Planet Formation

The favoured model of forming the cohort of currently discovered planets is by core accretion (Laughlin, Bodenheimer & Adams 2004) this is backed up by both observations (Fischer & Valenti 2005) and theory (Rafikov 2005). Core accretion forms planets by the build-up of particles in the stellar disk; the disk being leftover from the formation of the parent star. Initially the small dust grains in the disk begin to coagulate until they form larger structures whose gravity can affect their

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Fig. 6.10. The plot shows median MJUP versus Mstal0. The solid line of binned data suggests a relationship between host mass and planet mass, however, the dashed line of all planets above 1 MJUP sin i does not support this conclusion for a more 'complete' sample. The error bars are from y/ (number) statistics and are only indicative.

Fig. 6.10. The plot shows median MJUP versus Mstal0. The solid line of binned data suggests a relationship between host mass and planet mass, however, the dashed line of all planets above 1 MJUP sin i does not support this conclusion for a more 'complete' sample. The error bars are from y/ (number) statistics and are only indicative.

surroundings. The simplest method of planet formation occurs when a core grows by the accretion of a number of planetesimals at a fixed orbital radius from the star (Safranov 1969). Initially this core will be surrounded by a near-hydrostatic gaseous envelope, with the majority of the luminosity being generated by the continuous planetesimal accretion. The growth of the core will continue until a critical core mass is reached, at which point there is no stable core-envelope solution and the planetary envelope will collapse onto the protoplanetary core (Mizuno 1980; Pollack et al. 1996). The mass of the accreted planetesimals and accreted gas is almost equally distributed and as the envelope contracts the gas accretion rate increases, which increases the energy losses through radiation giving rise to a runaway accretion process. This formation scenario describes the types of planets discovered thus far but the picture is slightly different for terrestrial planets, where the core does not grow large enough for the surrounding gas to collapse onto it, and hence a significant planetary atmosphere must be generated by different processes like volcanism.

One of the long standing problems with the core accretion model relates to the timescale of formation of the planetary embryos. To generate a core massive enough for the surrounding disk gas to collapse onto requires timescales on the order of 107 years (Hayashi, Nakazawa & Adachi 1977), however the observational upper limits for the lifetimes of disks forming planets is in the range 105 -107 years. Thus the planetesimals will have had insufficient time to accrete enough gas to create the gas giant planets we observe today. More recently models of formation and the environment of the disk can somewhat account for this discrepancy. Alibert et al. (2005a) has shown that by including disk evolution, migration and gap formation into their models, the formation timescales for planetesimal cores become more compatible with typical disk lifetimes. This speed-up is mostly due to migration, which ensures there is no depletion of the feeding zones of forming planets, a problem for in situ models. They have also shown that these models are consistent with the wealth of information garnered on Jupiter and Saturn on their core masses and total amount of heavy elements in their envelopes (Alibert et al. 2005). However, another problem arises due to the fast migration timescales of small planetesimal cores.

An alternative model for planetary formation is the disk instability mechanism which is based on the formation and survival of self-gravitating clumps of gas and dust in a marginally gravitationally unstable protoplanetary disk (Boss 1997; reviewed by Durisen et al. 2007). In order for a disk instability to succeed, the disk must be able to cool its midplane as the clumps form allowing them to continue to contract to higher densities. The clumps must be able to survive indefinitely in the face of Keplerian shear, tidal forces and internal thermal pressure. Different resolution and sized calculations together with different boundary conditions and treatments of inner disk heating lead different authors (e.g. Boss 2007, Boley et al. 2007, Mayer et al. 2007) to almost rule out or even promote the disk instability mechanism. Thus further work is required before the importance of the disk instability formation mechanism can be established.

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