqueira and Estrada (Mosqueira andEstrada 2003a, b) assumed that satellite formation did not be-gin until the gas inflow through the gap wanes, at which pointturbulence in the subnebula may decay. Once turbulence de-cays, the circumplanetary gas disk becomes quiescent, so thatone can not rely on turbulence to dissipate the gas disk. Theyassumed this because there are currently no intrinsic mech-anisms that have been identified that can sustain turbulencein a relatively dense, mostly isothermal subnebula (Estradaet al. 2009). Instead, Mosqueira and Estrada (2003b) showedthat the largest satellites (e.g., Titan, Ganymede), and notdisk dissipation due to turbulence, may be responsible for gi-ant planet satellite survival. This represents a key differencethen between these authors’ model and those listed above.

Another key issue is the source of solids which is inti-mately tied to the explanation for the observed mass andangular momentum of the regular satellite systems. Thereare several ways in which solids can be delivered to thecircumplanetary disk environment, which basically fall intotwo categories: solid delivery via planetesimals (e.g., Estradaand Mosqueira 2006; Mosqueira and Estrada 2003a, b); andthe delivery of solids via dust coupled to the gas inflowduring planet accretion (e.g., Canup and Ward 2002, 2009;Makalkin and Dorofeeva 2006).

Planetesimal delivery mechanisms offer a straightforwardway to explain the mass and angular momentum of the reg-ular satellites, and help to explain the enhancement of heavyelements in the atmospheres of the giant planets (Estradaet al. 2009). At the time of planet formation, most of themass in solids are in planetesimals with sizes larger thana kilometer (Charnoz and Morbidelli 2003; Kenyon andLuu 1999; Wetherill and Stewart 1993). As the planet ap-proaches its final mass, planetesimals undergo an intense pe-riod of collisional grinding where a significant fraction ofmaterial is fragmented into smaller objects (Charnoz andMorbidelli 2003; Stern and Weissman 2001), but are stilltoo large to be coupled to the gas. Much of the mass inthese planetesimal fragments may then pass through the cir-cumplanetary gas disk where, depending on the gas surfacedensity, they may ablate, melt, vaporize, and/or be captured(Estrada et al. 2009; Mosqueira and Estrada 2003a, b). In thedecaying turbulence model of Mosqueira and Estrada, gassurface densities are sufficiently high that enough materialmay be deposited in this way to explain the mass and angu-lar momentum of the regular satellites. Furthermore, signif-icant increases in the ice/rock ratio can occur because thelargest first-generation planetesimal bodies (which containmost of the mass) can differentiate if a significant amountof short-lived radionuclides such as 26Al are present (Jewitt

et al. 2007; Merk and Prialnik 2003). Subsequent break up ofthese planetesimals may then lead to fractionation. A lowerdensity outer circumplanetary gas disk might then preferen-tially ablate much more ice relative to rock (although muchless total mass overall) compared to the denser, inner regions.Thus, for example, Iapetus’ low density (compared with Cal-listo at Jupiter) would suggest that the Saturnian outer gasdisk was less dense relative to that of Jupiter’s.

Other formation scenarios which consider satellite accre-tion in a gas-starved disk under turbulent conditions do notinclude planetesimal delivery mechanisms, and rely on dustcoupled to the gas inflow as the primary source of solids forgenerations of satellites, most of which are lost due to mi-gration into the planet (Canup and Ward 2002). Canup andWard (2009) note that the ablation of planetesimals cannotbe a significant source of solids in their gas-starved model,presumably because their gas disk density is too low. On theother hand, it has been argued that dust inflow cannot be thedominant source of material for the satellites for a varietyof reasons, and does not explain their angular momentum(see Estrada et al. 2009, and references therein). However,dust inflow may play a role in providing supplement mate-rial for the outermost satellites. For example, Mosqueira andEstrada (2003a) considered that dust inflow through the gapmay augment the mass of Callisto (and Iapetus), but pointedout that too much mass delivered this way would lead to theformation of satellites in very cold regions of the disk (wherethey are not observed) with their full complement of ices.

In an effort to circumvent the difficulties with coagula-tion of dust into satellitesimals in high orbital frequency gasdisks, an alternative view is that an intrinsic source of tur-bulence acts to remove most of the gas disk in a timescaleshorter than the satellite formation timescale (Estrada andMosqueira 2006). Although the residual gas density is leftunspecified in this model, it is assumed to be low enoughthat the angular momentum of the satellite system is largelydetermined by circumsolar planetesimal dynamics and notgas dynamics (i.e., dust may be entrained in the gas, butsatellitesimals and embryos remain largely unaffected). Inthis case, planetesimal material (mostly in the form of frag-ments) is captured and deposited in the circumplanetary diskthrough inelastic and gravitational collisions between plan-etesimals within the giant planet’s Hill sphere rather thanthrough ablation. Intermediate gas densities (where ablationwould still remain unimportant) would actually make the col-lisional capture of planetesimals easier. One consequence ofa slightly higher gas density is that, since most of the masscaptured is in the form of large fragments (which sweep upthe dust content in the disk), the planetesimal capture modelcan by-pass the problem of coagulative growth in a turbulentenvironment faced by models that rely only on dust inflow.However, the planetesimal capture model would still facemany of the same problems associated with the gas-starved


disk. For example, a similar problem faced by both Estradaand Mosqueira (2006) and Canup and Ward (2002, 2009) isthat they cannot explain the origin of Iapetus.

Explaining the current state of the satellite systems iscomplicated by many unknowns; yet a great deal of progresscontinues to be made thanks to Cassini/Huygens as well aspast spacecraft missions. A thorough treatment of the variouspossibilities and stages of satellite formation for the Jupitersystem in which planetesimal delivery mechanisms providethe source of solids has been done by Estrada et al. (Estradaet al. 2009), and applies in its basic aspects to the Saturn sys-tem as well. One perhaps key difference between the Jovianand Saturnian systems is the longer planetesimal ejectiontimes at Saturn relative to Jupiter (Goldreich et al. 2004).Although both Jupiter and Saturn would have experiencedthe possibility of disruptive impacts deep within their respec-tive potential wells, longer satellite accretion times coupledwith longer planetesimal ‘exposure’ times at Saturn relativeto Jupiter may have made it more likely for satellites deepwithin Saturn’s potential well to be disrupted by impactshelping to explain their present day configuration (Estradaand Mosqueira 2006; Mosqueira and Estrada 2003a).

