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A long-popular model for producing Ganymede's bright terrain involves flooding of low-lying graben with liquid water, slush, or warm, soft ice. The model suffers from major problems, however, including the absence of obvious near-surface heat sources, the negative buoyancy of liquid water, and the lack of a mechanism for confining the flows to graben floors. We present new models for cryovolcanic resurfacing to overcome these difficulties. Tidal heating within an ancient Laplace-like orbital resonance (Showman and Malhotra 1997, Icarus 127, 93; Showman et al., 1997, Icarus 129, 367) provides a plausible heat source and could allow partial melting to occur as shallow as 5–10 km depth. Our favored mechanism for delivering this water to the surface invokes the fact that topography—such as a global set of graben—causes subsurface pressure gradients that can pump water or slush upward onto the floors of topographic lows (graben) despite the negative buoyancy of the liquid. These eruptions can occur only within the topographic lows; furthermore, as the low areas become full, the pressure gradients disappear and the resurfacing ceases. This provides an explanation for the observed straight dark-bright terrain boundaries: water cannot overflow the graben, so resurfacing rarely embays craters or other rough topography. Pure liquid water can be pumped to the surface from only 5–10 km depth, but macroscopic bodies of slush ascending within fractures can reach the surface from much greater depths due to the smaller negative buoyancy of slush. A challenge for these models is the short predicted gravitational relaxation timescale of topographic features at high heat flows; the resurfacing must occur before the graben topography disappears. We also evaluate alternate resurfacing mechanisms, such as pumping of liquid water to the surface by thermal expansion stresses and buoyant rise of water through a silicate-contaminated crust that is denser than liquid water, and conclude that they are unlikely to explain Ganymede's bright terrain.

For a satellite to survive in the disk the time scale of satellite migration must be longer than the time scale for gas dissipation. For large satellites (∼1000 km) migration is dominated by the gas tidal torque. We consider the possibility that the redistribution of gas in the disk due to the tidal torque of a satellite with mass larger than the inviscid critical mass causes the satellite to stall and open a gap (W.R. Ward, 1997, Icarus 26, 261–281). We adapt the inviscid critical mass criterion to include gas drag, and m-dependent nonlocal deposition of angular momentum. We find that such a model holds promise of explaining the survival of satellites in the subnebula, the mass versus distance relationship apparent in the saturnian and uranian satellite systems, the concentration of mass in Titan, and the observation that the satellites of Jupiter get rockier closer to the planet whereas those of Saturn become increasingly icy. It is also possible that either weak turbulence (close to the planet) or gap-opening satellite tidal torque removes gas on a similar time scale (104–105 years) as the orbital decay time of midsized (200–700 km) regular satellites forming in the inner disk (inside the centrifugal radius (I. Mosqueira and P.R. Estrada, 2003, Icarus, this issue)). We argue that Saturn’s satellite system bridges the gap between those of Jupiter and Uranus by combining the formation of a Galilean-sized satellite in a gas optically thick subnebula with a strong temperature gradient, and the formation of smaller satellites, closer to the planet, in a disk with gas optical depth ≲1, and a weak temperature gradient. Using an optically thick inner disk (given gaseous opacity), and an extended, quiescent, optically thin outer disk, we show that there are regions of the disk of small net tidal torque (even zero) where satellites (Iapetus-sized or larger) may stall far from the planet. For our model these outer regions of small net tidal torque correspond roughly to the locations of Callisto and Iapetus. Though the precise location depends on the (unknown) size of the transition region between the inner and outer disks, the result that Saturn’s is found much farther out (at ∼3rcS, where rcS is Saturn’s centrifugal radius) than Jupiter’s (at ∼ 2rcJ, where rcJ is Jupiter’s centrifugal radius) is mostly due to Saturn’s less massive outer disk and larger Hill radius. However, despite the large separation between Ganymede and Callisto and Titan and Iapetus, the long formation and migration time scales for Callisto and Iapetus (I. Mosqueira and P.R. Estrada, 2003, Icarus, this issue) makes it possible (depending on the details of the damping of acoustic waves) that the tidal torque of Ganymede and Titan clears the gas disk out to their location, thus stranding Callisto and Iapetus far from the planet. Either way, our model provides an explanation for the presence of regular satellites outside the centrifugal radii of Jupiter and Saturn, and the absence of such a satellite for Uranus.

