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Lecture L13 + L14 – ASTC25 Formation of planets

Lecture L13 + L14 – ASTC25 Formation of planets Top-down (GGP) hypotheis and its difficulties 2. Standard (core-accretion) scenario ( a ) From dust to planetesimals - two ways ( b ) From planetesimals to planetary cores:

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Lecture L13 + L14 – ASTC25 Formation of planets

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  1. Lecture L13 + L14 – ASTC25 • Formation of planets • Top-down (GGP) hypotheis and its difficulties • 2. Standard (core-accretion) scenario • ( a ) From dust to planetesimals - two ways • ( b ) From planetesimals to planetary cores: • how many planetesimals were there?( c ) Safronov number, runaway growth, • oligarchic growth of protoplanets

  2. Or some other way??

  3. Only theories involving disks make sense...

  4. Note: the standard scenario on the left also looks like the r.h.s. pictures…. Major difference: time of formation of giant protoplanets: 3-10 Myr (left panel) 0.1 Myr (right panel)

  5. There are two main possible modes of formation of giant gaseous planets and exoplanets: #bottom-up, or accumulation scenario for rocky cores (a.k.a. standard theory) predicts formation time ~(3-10) Myr (V.Safronov, G.Kuiper, A.Cameron) @top-down, by accretion disk breakup as a result of gravitational instability of the disk. A.k.a. GGP = Giant Gaseous Protoplanets formation time < 0.1 Myr (I.Kant, G.Kuiper, A.Cameron) To understand the perceived need for @, we have to consider disk evolution and observed time scales. To understand the physics of @, we need to study the stability of disks against self-gravity waves.

  6. Gravitational Instability and the Giant Gaseous Protoplanet hypothesis

  7. Fragmentation of disks

  8. Self-gravity as a destabilizing force for the epicyclic oscillations (radial excursions) of gas parcels on slightly elliptic orbits To study waves in disks, we substitute into the equations of hydrodynamics the wave in a WKBJ (WKB) approximation, also used in quantum mechanics: it assumes that waves are sinusoidal, tightly wrapped, or that kr >>1. All quantities describing the flow of gas in a disk, such as the density and velocity components, are Fourier-analyzed as , Some history WKB applied to Schrödinger equation (1925) Gregor Wentzel (1898-1978)German/American physicist Hendrick A. Kramers (1894–1952) Dutch physicist 1926 Léon N. Brillouin (1889-1969) French physicist Harold Jeffreys (1891 –1989) English mathematician, geophysicist, and astronomer, established a general method of approximation of ODEs in 1923

  9. Example of a crest of the spiral wave ~X1 exp[ … ] for k=const >0 , m=2, (omega)=const. This spiral pattern has constant shape and rotates with an angular speed equal to (omega)/m The argument of the exponential function is constant on the spiral wavecrest

  10. Dispersion Relation for non-axisymmetric waves in disks tight-winding (WKB) local approximation Doppler-shifted frequency epicyclic frequency self- gas gravity pressure In Keplerian disks, i.e.disks around point-mass objects

  11. The dispersion relation is, as in all the physics, a relation between the time and spatial frequencies, Though it looks more frightening than the one describing the simple harmonic (sinusoidal) sound wavein air: you can easily convince yourself that in the limit of vanishing surface density in the disk (no self-gravity!) and vanishing epicyclic frequency (no rotation!), the full dispersion relation assumes the above form. So the waves in a non-rotating medium w/o gravity are simply pure pressure (sound) waves. The complications due to rotation lead to a spiral shape of a sound wave, or full self-gravitating pressure wave. Don’t use THIS formula in disks!

  12. Dispersion Relation in disks with axisymmetric (m=0) waves (1960,1964)

  13. Gravitational stability requirements Local stability of disk, spiral waves may grow Local linear instability of waves, clumps form, but their further evolution depends on equation of state of the gas. It turns out that even at Q~1.5 there are unstable global modes.

  14. Question: Do we have to worry about self-gravity and instability? Ans: No z/r~0.1 qd>0.1 Ans: Yes

  15. From: Laughlin & Bodenheimer (2001) Disk in this SPH simulation initially had Q~1.8. The m-armed global spiral modes of the form grow and compete with each other. But the waves in a stable Q~2 disk stop growing and do not form small objects (GGPs).

  16. Recently, Alan Boss revived the half-abandoned idea of disk fragmentation Clumps forming in a gravitationally unstable disk (Q < 1) GGPs?

