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Empirical' models (pressure vs. density) of Uranus and Neptune interiors constrained by the gravitational coefficients J2, J4, the planetary radii and masses, and Voyager solid-body rotation periods are presented. The empirical pressure-density profiles are then interpreted in terms of physical equations of state of hydrogen, helium, ice (H2O), and rock (SiO2) to test the physical plausibility of the models. The compositions of Uranus and Neptune are found to be similar with somewhat different distributions of the high-Z material. The big difference between the two planets is that Neptune requires a non-solar envelope while Uranus is best matched with a solar composition envelope. Our analysis suggests that the heavier elements in both Uranus' and Neptune's interior might increase gradually towards the planetary centers. Indeed it is possible to fit the gravitational moments without sharp compositional transitions. * Retiree

Numerical simulations, based on the core-nucleated accretion model, are presented for the formation of Jupiter at 5.2 AU in 3 primordial disks with three different assumed values of the surface density of solid particles. The grain opacities in the envelope of the protoplanet are computed using a detailed model that includes settling and coagulation of grains and that incorporates a recalculation of the grain size distribution at each point in time and space. We generally find lower opacities than the 2% of interstellar values used in previous calculations [Hubickyj, O., Bodenheimer, P., Lissauer, J. J., 2005. Icarus 179, 415--431; Lissauer, J. J., Hubickyj, O., D'Angelo, G., Bodenheimer, P., 2009. Icarus 199, 338-350]. These lower opacities result in more rapid heat loss from and more rapid contraction of the protoplanetary envelope. For a given surface density of solids, the new calculations result in a substantial speedup in formation time as compared with those previous calculations. Formation times are calculated to be 1.0, 1.9, and 4.0 Myr, and solid core masses are found to be 16.8, 8.9, and 4.7 Earth masses, for solid surface densities, sigma, of 10, 6, and 4 grams per squared centimeter, respectively. For sigma=10 and sigma=6 g/cm^2, respectively, these formation times are reduced by more than 50% and more than 80% compared with those in a previously published calculation with the old approximation to the opacity

In a protoplanetary disk, the inner edge of the region where the temperature falls below the condensation temperature of water is referred to as the 'snow line'. Outside the snow line, water ice increases the surface density of solids by a factor of 4. The mass of the fastest growing planetesimal (the 'isolation mass') scales as the surface density to the 3/2 power. It is thought that ice-enhanced surface densities are required to make the cores of the gas giants (Jupiter and Saturn) before the disk gas dissipates. Observations of the Solar System's asteroid belt suggest that the snow line occurred near 2.7 AU. In this paper we revisit the theoretical determination of the snow line. In a minimum-mass disk characterized by conventional opacities and a mass accretion rate of 10^-8 solar masses per year, the snow line lies at 1.6-1.8 AU, just past the orbit of Mars. The minimum-mass disk, with a mass of 0.02 solar, has a life time of 2 million years with the assumed accretion rate. Moving the snow line past 2.7 AU requires that we increase the disk opacity, accretion rate, and/or disk mass by factors ranging up to an order of magnitude above our assumed baseline values.Comment: Accepted for publication in ApJ, 9 pages, 4 figure

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The large number of detected giant exoplanets offers the opportunity to improve our understanding of the formation mechanism, evolution, and interior structure of gas giant planets. The two main models for giant planet formation are core accretion and disk instability. There are substantial differences between these formation models, including formation timescale, favorable formation location, ideal disk properties for planetary formation, early evolution, planetary composition, etc. First, we summarize the two models including their substantial differences, advantages, and disadvantages, and suggest how theoretical models should be connected to available (and future) data. We next summarize current knowledge of the internal structures of solar- and extrasolar- giant planets. Finally, we suggest the next steps to be taken in giant planet exploration.Comment: Accepted for publication as a chapter in Protostars and Planets VI, to be published in 2014 by University of Arizona Pres

We have computed the size distribution of silicate grains in the outer radiative region of the envelope of a protoplanet evolving according to the scenario of Pollack et al. (1996). Our computation includes grain growth due to Brownian motion and overtake of smaller grains by larger ones. We also include the input of new grains due to the breakup of planetesimals in the atmosphere. We follow the procedure of Podolak (2003), but have speeded it up significantly. This allows us to test the sensitivity of the code to various parameters. We have also made a more careful estimate of the resulting grain opacity. We find that the grain opacity is of the order of 10 −2 cm 2 g −1 throughout most of the outer radiative zone as Hubickyj et al. (2005) assumed for their low opacity case, but near the outer edge of the envelope, the opacity can increase to ∼ 1 cm 2 g −1 . We discuss the effect of this on the evolution of the models.

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