Planetesimal formation

Many of the observations described in the previous section can be explained by the planetesimal theory, in which rocky planets form by the aggregation of huge numbers of asteroid-sized building blocks called planetesimals (Safronov 1969). This 

Most young stars with optically thick disks appear to be accreting gas, which suggests that gas is continually flowing inwards through the disk and onto the star (Calvet et al. 2004). This inward viscous accretion is probably caused by turbulence in the differentially rotating disk. The existence of turbulence and the strength of the turbulent eddies can have a substantial impact on planetesimal formation, as we will see in Section 10.2.2. 

The turbulent viscosity v is often parameterized using v = acsH, where it is assumed that turbulent eddies will be smaller than the scale height H of the gas and will rotate more slowly than the sound speed cs. The observed rates of gas accretion onto young stars suggest that 10-2 a 10-4 (Hueso Guillot 2005). Most studies assume that a is a constant, although in practice it probably varies with time and location. 

Laboratory experiments show that micrometer-sized dust grains, like those in proto-planetary disks, readily stick together during collisions, forming porous aggregates up to 1 cm in size (Poppe et al. 2000 Marshall Cuzzi 2001 Krause Blum 

Compact objects like chondrules rebound during low-speed collisions and undergo some fragmentation at speeds 20 m s-1 (Ueda etal. 2001). 

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The Solar System, in comparison, offers two lines of evidence to constrain the timescale for the lifetime of the proto-solar nebula and the epoch of planetesimal formation. On one side, relatively unaltered chondritic components preserve traces of their chemical and thermal history on the other side, dynamical information is imprinted on the hierarchy of the Solar System. 

On a more local scale, the effects of turbulence are to impart random motions to particles that are then added to the motions that those particles experience owing to other effects. The details of the resulting motions are beyond the scope of this chapter (comprehensive discussion and equations are provided in Cuzzi Weidenschilling 2006 Ormel Cuzzi 2007). A characteristic velocity that describes these motions is the root-mean-square turbulent velocity, /acs which is the overturn velocity of the largest eddies, and an estimate of the maximum velocity two particles (both of St = 1) would develop with respect to one another (Cuzzi Hogan 2003). Under nominal conditions, such velocities can exceed 100 m s-1 (for a = 0.01). These velocities are important to consider when developing models for dust coagulation and planetesimal formation (see Section 3.4.1 and Chapters 7 and 10). 

Intrinsic turbulence in the disk poses severe problems for planetesimal formation by either pairwise sticking or gravitational instability. This has led to the development of a third class of model that embraces turbulence as a necessary ingredient for planetesimal formation. 

In this model, large gravitationally bound clumps form only occasionally, which would explain why planetesimal formation in the Asteroid Belt continued for several million years. This mechanism also reproduces the narrow size distribution of chondrules seen in chondritic meteorites (Cuzzi et al. 2001). While the mean chon-drule size differs from one meteorite type to another, the size distributions closely... 

Woolum D. S. and Cassen P. (1999) Astronomical constraints on nebula temperatures implications for planetesimal formation. Meteorit. Planet. Sci. 34, 897-907. 

Bridges F. G., Supulver K. D., Lin D. N. C., Knight R., and Zaffa M. (1996) Energy loss and sticking mechanisms in particle aggregation in planetesimal formation. Icarus 123, 422-435. 

Goodman J. and Pindor B. (1999) Secular instability and planetesimal formation in the dust layer. Icarus 148, 537-549. 

Ward W. R. (2001) On planetesimal formation the role of collective particle behaviour. Origin of the Earth and Moon (eds. R. M. Canup and K. Righter). University of Arizona Press, Tucson, pp. 75-84. 

Class III objects are weak line T Tauri stars (T Tauri stars in which the characteristic emission lines are only weakly observed in their optical spectra) and have little or no evidence of a disk. At this stage of solar nebula evolution, which may last between 3 and 30 Ma, the sun has formed, and the material of the nebula is being dissipated by solar winds in the inner part. In the outer part of the nebula material is dissipated by photo-evaporation caused by UV radiation from the solar wind. A positive pressure gradient near the inner edge of the nebula facilitates planetesimal formation. 

Markowski, A., Leya, 1., Quitte, G., Ammon, K., Halliday, A.N., and Wieler, R. (2006) Correlated heliiun-3 and tungsten isotopes in iron meteorites quantitative cosmogenic corrections and planetesimal formation times. Earth Planet. Sci. Lett., 250, 104-115. 

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