  17. Two examples of formally unstable disks not willing to form objects immediately Durisen et al. (2003) Break-up of the disk depends on the equation of state of the gas, and the treatment of boundary conditions.

  18. Simulations of self-gravitating objects forming in the disk (with grid-based hydrodynamics) shows that rapid thermal cooling is crucial Armitage and Rice (2003) Disk not allowed to cool rapidly (cooling timescale > 1 P) Disk allowed to cool rapidly (on dynamical timescale, <0.5 P)

  19. Mayer, Quinn, Wadsley, Stadel (2003) SPH = Smoothed Particle Hydrodynamics with 1 million particles Isothermal (infinitely rapid cooling)

  20. GGP (Giant Gaseous Protoplanet) hypothesis = disk fragmentation scenario (A. Cameron in the 1970s) Main Advantages: forms giant planets quickly, avoids possible timescale paradox; planets tend to form at large distances amenable to imaging. MAIN DIFFICULTIES: 1. Non-axisymmetric and/or non-local spiral modes start developing not only at Q<1 but already when Q decreases to Q~1.5…2 They redistribute mass and heat the disk => increase Q (stabilize disk). 2. Empirically, this self-regulation of the effects of gravity on disk is seen in disk galaxies, all of which have Q~2 and yet don’t split into many baby gallaxies. 3. The only way to force the disk fragmentation is to lower Q~c/Sigma by a factor of 2 in just one orbital period. This seems impossible. 4. Any clumps in disk (a la Boss’ clumps) may in fact shear and disappear rather than form bound objects. Durisen et al. Have found that the equation of state and the correct treatment of boundary conditions are crucial, but could not confirm the fragmentation except in the isothermal E.O.S. case. 5. GGP is difficult to apply to Uranus and Neptune; final masses: Brown Dwarfs not GGPs 6. Does not easily explain core masses of planets and exoplanets, nor the chemical correlations (to be discussed in lecture L23)

  21. Standard Accumulation Scenario

  22. Two-stage accumulation of planets in disks

  23. Planetesimal = solid body >1 km Mcore=10 ME(?) => contraction of the atmosphere and inflow of gas from the disk (issues not addressed in the standard theory so far)

  24. Two scenarios proposed for planetesimal formation Particles settle in a very thin sub-disk, in which Q<1, then gravitational instability in dust layer forms planetesimals Particles in a turbulent gas not able to achieve Q<1, stick together via non-gravity forces. gas gas Q<1 Q>1 Solid particles (dust, meteoroids)

  25. How many planetesimals formed in the solarnebula?

  26. How many planetesimals formed in the solarnebula? Gas mass ~ 0.02-0.1 Msun > 2e33g*0.02=4e31 g ~ 1e32 g Dust mass ~ 0.5% of that ~ 1e30 g ~ 100 Earth masses (Z=0.02 but some heavy elements don’t condense) Planetesimal mass, s =1 km = 1e5 cm, ~(4*pi/3) s^3*(1 g/cm^3) => 1e16 g Assuming 100% efficiency of planet formation N = 1e30g /1e16g = 1e14, s=1 km N = 1e30g /1e19g = 1e12, s=10 km (at least that many initially) N = 1e30g /1e22g = 1e8, s=100 km N = 1e30g /1e25g = 1e5, s=1000 km N = 1e30g /1e28g = 100, s=10000 km (rock/ice cores & planets) N = 1e30g /1e29g = 10, (~10 ME rock/ice cores of giant planets)

  27. Gravitational focusing factor

  28. Stopping the runaway growth of planetary cores Roche lobe radius grows non-linearly with the mass of the planet, slowing down the growth as the mass (ratio) increases. The Roche lobe radius rL is connected with the size of the Hill-stable disk region via a factor 2*sqrt(3), which, like the size of rL, we derived earlier in lecture L12. This will allow us to perform a thought experiment and compute the maximum mass to which a planet grows spontaneously by destabilizing further and further regions.

  29. Isolation mass in different parts of the Minimum Solar Nebula *Based on Minimum Solar Nebula (Hayashi nebula) = a disk of just enough gas to contain the same amount of condensable dust as the current planets; total mass ~ 0.02 Msun, mass within 5AU ~0.002 Msun Conclusions: (1) the inner & outer Solar System are different: critical core=10ME could only be achieved in the outer sol.sys. (2) there was an epoch of giant impacts onto protoplanets when all those semi-isolated ‘oligarchs’ where colliding

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