Finally, in a scenario in which the giant planets of theSolar System start out much closer to each other and sub-sequently migrate outwards as in the Nice model (Tsiganiset al. 2005) the two planets jointly open a gap (if they growalmost simultaneously; see, e.g, Bryden et al. (2000)). Insuch a scenario, Saturn’s local influence would dominateover that of turbulence in a moderately viscous circumsolarnebula (Morbidelli and Crida 2007). As a result, gap forma-tion could be delayed for both planets which may influencethe amount of high specific angular momentum gas they ac-creted. Indeed, the amount of gas left around after the open-ing of a combined gap may be a key factor in explaining thedifferences between the giant planet satellite systems, partic-ularly their outermost satellites. In the next section we willexplore how some of the results from the Cassini/Huygensmission bear on these issues and possibly provide better con-straints for current models.

3.3 Cassini Results and Discussion

3.3.1 Saturn and Rings

Saturn’s composition and structure provide clues to its originand evolution. Cassini data on the atmosphere and interior ofSaturn are covered elsewhere in this book (Chapters 2, 4 and5). With respect to the issues connected with gas instabilityvs core accretion formation models, the strong enrichment incarbon over solar abundances, known from ground-based andVoyager observations, has been confirmed by Cassini and

explored in greater detail. The existence of massive stormsystems associated with evidence for lightning argues, as forJupiter, for significant amounts of water in a condensablelayer at depth, but do not provide a strong quantitative con-straint on its relative abundance (Dyudina et al. 2007). Like-wise, the absence of measurements of noble gases prevents adirect comparison with Galileo Probe results.

To date, no new strong gravitational constraints on interiorstructure have been obtained. However the possibility of ahigh-inclination, low altitude phase for the Cassini orbiter atthe end of its extended operations may ultimately greatly im-prove this situation and allow comparison with similar mea-surements expected from the planned Juno mission to Jupiter.

How Saturn’s rings are linked to the origin of the sys-tem depends on their own origin and evolution. Generally,two classes of model have been discussed in the literature:1. Primordial models where the material in the rings is re-garded as material left over from the circumplanetary diskthat never accreted into satellites due to tidal forces (i.e.,inside the Roche limit), and 2. Satellite disruption models,where it is hypothesized that a satellite’s orbit decayed dueto tidal forces or drag, taking it inside the Roche limit whereit was subsequently disrupted by tidal forces, or, alternativelywhere a satellite was strong enough to exist inside the Rochezone but was disrupted by a large impact (either very early orrecently depending on the version) and the resulting debrisdisk could not re-accrete into a new satellite.

Following the discovery of striking gravitationally pro-duced wave structures in the ring system by Voyager, anal-yses of the angular momentum transfer involved in theirproduction suggested that the current ring system could notsurvive for the age of the solar system, strongly favoring thedisruption type models of recent origin. Cassini has mademany detailed observations of the rings using multiple tech-niques. These observations and results are described in detailin the relevant chapter in this book (Chapters 13–17). Someof these results imply that the ring system is more massiveand longer-lived than thought from Voyager results, focusinginterest again on primordial models, possibly involving con-tinual ‘renewal’ of locally ephemeral structures in the rings(Chapter 17).

3.3.2 Satellite Composition

3.3.2.1 Satellite Bulk Densities

The most important single constraint on the bulk composi-tion of the satellites comes from measurements of their av-erage densities. These can be used to infer the approximateproportions of condensed volatiles (primarily water ice) and‘rock’ (silicates plus metal content) in the body and compare


Fig. 3.3 Density of icy outer solar system bodies as a function ofradius: Most of the planetary satellite values are from spacecraft ob-servations of radius (volume) and mass from radio tracking and grav-ity science analyses of flybys. Data sources for the Ganymede andCallisto (orange diamonds), Uranus satellites (purple triangles), andTriton (green circle) and Titan (blue square) are found in (Yoder 1995).Amalthea data are from (Anderson et al. 2005). Mid-sized and smallSaturn satellite data (radii �10–1; 000 km, blue squares) are from(Jacobson 2004; Jacobson et al. 2006; Porco et al. 2005a, b, 2007;Thomas et al. 2007). Pluto and Charon and KBO’s (green circles) dataare, for Pluto/Charon from (Buie et al. 2006), Eris (Brown and Schaller

2007), Varuna (Jewitt and Sheppard 2002), 2003EL61 (Rabinowitzet al. 2006), 1999TC36 (Stansberry et al. 2006).The curves labeled“60/40 ice/rock” (by mass) and “Pure Ice” are theoretical calculations ofdensities of a sphere with a rock-ice (‘rock’ density D 3; 662 kg m�3)mixture and pure water ice composition, taking into account the pres-ence of high density phases of water ice in the interior of a larger body(Lupo and Lewis 1979). The curve labeled “Porous Ice” represents thedensity of an ice sphere with a porosity throughout equal to that allowedat the central pressure of the object, based on laboratory compressionexperiments with cold ice compaction (Durham et al. 2005)

these with models of material composition expected fromdifferent formation models and circumplanetary nebula con-ditions. Figure 3.3 shows measured or derived densities foricy bodies in the outer solar system including some KuiperBelt Objects and comets.