We model the subnebulae of Jupiter and Saturn wherein satellite accretion took place. We expect a giant planet subnebula to be composed of an optically thick (given gaseous opacity) inner region inside of the planet's centrifugal radius (located at r(sub c, sup J) = l5R(sub J) for Jupiter and r(sub c, sup S) = 22R(sub S) for Saturn), and an optically thin, extended outer disk out to a fraction of the planet's Roche lobe, which we choose to be R(sub roche)/5 (located at approximately 150R(sub J) near the inner irregular satellites for Jupiter, and approximately 200R(sub S) near Phoebe for Saturn). This places Titan and Ganymede in the inner disk, Callisto and Iapetus in the outer disk, and Hyperion in the transition region. The inner disk is the leftover of the gas accreted by the protoplanet. The outer disk results from the solar torque on nebula gas flowing into the protoplanet during the time of giant planet gap opening. For the sake of specificity, we use a cosmic mixture 'minimum mass' model to constrain the gas densities of the inner disks of Jupiter and Saturn (and also Uranus). For the total mass of the outer disk we use the simple scaling M(sub disk) = M(sub P)tau(sub gap)/tau(sub acc), where M(sub P) is the mass of the giant planet, tau(sub gap) is the gap opening timescale, and tau(sub acc) is the giant planet accretion time. This gives a total outer disk mass of approximately 100M(sub Callisto) for Jupiter and possibly approximately 200M(sub Iapetus) for Saturn (which contain enough condensables to form Callisto and Iapetus respectively). Our model has Ganymede at a subnebula temperature of approximately 250 K and Titan at approximately 100 K. The outer disks of Jupiter and Saturn have constant temperatures of 130 K and 90 K respectively.

An explanation of the dynamical mechanism for apse-alignment of the eccentric Uranian rings is necessary before observations can be used to determine properties such as ring masses, particle sizes, and elasticities. The leading model relies on the ring self-gravity to accomplish this task, yet it yields equilibrium masses which are not in accord with Voyager radio measurements. We explore possible solutions such that the self-gravity and the collisional terms are both involved in the process of apse-alignment. We consider limits that correspond to a hot and a cold ring, and show that pressure terms may play a significant role in the equilibrium conditions for the narrow Uranian rings. In the cold ring case, where the scale height of the ring near periapse is comparable to the ring particle size, we introduce a new pressure correction pertaining to a region of the ring where the particles are locked in their relative positions and jammed against their neighbors, and the velocity dispersion is so low that the collisions are nearly elastic. In this case, we find a solution such that the ring self-gravity maintains apse-alignment against both differential precession (m = 1 mode) and the fluid pressure. We apply this model to the Uranian alpha ring, and show that, compared to the previous self-gravity model, the mass estimate for this ring increases by an order of magnitude. In the case of a hot ring, where the scale height can reach a value as much as fifty times larger than a particle size, we find velocity dispersion profiles that result in pressure forces which act in such a way as to alter the ring equilibrium conditions, again leading to a ring mass increase of an order of magnitude; however, such a velocity dispersion profile would require a different mechanism than is currently envisioned for establishing heating/cooling balance in a finite-sized, inelastic particle ring. Finally, we introduce an important correction to the model of Chiang and Goldreich.