There are two major factors affecting the interpretationof bulk density in icy satellites, high density ice phases andporosity. The pressures in the interiors of even the largestsatellites are too low for high pressure silicate phases to besignificant (Titan with a radius of �2; 575 km has a cen-tral pressure of �3:3GPa). However, water ice has sev-eral high density phases that can be present at pressuresof only a few hundred MPa. The temperature and pres-sure conditions for these phases (particularly Ice-II and IceIII) to exist can be reached in the interiors and icy crustsof satellites with radii larger than about 103 km, dependingon their internal structure and temperatures. For these ob-jects estimates of the uncompressed density from interiormodeling are needed to constrain their composition (Lupoand Lewis 1979; Schubert et al. 1986, 2004). Two exam-ple models from Lupo and Lewis, one for a pure water icesphere and the other for a rock fraction of 0.40 are shownin Fig. 3.3 for comparison. Of the Saturnian satellites onlyTitan (very similar to Ganymede and Callisto in bulk proper-

ties) is expected to have significant amounts high pressure icein its interior resulting in model values for its uncompresseddensity of �1; 500 kg m�3 (or about 0.50 rock fraction)(Schubert et al. 1986).

At the other end of the size spectrum, very small bodieswith low internal pressures may sustain a high degree of bulkporosity if the internal stresses do not exceed the materialstrength of their constituents. Such small bodies may (andin fact, do, generally) also have arbitrarily irregular shapes,from similar considerations (Johnson and McGetchin 1973).For icy satellites, the transition from essentially spherical,non-porous objects to irregular, porous bodies occurs atabout a radius of 200 km (central pressures of 1–10 MPa).Densities measured for all icy objects with radii smaller thanthis transition radius, except Phoebe, have densities well be-low 1;000 kg m�3, indicative of high porosity even if theywere made of pure water ice. Phoebe is also likely porousbut with a significantly higher rock fraction, (Johnson andLunine 2005).

Laboratory measurements of low temperature ice poros-ity confirm this general picture (Durham et al. 2005). Thesedata show that ice even at pressures of �140MPa can sus-tain significant porosity .�0:1/ and at lower pressures, below20 MPa, may have porosities of 0.3–0.5. Note that this means


that even though it may not affect their bulk density signifi-cantly, large satellites may have significant porosity in theirouter icy layers, affecting lithospheric strength and thermalproperties (e.g., Castillo-Rogez et al. 2007). The line labeled‘Porous ice’ in Fig 3.3 was calculated using the Durham etal. measurements (Durham et al. 2005), converted to the den-sity of an ice sphere with porosity set by the central pressure.Given the lower pressures (and higher porosities) in the outerlayers of a real body as noted above, pure ice, porous objectscould have densities at or below these levels. The small icySaturnian satellite densities are in general agreement with aprojection from this pure ice model and must have little or norock fraction.

The six ‘mid-sized’ icy satellites in the Saturnian sys-tem (Mimas, Enceladus, Tethys, Dione, Rhea and Iapetus)fall between the extremes discussed above. Their bulk densi-ties are not strongly affected by either porosity or compres-sion effects, and so should reflect fairly closely the densityand compositional mix of their constituents. With accurateCassini measurements of their radii and masses, the errorsin satellite densities are extremely small, and the wide rangeof density values, from Enceladus

�1;600 kg m�3� to Tethys�

991 km m�3� must reflect real and significant variations intheir compositions as reflected in rock/ice ratios. This patterncontrasts sharply with the systematic density variation of thelarge Galilean satellites, where the most distant, Callisto andGanymede, are ice rich and the inner two satellites, Europaand Io, are essentially rocky objects, although tiny Amalthea,interior to Io, has a low density indicative of an ice-rich andporous small body (Anderson et al. 2005).

These compositional variations, as well as the apparentvery ice-rich nature of the Rings and the inner co-orbitaland ring-related moons, pose difficult questions for satelliteformation models, particularly for earlier models which de-pended on early giant planet luminosity to explain the radialdensity gradient in the Galilean satellites.

3.3.2.2 Equilibrium Condensationand Solar Composition

Background

Since the 1970s, most discussions of the composition of theplanets and satellites have been in the context of what is gen-erally referred to as ‘equilibrium condensation’ from a solarnebula with a ‘solar composition’. The foundation for thisapproach was laid by John Lewis and his colleagues usingtheoretical calculations of pressure and temperature in theearly solar nebula to model the expected composition of solidmaterial as a function of heliocentric distance (Grossman1972; Grossman and Larimer 1974; Lewis 1971–1973). The

underlying assumption in this work was that gas and solidsat any point in the solar nebula could be treated as an equi-librium assemblage based on elements being present in thesame as abundances as found in the Sun and primitive mete-orites (‘solar’ or ‘cosmic’ abundance).

It was recognized that the equilibrium condensationmodel undoubtedly oversimplified processes in the real solarnebula. However, it provided an extremely powerful workinghypothesis for understanding major features of what was thenknown of the distribution of different materials in the solarsystem. In particular, it explained the major dichotomy be-tween the rock and metal rich inner terrestrial planets andthe gas rich outer planets, the apparent spectral and compo-sitional gradient in the asteroid belt, and the existence of hy-drated minerals in meteoritic samples from (presumably) theouter asteroid belt. In addition a key feature of the model wasthat beyond the asteroid belt the composition of condensedmaterials should shift abruptly from being rock/metal domi-nated to being a roughly equal mixture of rock and water ice,producing objects with densities of �2;000 kg=m�3. Lewisfurther suggested in his early papers that heating from therock portion of such bodies might produce melting, differen-tiation and volcanic activity, aided possibly by small amountsof ammonia hydrate lowering the melting point.

At the time, few satellite densities were known accuratelyand none had been studied by spacecraft. By 1989, follow-ing the initial reconnaissance of the outer planets by thePioneer and Voyager missions it was clear that ice-rich satel-lites were the norm, with the three largest having almost ex-actly the rock/ice fractions suggested by Lewis, and manyrelatively small icy bodies showed evidence of geologic ac-tivity. More than 30 years after Lewis’ initial papers on thetopic, both ‘solar composition’ ice/rock satellites and cry-ovolcanism remain at the center of discussions of satellitesystem formation.

Complications with the Equilibrium Model

Subsequent work has elaborated on the basic equilibriumcondensation concept and modified it significantly in somerespects. Although the predicted compositions of planetarymaterials from equilibrium condensation seem to be in gen-eral agreement with the bulk properties of planetary objects,at smaller scales there is considerable evidence for strikinglydisequilibrium conditions in the formation of planetary ma-terials. In meteorites, the existence of pre-solar grains andrefractory materials such as CAIs in otherwise ‘primitive’samples provides clear evidence that the nebular gas and re-fractory solids cannot have remained in complete equilibriumat all times (Grossman 1972). Recently returned samplesfrom Comet Wild 2 by the Stardust mission also contain


refractory grains that must have been formed under muchhotter conditions close to the early Sun than the bulk of thevolatile-rich cometary matter (Joswiak et al. 2008; Westphalet al. 2008).

Recent work modeling the solar nebula has also modifiedthe earlier concepts. In current models, gas and particle inter-actions are modeled in more detail, resulting in the potentialtransport of material within the nebula by drag and turbulentprocesses. This may lead for instance to local enhancementin water and other volatiles due to evaporation of inwardlymigrating grains at the ‘snow line’ (Cuzzi et al. 1993; Cuzziand Zahnle 2004; Stevenson and Lunine 1988). Planetesi-mal migration to the planet forming regions can also producevolatile enhancement in the Jupiter and Saturn systems (seeEstrada et al. 2009). Cassini observations of cometary H/Dvalues in the gases in Enceladus’ plumes support the sugges-tion that the Saturnian satellites incorporated cold icy ma-terial from the outer solar nebula (Waite et al. 2009). Thisobservation may imply that low-temperature volatiles thatcondensed in the outer nebula were incorporated into migrat-ing material, and eventually planetesimals that found theirway into the feeding zones of the giant planets. Further-more, it would tend to support the idea of planetesimal de-livery mechanisms (see Section 3.2.3.2) providing the massfor the regular satellites. Whether these or other processescan fully explain compositional constraints from the GalileoProbe Jovian measurements and the new Titan and Enceladusdata is still to be determined.

One of the most significant modifications of the ‘pure’equilibrium condensation scheme was the recognition thatsolar nebula conditions in some regions might not permitequilibrium to be achieved in the age of the solar system.Studies of the kinetics of reactions in refined models of thesolar nebula suggested that, in the outer portions of the so-lar nebula, CO and N2 might be the primary C and N bear-ing species, as opposed to CH4 and NH3, in the warmer,denser regions (Lewis and Prinn 1980). In the satellite form-ing zones around the giant planets, calculations suggest thatthe reduced species should still predominate and produceequilibrium rock/ice ratios (Prinn and Fegley 1981, 1989).

Despite the issues and complications noted above, equi-librium condensation remains the framework for most dis-cussions of outer solar system bodies, particularly with re-spect to expected rock/ice fractions in the materials formingthe satellites. In this modified view of solar abundance equi-librium condensation, the mix of materials available forforming satellites depends very heavily on the amount ofoxygen which is available in the form of water ice. This inturn depends critically on the partition of carbon betweenCO, CH4, and solid carbon (organics and graphite). If sig-nificant carbon is tied up in CO, which does not form solidsunder the conditions around the giant planets, then a signifi-cant fraction of the nebular O is not available to form water

ice in these regions, and condensates would have a higherrock/ice fraction. These effects are quantitatively describedin the following sections, along with their dependence on theassumed solar C and O abundances.

Effects of Carbon in the Gas Phase

Estimates of protosolar C and O abundances have varied sig-nificantly over the last 30 years due to refinements in theinterpretation of solar spectra and structure. These are dis-cussed in detail in Wong et al. (2008). The most recent valuesare those of Grevesse and colleagues (Asplund et al. 2006;Grevesse et al. 2007). The resultant densities of condensedmaterial for these values are compared with earlier solar es-timates in Fig. 3.4 (Wong et al. 2008), which shows the ex-pected material density of condensates as a function of thefraction of total carbon in the form of CO in the nebular gas.Also noted on the plot are the uncompressed densities of var-ious outer planet satellites and Pluto and Triton (as examplesof far outer solar system bodies).

Effects of Carbon in the Solid Phase

The effects of solid carbon on expected condensate and satel-lite density depend on the gas phase chemistry discussed inthe previous section. The densities of potential solid carbonspecies range from kerogens (�1;100 kg m�3) to graphite(�2;000 kg m�3). Since these densities are not greatly differ-ent from the equilibrium (CO-poor) condensate, there is littlechange in the expected density of materials formed in reduc-ing conditions even if essentially all the carbon is in solidform. In contrast, for CO-rich conditions, the primary effectof solid carbon is to reduce the amount of oxygen tied up inCO, resulting in more water ice production and a decreaseof the condensate density. In the extreme case where all thecarbon is in solid form, gas phase carbon is irrelevant andthe resultant condensate would be a mixture of rock, ice andsolid carbon in roughly 40/40/20 proportions with a densityranging from �1;400 to �1;500 kg m�3, depending on thesolid carbon density. Figure 3.5 (Wong et al. 2008) illustratesthe effects of the fraction of carbon in the solid form on con-densate density for a solid carbon density of 1;700 kg m�3.

Discussion of Saturn Satellite Densities

Saturn’s satellites exhibit a significant range of measureddensities (Fig. 3.3), implying a similar variation in their bulkcomposition in terms of rock/ice fraction. The only irregu-lar satellite with an accurately determined density is Phoebe.Several lines of evidence suggest that Phoebe and the other


Den

sity

(K

g m

–3 x

103 )

Phoebe

0%

Pluto, Triton

EnceladusGanymede, Callisto,Titan

Saturn icy satelliteaverage

Tethys

CO / Total C

2.4

2.2

2.0

1.8

1.6

1.4

1.2

1.00 0.2 0.4 0.6 0.8 1

PROTOSOLAR VALUES

Grevesse et al. (2005)

Anders & Grevesse (1989)Cameron (1981)

30% porosity

Fig. 3.4 Calculated uncompressed density of material condensed froma gas of solar composition as a function of the carbon oxidationstate: Calculations for the resultant density for historical estimatesof solar carbon and oxygen abundances (Anders and Grevesse 1989;