We investigate the global dynamics of the uranian rings using a modified 2-D smoothed particle hydrodynamic code combined with a 2-D tree code used to compute the particle-to-particle gravitational interactions. This code includes epicyclic fluid motion, nonaxisymmetric flow, local and nonlocal shear viscosity, self-consistent scale height evolution, ring–satellite gravitational interaction and coevolution, and ring self-gravity. To follow the scale height of each SPH particle we solve the vertical momentum equation for the flow using a Runge-Kutta scheme with a second order polynomial fit to the vertical behavior of the fluid pressure (N. Borderies et al. , 1985, Icarus 63 , 406–420). The behavior of the fluid viscosity is obtained from I. Mosqueira (1996, Ph.D. thesis, Cornell University, Ithaca, NY), who found good agreement between a low optical depth extension of the high optical depth nonlocal viscosity model of Borderies et al. (1985) with the results of a local patch-code ring simulation. Our present viscosity model includes a further correction which accounts for the epicyclic limit to the mean free path (P. Goldreich and S. Tremaine, 1978, Icarus 34 , 227–239). This treatment covers both the high and the low ring density regimes so long as the fluid treatment remains valid. All energy source and sink terms in the presence of a nonzero fluid velocity divergence are self-consistently considered, including terms not considered in prior studies of rings. Furthermore, a new method has been tested and used to remove the noise contribution to the viscous momentum flux which occurs in traditional treatments of the viscosity within the SPH framework. In the present context this correction is needed to assure a physical behavior for the ring. Even within a 2-D framework the uranian rings are so narrow compared to their semimajor axes that radial resolution requires too many particles given our present computer resources. To address this issue we have developed a physical scaling that reduces the semimajor axis of the ring while preserving its width and, we believe, retains the relevant global satellite–ring dynamics. In the case of ring resolution on the order of kilometers, with a value of the scaling parameter that reduces the ring's semimajor axis by a factor of 20, our scaling allows for savings between a factor of 500 in the case of synodic time scales to a factor of 10,000 for viscous time scales. We expect this scaling to figure significantly in numerical studies of the global dynamics of planetary rings in the future. In the present study we use this scaling to follow the evolution of an epsilon-like ring in the presence of an m =2 mode (but no satellites) with and without self-gravity. In these runs, chosen as an initial demonstration of the usefulness and applicability of our scaling, the ring can be seen to evolve in the presence of shear.

Using a modifiedN-body code to include periodic boundary conditions in a perturbed shear flow, we investigate the role of viscosity on the dynamics of perturbed rings with optical depth τ ∼ 1. In particular, we are concerned with rings such thatq=a(de/da) ≠ 0, whereais the semi-major axis andeis the eccentricity. We confirm the possibility that, for a sufficiently perturbed ring, the angular momentum luminosity may reverse direction with respect to the unperturbed ring (Borderies, N., P. Goldreich, and S. Tremaine 1983a.Icarus55, 124–132). We use observationally constrained parameters for the δ and ϵ uranian rings, as well as the outer portion of Saturn'sBring. We find that understanding the effects of viscosity for the uranian rings requires that both local and non-local transport terms be considered if the coefficient of restitution experimentally obtained by Bridgeset al.(Bridges, F. G., A. Hatzes, and D. N. C. Lin 1984.Nature309, 333–335.) is appropriate for ring particles. We also find evidence that the criterion for viscous overstability is satisfied in the case of high optical depth rings, as originally proposed by Borderieset al.(Borderies, N., P. Goldreich, and S. Tremaine 1985.Icarus63,406–420.), making viscous overstability a leading candidate mechanism to explain the non-axisymmetric structure present in the outer portion of Saturn'sBring. To better understand our patch-code results we extend a non-local and incompressible fluid model used by Borderieset al.for dense rings. We incorporate local and non-local transport terms as well as compressibility, while retaining the same number of arbitrary model parameters.