Cameron 1981) and current values (Grevesse et al. 2007). Estimates ofthe uncompressed densities of the outer planet icy bodies are indicatedby arrows. From Wong et al. (2008)

irregular satellites are captured bodies which originated inthe outer solar nebula, but not within the circumplane-tary subnebula of Saturn (Nicholson et al. 2008; Turriniet al. 2008, 2009). Phoebe’s small size suggests that it couldbe highly porous, which means that its measured density of�1;600 kg m�3 is probably an underestimate of the densityof its constituents. Its density is in general agreement withexpectations for condensates from a CO-rich outer solar neb-ula, as described above (Johnson and Lunine 2005). Otherlarger bodies in the outer solar system, Pluto, Charon, Tri-ton, and the KBO Eris, share similarly high densities andestimated rock/ice fractions when compression is accountedfor. Lower density KBOs have also been reported, how-ever, suggesting ice-rich compositions and/or high porosi-ties (Stansberry et al. 2006). Thus the general picture of acold outer solar nebula where achievement of full chemicalequilibrium is inhibited, producing rock-ice condensates, isconsistent with the available data. However, the sample ofobjects in this region with accurately measured densities isstill small, and the existence of low density KBOs suggests,at the least, that processing and fractionation of ice and rockin proto-KBO bodies may have played a role in the currentlyobserved compositions.

The regular satellite system of Saturn is dominated byone massive satellite, Titan, in contrast to the Jupiter systemwith its four large Galilean moons. The rest of the icy satel-lites taken together comprise only about 4.5% of the mass ofTitan. Since Titan’s uncompressed density is �1;500 kg m�3this means that most of the condensed material in the regularsatellite system is consistent with near equilibrium condensa-tion in a reducing circumplanetary nebula with a CO fractionbelow 0.2 (see Fig. 3.4).

The smaller ‘icy’ satellites present a more complex prob-lem for any simple satellite formation model in a subnebulawith temperature decreasing with distance from the planet.As illustrated in Fig. 3.3, these satellites have significant dif-ferences in their densities. In addition, taken together theyhave significantly lower rock/ice fractions than Titan, witha mass-weighted composite density of �1;220 kg m�3. Thisis lower than expected from a solar composition mixture, nomatter what the reduction state or phase of carbon in the sub-nebula (Figs. 3.4 and 3.5). The small, very low density, highporosity satellites closer to the ring system and the very icerich composition of the Rings themselves (Chapter 15) areeven more difficult to reconcile with a solar composition con-densate.

Several possibilities to explain the low icy satellite densi-ties and the large variations among the Saturnian mid-sizedsatellites have been discussed (Wong et al. 2008). One pos-sibility is that subsequent to the formation of Titan, the lo-cal Saturn environment inside of Titan’s orbit became en-riched in oxygen or water (Mosqueira and Estrada 2003a).These authors argue that the satellites inside of Titan do notform and have a chance to survive until the inner regions be-come optically thin and cool. The planetary cooling time forSaturn is �105 years (Pollack et al. 1976), which is consider-ably shorter than the time for disk dissipation, e.g., �106–107years by photoevaporation (Shu et al. 1993). Condensed wa-ter may not have been available in the inner regions for thislong, but as the disk cooled enough for water condensation,the inner regions would have been water-enriched due to thesublimation of inward drifting bodies, and subsequent pref-erential loss of silicates. A similar enrichment may not haveoccurred, or have been as relevant for Jupiter because the


2.2

26%

74%

31%

7%

53% 47% 51% 44%

43% 37%

20%

5%

62%

0%

0%

2.0

1.8

1.6

1.4

1.2

1.00 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1

Fraction of carbon in heavy organics (ρ = 1700 Kg m−3)

Gas-phase CarbonCOCH4

Rock+MetalWater Ice

Solid-phase Carbon

Mat

eria

l Den

sity

(kg

m−3

x 1

03 )

Fig. 3.5 Protosolar condensate density versus fraction of carbon in theform of refractory organics: Organic material assumed to have densityof 1;700 kg m�3. The upper line is the case for the rest of the carbon be-ing in the form of CO gas, the lower line shows the result of the gaseouscarbon being in the form of CH3 gas. Models for intermediate oxidation

states are in the shaded zone between the lines. Example condensatecompositions (fraction of ice, rock, and solid carbon) for no solid car-bon, 20% solid carbon and all carbon in the solid state are illustrated bythe inserted pie diagrams. From Wong et al. (2008)

planetary cooling time is much longer (Pollack et al. 1976).Once these satellites form, their survival hinges on their mi-gration times being longer than the disk dissipation time(see Mosqueira and Estrada (2003b) for a discussion of sur-vival mechanisms). Another possibility is that the disk wasenhanced in water through the ablation of large, differenti-ated planetesimals, and fractionated (icy) planetesimal frag-ments that pass through the circumplanetary gas disk (seeSection 3.2.3.2).

In either case, subsequent processes are still required tocreate the wide range in rock/ice fractions found in the satel-lites (possibly including distant Iapetus). Thus the role thatcollisions may have played in the history of the satellites’evolution should be considered. As discussed earlier (seeSection 3.2.3.2), satellites deep within the potential well ofSaturn may have been more subject to disruption or fragmen-tation due to collisions with heliocentric interlopers than theJovian satellites. Hyperbolic collisions might conceivably re-move volatiles from the mantle of a differentiated satelliteand place them on neighboring ones (analogous to the im-pact that may have stripped Mercury’s mantle, Benz et al.(1988)).

Smaller satellites may have been disrupted and then reac-creted. Such a scenario may explain the near total absenceof rock from the innermost satellites (and possibly the rings)if the collisional process leaves behind the remnants of icy

shells of composite bodies whose silicate cores have sincedisappeared (see also Chapter 17). In a similar vein, if theLate Heavy Bombard at �3:9Ga was in fact a solar system-wide event as suggested by the Nice model, there would alsobe a significant probability that smaller satellites in the in-ner system might have been completely disrupted (Charnozet al. 2009) further modifying the distribution of materialsamong the satellites. However, by this time the subnebulawould long have dissipated, meaning there would be no gasaround to assist in migration or circularization of their orbits.While intriguing, none of the above suggestions have yet re-ceived a detailed treatment explaining the current distributionof satellite density and the effects on the dynamical historiesof the satellites, and the collisional mechanisms remain to bequantitatively evaluated.

On the other hand, the explanation for Iapetus’ low den-sity likely depends on whether it formed in-situ near itscurrent location, or was also a product of collisional pro-cesses. The most straightforward explanation is that Iape-tus formed near its present location in a low density, outerSaturnian subnebula that preferentially captured icy materialversus rock through the ablation of disk-crossing planetesi-mal fragments (see Section 3.2.3.2). An alternative scenariowhich is dynamically feasible, but low probability, is theimpact of a Triton-sized differentiated interloper with Titan(Mosqueira and Estrada 2005). Such a collision may produce


a volatile-rich debris disk, some of which may be reaccretedby Titan, or drift inward (due to gas drag) while still leav-ing sufficient material to form an Iapetus-sized object. Theseauthors estimated that its formation time is shorter than thetimescale for a collision with Titan, or the timescale to scatterIapetus to its present orbital location. If Iapetus is success-fully scattered, it would have a large eccentricity (e � 0:7)requiring gas drag to circularize its orbit, while subsequenttidal circularization could explain Titan’s anomalously largeprimordial eccentricity.

The low rock/ice fractions of the bulk of the material inSaturn’s satellites and the similarity in physical propertiesbetween Titan, Ganymede and Callisto are consistent withroughly solar system abundance condensates in reducing ci-crumplanetary environments. In contrast, the higher density,higher rock fraction composition of captured Phoebe andthe larger outer solar system bodies Pluto, Charon, Tritonand several KBOs are indicative of the more oxidizingconditions in outer solar nebula suggested by Lewis andPrinn (Lewis and Prinn 1980). However, the existence ofsmaller icy satellites with large differences in their rock/icefractions and the extreme ice-rich innermost satellites, aswell as evidence for ice-rich KBOs suggest that processessubsequent to formation, including collisional redistributionof material, may have altered their compositions. Thereforethe small satellite densities and compositions may notprovide direct clues to the composition and state of the earlycircumplanetary nebula.

3.3.3 Satellite Structures

3.3.3.1 Background

Following formation, the two factors which dominate a satel-lite’s subsequent evolution and thermal history are its compo-sition (i.e., its rock/ice fraction) and size. The amount of heatproduced is simply a function of the rock fraction in the satel-lite, which carries the radioactive elements, while the degreeto which the satellite is heated is proportional to its radius(i.e., heat production ˛ volume, while cooling ˛ to surfacearea). Lewis noted in his early papers on outer planet ice-rocksatellites that they would be easier to melt and differentiatethan their silicate counterparts simply because their large icecomponent would melt at a much lower temperature. If minorconstituents such as ammonia or salts are present, he pointedout that even lower temperature melting might be achieved.

Simple thermal models for heat conduction in icy satel-lites confirmed that heating from long-lived radioactive iso-topes of U, Th, and K could produce temperatures nearthe melting point of water at relatively shallow depthin large satellites (e.g., Ganymede, Callisto, and Titan)(Consolmagno and Lewis 1977, 1978; Fanale et al. 1977).

The major complicating factor for these satellites is the is-sue of conduction versus convection. If sub-solidus convec-tion begins as the ice viscosity decreases with temperature,their interiors can remain below the melting point, since con-vection is more efficient at removing internal heat than isconduction (Reynolds and Cassen 1979). Because of this un-certainty as well as uncertainties in the satellites’ initial inter-nal temperatures following accretion, thermal models of icysatellites prior to Galileo and Cassini/Huygens usually con-sidered both homogenous and differentiated cases as initialconditions (see Schubert et al. (1986) for a review of meth-ods and results).

3.3.3.2 Observations

In the absence of seismic data from the surface, the bestfirst-order indication of the internal structure of a satelliteis its axial moment-of-inertia, which constrains the degree ofconcentration of mass towards the center. Satellites orbitinggiant planets undergo distortions in their shape and gravityfield due to both spin and tidal forces with the amount ofdistortion depending on the moment of inertia. Hubbard andAnderson (Hubbard and Anderson 1978) showed that space-craft tracking measurements of the J2 term of the sphericalharmonic expansion of a satellite’s gravity field would en-able a useful determination of the moment of inertia. Themajor assumption required is that the body be in hydro-static equilibrium, which fixes the relation between J2 andC22 (the zonal harmonic). Although believed to be a reason-able assumption for large ice-rich objects with warm interi-ors, it may not be valid for small and/or very cold satellites(McKinnon 1997; Nimmo and Matsuyama 2007).

The Celestial Mechanics radio science team on Galileosuccessfully applied this technique to the four Galileanmoons during multiple flybys and determined that Io,Europa and Ganymede all had moments of inertia indicativeof significant differentiation of heavy material (rock andmetal) from water in their interiors, with the number ofdistinct layers and their densities being less well constrained(Anderson et al. 1996a, b, 1997a, b, 1998a, b, 2001a, b;Schubert et al. 2004). The existence of a dynamo magneticfield at Ganymede is also strong evidence for differentiationof this satellite (Kivelson et al. 1996, 1997). Inductionmagnetic fields produced by interaction with the time vari-able Jovian field have also been interpreted as evidence forsalty, electrically conducting subsurface oceans at Europa,Ganymede and Callisto (Khurana et al. 1998; Kivelsonet al. 1999, 2000). However, Callisto, while showing aclear magnetic induction signature, has an inferred momentof inertia suggestive of only partial differentiation. Thus,while its outer layers may have been heated and melted, thedeep interior has apparently not completely differentiated(Schubert et al. 2004).