The 11 July 1992 occultation of Uranus by U103 (mK= 10.1) was observed using the 5-m Hale telescope on Palomar Mountain in California, the 4-m telescope at Cerro Tololo Inter-American Observatory (CTIO) in Chile, and the nearby 2.2-m telescope at European Southern Observatory (ESO). All observations were taken using InSb aperture photometers andK(λ ∼ 2.2 μm) filters and provided high-quality profiles of the uranian rings. The ESO and Palomar ring event times have been incorporated into a comprehensive orbit model for the elliptical rings using all available Earth-based and Voyager occultation data. The kinematical model was used to determine the radius scale and geometry of the U103 occultation, as well as to provide absolute timing calibration for the CTIO light curve, which was reconstructed from two-dimensional digital scans of a high-speed strip chart. All nine classical rings were clearly observed during both ingress and egress. In addition, a sharp feature appeared in the CTIO ingress light curve with an orbital radius of 50,026.89 ± 0.39 km, very close to the mean radius of the λ ring of 50,026.3 ± 0.6 km found by M. R. Showalter (Science267, 490–493, 1995) from an analysis of Voyager 2 images. The CTIO feature is well above the noise: its equivalent width ofE= 0.32 ± 0.14 km corresponds to a 5–σ detection. It is the only such signature in the entire CTIO lightcurve, other than the previously known rings. The characteristic ring-like shape of the profile, its high SNR, and its radial location all support the identification of this feature as an Earth-based detection of the λ ring. At the same time, if the ring were azimuthally homogeneous withE≳ 0.32 km, it would have been clearly visible in the CTIO egress and Palomar light curves, but no such features were found. The λ ring is thus azimuthally clumpy, resembling to some degree Neptune's arc-like Adams ring. At the nearby ESO station, a suggestive ingress feature appeared with a radius of 50,023.01 ± 0.28 km, but the data are too noisy to permit a positive identification. The orbital characteristics of the λ ring and its azimuthal brightness profile cannot be uniquely extrapolated from the Voyager encounter epoch to the time of the U103 occultation because the mean motion and orbital radius of the ring are too uncertain. However, additional detections in archival Earth-based occultation data or from future observations could allow a much better determination of the dynamics of this dusty and clumpy ring.

Data from eight stellar occultations by Neptune between 1984 and 1988 are analyzed to set limits on the optical depths of the continuous Adams and Le Verrier Rings, and to search for previously unidentified occultations by the arcs in the Adams Ring. We employ the approach introduced by Sicardy et al . (1991), Icarus 89, 220-243) and convert the raw lightcurves to radial profiles of equivalent width E ( r ), the fraction of incident starlight removed by ring material at normal incidence integrated over a specified radial range, or "window." In order to optimize the search for narrow rings, the window length is fixed at 15 km. Radii in Neptune's equatorial plane are calculated using the Neptune pole direction determined by Jacobson et al . (1990), AIAA/AAS Astrodynamics Conference , pp. 157-167). Using data from the occultation of 20 August 1985, we set a 3 σ upper limit of 75 m on the equivalent width of the Adams Ring at a longitude 120° away from tile arc region, near the location of minimum ring brightness in the Voyager images (Showalter and Cuzzi 1994, in preparation). This limit is for a maximum ring width of 15 km; for a maximum ring width of 50 km, the corresponding upper limit is E ≤ 210 m. Combining the results from this data set with similar upper limits obtained by Sicardy et al . (1991) with a 50-km window size, and assuming that the Adams Ring is azimuthally homogeneous away from the arc region, we find that the equivalent width of this ring is ≤90 m at the 90% confidence level, or ≤ 5% of the measured equivalent width of the Egalite arc. A comparison of this result with the relative brightnesses of ring and arcs observed at low phase angles in Voyager images (Showalter and Cuzzi 1994) suggests either that the particles in the Adams Ring have higher albedos than those confined within the arcs, or that the width of the continuous ring significantly exceeds 50 km. The 20 August 1985 lightcurve has also yielded the first tentative earth-based detection of the Le Verrier Ring, with an overall radial width of 135 km and an equivalent width of 520 ± 55 m at λ2.2 μm which are similar to results obtained by the Voyager Photopolarimeter occultation experiment at λ0.27 μm (Horn et al . 1991, Geophys. Res. Lett. 17, 1745-1748). The ratio E PPS / E 2.2 μm = 1.3 ± 0.3. The mean radii of both occultation profiles—53055 and 53125 km, respectively, projected into the ring plane determined by Porco et al . ( Neptune and Triton Conference, Tucson, AZ, 1992 ) and Showalter and Cuzzi (1992), Bull. Am. Astron. Soc. 24, 1029)—are significantly less than the image-derived radius of 53200 ± 20 km (Porco et al . 1992). It is conceivable that this discrepancy could be due to an unrecognized inclination and/or eccentricity. A comparison of our putative detection of the Le Verrier Ring with upper limits set by Sicady et al . (1991) suggests that this ring must be azimuthally inhomogeneous on scales of ∼10°, if our detection is real. A feature in the vicinity of the Adams Ring at the predicted longitude of the Liberte arc was seen in data obtained on 9 July 1988, but a discrepancy of 270 km between the calculated radius and the known semi-major axis of the ring seems to rule out an identification with the arc. Attempts to fit a precessing ellipse model to all available occultation data for the Adams Ring which will also accommodate this observation are in vain. The remaining occultation-derived radii for this ring are found to be consistent with the resonant-perturbation model of Porco (1991, Science 253, 995-1001), and suggest strongly that the Adams Ring lies in, or very close to, Neptune's equatorial plane.