At Saturn, one of Cassini/Huygens goals is to determinethe interior structures of the satellites where possible. Thishas proved a challenging task for several reasons. The designof the Cassini Orbiter spacecraft makes it very difficult toobtain precision radio tracking data near a satellite simulta-neously with remote sensing observations. This has limiteddedicated radio tracking for gravity purposes to a handfulof flybys of Titan and Rhea. In addition, as noted above,the validity of the hydrostatic assumption has been calledinto question, particularly for the smaller, very cold satellites,which may have sufficiently rigid crusts to prevent relaxationto the hydrostatic state. Therefore, several different lines ofevidence must be brought to bear on the problem of internalstructure and degree of differentiation. The companion vol-ume, Titan from Cassini-Huygens, and Chapter 18 discussthe Cassini results in detail. We provide a brief summary oftheir findings below.

Titan

Based on its similar size and density to Ganymede andCallisto, Titan is expected to be at least partially differen-tiated (e.g., Schubert et al. 2004). Radio tracking data fromCassini show that Titan’s quadrupole gravity field appearsto be non-hydrostatic, making it difficult to infer internalstructure from the simple hydrostatic approximation [Radioscience team paper in preparation]. There is other evidenceof internal heating and a subsurface liquid layer, however.This includes theoretical estimates the loss of methane fromTitan’s atmosphere by photodissociation and hydrogen es-cape (Lunine and Atreya 2008), which implies outgassingof significant quantities of methane during its history (Tobieet al. 2006). The geologically young, eroded surface of themoon also suggests internal heating and activity. Finally, car-tographic comparisons of features observed by the CassiniRadar investigation show that Titan’s ice crust is not rotat-ing synchronously. This result has been interpreted as due todrag forces from atmospheric winds, requiring a liquid layerto decouple the crust from the interior (Lorenz et al. 2008).Thus the weight of the evidence suggests that Titan’s inte-rior has been warm enough for at least partial differentiationand melting. [see Titan after Cassini/Huygens for completediscussion]

Enceladus

Voyager images revealed that tiny Enceladus (R D 250 km,smaller than the largest asteroids) was a geologically ac-tive moon at some point in its history, with heavily cratered

regions adjacent to smooth areas evidently resurfaced bysome process. Its location close to the densest part of thediffuse E-ring and the short estimated lifetime of E-ring par-ticles suggested that geologic activity might be on-going andresponsible for replenishing the E-ring. A number of effortswere made to construct thermal models which might ex-plain this putative activity by raising the internal tempera-tures to close to the water melting point, possibly with thehelp of ‘anti-freeze’ such as small amounts of NH3 (Squyreset al. 1983). A major difficulty with all these models isthe relatively low amount tidal energy available from Ence-ladus’ current eccentricity, forced by a orbital resonancewith Dione. In addition, prior to Cassini Huygens, Ence-ladus’ mass, and therefore density and rock fraction werevery poorly constrained, with most estimates suggesting alow density. Thus the amount of radiogenic heating availablewas highly uncertain.

Cassini discoveries have spectacularly confirmed cur-rent, geyser-like activity on Enceladus and its role supply-ing material to the E-ring (Dougherty et al. 2006; Hansenet al. 2006; Porco et al. 2006; Spencer et al. 2006; Waiteet al. 2006). The temperature and heat flow in the SouthPolar Region has been measured from infrared observation,with estimates of 3–7 GW of power being emitted fromthe fracture system dubbed ‘tiger stripes’ from which theplumes originate (Spencer et al. 2006) (Chapter 21). Themass and density of Enceladus is now measured accuratelyfrom Cassini tracking and radio science data, yielding ahigh density of 1;600 kg m�3, providing potentially higherlevels of radiogenic heating from the rock fraction. How-ever, long-lived radiogenic heating cannot currently producethe observed power levels. Tidal heating at some point inEnceladus’ history is therefore very likely and this com-plicates understanding of its initial thermal and interiorstate.

Although the active plumes in the South Polar Regionand the measured high heat flow values are commonly takenas prima facie evidence that there is sufficient energy avail-able to produce differentiation, there are not yet any mea-surements of gravity structure against which to test thisassumption. Thermal models using the new density valuesshow that early differentiation is possible due to heating fromlong lived radioactive isotopes, and highly probable if Saturnformed early enough for heating for short lived radioactiveisotopes to be significant (Matson et al. 2007; Schubertet al. 2007). However, there is not yet a satisfactory theoryfor ultimate source of the total power required to explain theheat flow from the SPR either directly from global tidal dis-sipation (Porco et al. 2006) or from friction along the frac-tures in an ice shell over a liquid layer (Nimmo et al. 2007)(Chapter 21).


Rhea

Rhea’s surface is relatively bright with regions of heavy tec-tonic fracturing exposing bright ice surface. This impliessome degree of separation of ice and dark rock material in atleast the upper surface layers (many meters to perhaps kilo-meters). Gravitational data from the radio science experimenthowever are consistent with either no, or very limited differ-entiation, although other interpretations of the single flybyradio data are possible (Anderson and Schubert 2007; Iesset al. 2007; Mackenzie et al. 2008). For a full discussion seeChapter 18. Rhea’s density is close to that for Enceladus andalso implies significant radiogenic heating, so a thermal his-tory which prevents full differentiation may be required, sim-ilar to the problem of explaining Callisto’s apparent liquidlayer combined with a derived moment of inertia indicatedlimited differentiation (Anderson et al. 2001b).

Iaptetus

Iapetus’ low density (1;200 kg m�3) and low rock fraction,make it intrinsically difficult to differentiate from radiogenicheating. It also has an extremely non-equilibrium shape forits synchronously locked spin period of almost 80 days, im-plying a thick rigid lithosphere at the current time. Thermalmodels for Iapetus which allow enough dissipation to de-spin the satellite while cooling rapidly enough retain a ‘fos-sil’, oblate shape from an earlier more rapid spin state havebeen developed by Castillo-Rogez et al. (2007). These mod-els require that Iapetus, and therefore the Saturn system, wasformed about 2–5 Myr following the condensation of the firstsolids in the solar nebula (CAIs), consistent with general dis-cussion of rapid giant planet formation in earlier sections.This time frame is also broadly consistent with that neces-sary to explain Callisto’s apparent lack of complete differen-tiation (Barr and Canup 2008). The heating from short livedradioactive isotopes, primarily 26Al, in these models is suffi-cient to warm the deep interior but not to trigger significantdifferentiation in the bulk of the satellite.