The origin of the regular satellites ties directly to planetary formation in that the satellites form in gas and dust disks around the giant planets and may be viewed as mini-solar systems, involving a number of closely related underlying physical processes. The regular satellites of Jupiter and Saturn share a number of remarkable similarities that taken together make a compelling case for a deep-seated order and structure governing their origin. Furthermore, the similarities in the mass ratio of the largest satellites to their primaries, the specific angular momenta, and the bulk compositions of the two satellite systems are significant and in need of explanation. Yet, the differences are also striking. We advance a common framework for the origin of the regular satellites of Jupiter and Saturn and discuss the accretion of satellites in gaseous, circumplanetary disks. Following giant planet formation, planetesimals in the planet's feeding zone undergo a brief period of intense collisional grinding. Mass delivery to the circumplanetary disk via ablation of planetesimal fragments has implications for a host of satellite observations, tying the history of planetesimals to that of satellitesimals and ultimately that of the satellites themselves. By contrast, irregular satellites are objects captured during the final stages of planetary formation or the early evolution of the Solar System; their distinct origin is reflected in their physical properties, which has implications for the subsequent evolution of the satellites systems., 19 pages, 1 figure. The chapter is to be in the Space Science Reviews, Europlanet vol. on Icy Satellites

Here we show that Iapetus can serve to discriminate between satellite formation models. Its accretion history can be understood in terms of a two-component gaseous subnebula, with a relatively dense inner region, and an extended tail out to the location of the irregular satellites, as in the SEMM model of Mosqueira and Estrada (2003a,b). Following giant planet formation, planetesimals in the feeding zone of Jupiter and Saturn become dynamically excited, and undergo a collisional cascade. Ablation and capture of planetesimal fragments crossing the gaseous circumplanetary disks delivers enough collisional rubble to account for the mass budgets of the regular satellites of Jupiter and Saturn. This process can result in rock/ice fractionation provided the make up of the population of disk crossers is non-homogeneous, thus offering a natural explanation for the marked compositional differences between outer solar nebula objects and those that accreted in the subnebulae of the giant planets. Consequently, our model leads to an enhancement of the ice content of Iapetus, and to a lesser degree those of Ganymede, Titan and Callisto, and accounts for the (non-stochastic) compositions of these large, low-porosity outer regular satellites of Jupiter and Saturn. (abridged), Comment: 33 pages, 7 figures, 2 tables, Accepted for publication to Icarus

Assuming that an unknown mechanism (e.g., gas turbulence) removes most of the subnebula gas disk in a timescale shorter than that for satellite formation, we develop a model for the formation of regular (and possibly at least some of the irregular) satellites around giant planets in a gas-poor environment. In this model, which follows along the lines of the work of Safronov et al. (1986), heliocentric planetesimals collide within the planet's Hill sphere and generate a circumplanetary disk of prograde and retrograde satellitesimals extending as far out as $\sim R_H/2$. At first, the net angular momentum of this proto-satellite swarm is small, and collisions among satellitesimals leads to loss of mass from the outer disk, and delivers mass to the inner disk (where regular satellites form) in a timescale $\lesssim 10^5$ years. This mass loss may be offset by continued collisional capture of sufficiently small $< 1$ km interlopers resulting from the disruption of planetesimals in the feeding zone of the giant planet. As the planet's feeding zone is cleared in a timescale $\lesssim 10^5$ years, enough angular momentum may be delivered to the proto-satellite swarm to account for the angular momentum of the regular satellites of Jupiter and Saturn.(abridged), 45 pages, 11 figures, 3 appendices, uses rgfmacro.tex, accepted for publication to Icarus