Solar wind magnetic field deflections in the vicinity of Ia-petus have been interpreted (Leisner et al. 2008) as evidencefor either significant crustal magnetization or a conductinglayer, such as a liquid ocean with salt content (similar tothose inferred from induction magnetic field signatures inthe icy Galilean satellites). Based on the above discussion,a liquid layer appears highly unlikely, both from basic en-ergy considerations and from the observed non-equilibriumshape, implying very low levels of subsurface heating in thelast three to four Gyr.

3.3.4 Satellite Geological History

3.3.4.1 Crater ages

Determining absolute ages for outer planet satellite surfacesis greatly complicated by the on-going debate on the natureof the impactor population and the appropriate impact fluxas a function of time and impactor size (Porco et al. 2005a;Zahnle et al. 2003). In the absence of isotopic age determina-tions of samples, such as are available to calibrate the lunarcrater record, satellites ages are dependent on large extrap-olations from current observations and uncertainties in thehistory of the impact flux. Very heavily cratered surfaces,at or near ‘crater saturation’, are generally agreed to be ex-tremely ancient, dating back 3–4 Ga or more when essen-tially all models predict much higher fluxes of impacting ob-jects than in the current solar system. Likewise, uncrateredsurfaces or very low crater densities certainly imply geologi-cally young surfaces (years to millions of years old). Iapetus,discussed in more detail in the next section, is the best exam-ple of an ancient, highly cratered surface in the system, whileEnceladus’ South Polar Region, with its on-going plume ac-tivity is an example of an archetypically geologically youngregion, as is most of Titan’s very sparsely cratered surface,which shows ample evidence for recent fluvial, Aeolian, tec-tonic, and possibly cryovolcanic modification. For most ofthe other icy satellites, while relative chronologies can be es-tablished based on crater statistics, estimates of absolute agesfor different assumptions of impact flux can vary by factorsof 2–10 (see Chapter 19 for a full discussion).

3.3.4.2 Iapetus’ Formation Time

The Nice model for early solar system dynamical evolutiondiscussed earlier provides a possible path to linking the sam-ple based chronology of the inner solar system, including theMoon, to at least some portions of the history of the Saturnsystem. If the multiple large impact basins on Iapetus formedcontemporaneously with the lunar basins at 3.9 Ga as a re-sult of a solar system wide event associated with the LHB,it provides a tie point for evaluating other age estimates inthe system. Such an age is basically in agreement with theinterpretation of crater statistics for either of the two con-tending models within uncertainties at present. It is also con-sistent with the chronology proposed by Castillo-Rogez et al.for Iapetus, whose thermal models imply a rapidly coolinglithosphere rigid enough to retain both the satellite’s oblateshape and the record of large impact basins by the time ofthe LHB (Johnson et al. 2007). Figures 3.6 and 3.7 illustrate


10−2 10−3 10−4 10−5 10−6 10−7 10−8 10−9 10−10 10−11 10−12

lapetus forms

Normal CAIs

Pb – Pb(Ma)

4,575 4,570 4,565 4,560 4,555 4,550

26Al/27Al Anomalous CAIs

Eucrites

Chondrules

Angrite

Angrites

Phosphates

? ?

Fig. 3.6 Short lived radioisotope chronometry compared with ab-solute ages from Pb-Pb dating for various early chondritic and ig-neous (Eurcrite) meteorite samples: Modified after (McKeegan andDavies 2007), reprinted by permission from Elsevier, and compared

with the estimated time of formation of Iapetus from Castillo-Rogezet al. (2007). The upper scale gives the inferred initial 26Al=27Al valuetied to absolute Pb/Pb ages (bottom scale) using the dating of CAI for-mation of 4:567:2 ˙ 0:6Ga (Amelin et al. 2002)

Initiation of Solar System FormationFirst solids, CAIs ~ 1 Myr later

Iapetus and the Lunar Late Heavy BombardmentSimplified Chronology

Iapetus despins to synchronous rotation: model range from (Castillo-Rogez et al., 2007)Iapetus forms:

~3.5-4.6 Myr after CAIs

35003600370038003900400041004200430044004500

Age, Ma

Iapetus and the Lunar Late Heavy BombardmentSimplified Chronology

Lunar Late Heavy Bombardment

Earth-Moon forms:~40 Myr after CAIs

Iapetus' lithosphere retains large impact basins(Castillo-Rogez et al., 2007)

Fig. 3.7 Solar system chronology: Shows the time ranges for possi-ble de-spinning Iapetus and for Iapetus’ lithosphere to become rigidenough to retain the record of large impact basins and compares this tothe time of the LHB as dated from lunar samples. From (Castillo-Rogezet al. 2007)

how this proposed chronology relates Iapetus’ formation andearly thermal evolution to the formation of CAIs, meteoriteparent bodies and the lunar record.

3.4 Summary and Issues for the Future

This chapter summarizes how the Cassini Huygens resultshave constrained and added to our understanding of the for-mation and evolution of the Saturn system, its rings and satel-

lites. As the Cassini orbiter continues to make new observa-tions in the extended phase of its mission, new discoveriesare likely in the future related to Saturn’s internal structure,the structure of the icy satellites and Titan, the history ofTitan’s surface and atmosphere, the evolution of the rings,and the on-going plume activity on intriguing Enceladus.NASA and ESA are currently in the process of developingplans for future outer planet exploration which will undoubt-edly include following up on many of Cassini Huygens ex-citing discoveries with future missions to the Saturn system.

Acknowledgements The authors wish to acknowledge Jack J. Lissauer(NASA Ames Research Center) for many useful discussions and com-ments on an earlier draft of this work.

A portion of this work (TVJ) has been conducted at the Jet Propul-sion Laboratory, California Institute of Technology, under a contractwith the National Aeronautics and Space Administration.

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