Giant Planet Accretion and Migration : Surviving the Type I Regime Edward Thommes Norm Murray CITA, University of Toronto Edward Thommes Norm Murray CITA, University of Toronto The Western Workshop, UWO, May 19, 2006 JPL

Gas giant formation: The core accretion model  Gas disk lifetime sets upper limit on gas giant formation: ~1-10 Myrs from observations (e.g. Haisch, Lada & Lada 2001)  The core accretion model (Mizuno 1980, Pollack et al 1996): 1.Solid core grows, ~10 M Earth 2.Core accretes massive gas envelope, 100+ M Earth  Observational support for core accretion:  planet-metallicity correlation (Gonzalez 1997, Fischer & Valenti 2003)  HD planet (Saturn mass, ~70 M Earth core; Sato et al. 2005, Charbonneau et al 2006)  Gas disk lifetime sets upper limit on gas giant formation: ~1-10 Myrs from observations (e.g. Haisch, Lada & Lada 2001)  The core accretion model (Mizuno 1980, Pollack et al 1996): 1.Solid core grows, ~10 M Earth 2.Core accretes massive gas envelope, 100+ M Earth  Observational support for core accretion:  planet-metallicity correlation (Gonzalez 1997, Fischer & Valenti 2003)  HD planet (Saturn mass, ~70 M Earth core; Sato et al. 2005, Charbonneau et al 2006) Marcy et al 2005

Planet-disk interaction  Presence of substantial gas disk means planet-disk interactions important!  Bodies in gas disk launch density waves  repulsive torque between body and inner, outer disk  Jupiter-mass planets open a gap, locked into viscous evolution of disk: “Type II” inward migration  Smaller bodies: no gap, outer torques stronger: “Type I” inward migration  Presence of substantial gas disk means planet-disk interactions important!  Bodies in gas disk launch density waves  repulsive torque between body and inner, outer disk  Jupiter-mass planets open a gap, locked into viscous evolution of disk: “Type II” inward migration  Smaller bodies: no gap, outer torques stronger: “Type I” inward migration Density of planet  disk torque Ward 1997 Geoff Bryden

Migration and accretion rates

Comparing the timescales  Scary result! Thus people tend to ignore/greatly reduce Type I (e.g. Thommes, Duncan & Levison 2003, Ida & Lin 2004, Alibert et al. 2005) But is there a way to make the worst-case scenario work...?  Scary result! Thus people tend to ignore/greatly reduce Type I (e.g. Thommes, Duncan & Levison 2003, Ida & Lin 2004, Alibert et al. 2005) But is there a way to make the worst-case scenario work...?

Accretion, no migration Thommes & Murray 2006

Accretion + Migration Thommes & Murray 2006

A viscously evolving disk t=0 t=1 Myr t=10 Myrs

Accretion + Migration in a viscously evolving gas disk Thommes & Murray 2006  =10 -2 M disk M 100 AU M 30 AU Disk gas mass

Winners and losers  Inner region: growth too fast, cores lost onto star  Outer region: growth too slow relative to disk lifetime  In between: An annulus where the growth rate turns out just right  Inner region: growth too fast, cores lost onto star  Outer region: growth too slow relative to disk lifetime  In between: An annulus where the growth rate turns out just right Thommes & Murray 2006

 Method:  Vary disk mass, metallicity,   For each set (M D, ,[Fe/H]), compute largest protoplanet mass when 1 M Jup of gas left inside 100 AU  Results  Disks with higher M D, [Fe/H] do better  There is always an “optimal” , ~ ; consistent with fits to T Tauri disks (Hartmann et al 1998)  Method:  Vary disk mass, metallicity,   For each set (M D, ,[Fe/H]), compute largest protoplanet mass when 1 M Jup of gas left inside 100 AU  Results  Disks with higher M D, [Fe/H] do better  There is always an “optimal” , ~ ; consistent with fits to T Tauri disks (Hartmann et al 1998) Thommes & Murray 2006 Disk properties and core formation

Summary  In the worst-case scenario of unmitigated Type I migration:  protoplanets in a young, massive gas disk fall onto central star long before they can reach gas giant core size (~10 M Earth )... ...but as the gas disk dissipates, a window may open for cores to form and survive  endgame: gas envelope accretion plays large role in cleaning up rest of disk (cf. Lecar & Sasselov 2003)  Predictions  Favourable disk properties: high M(0), high [Fe/H], and  ~  no giant planets (i.e. for ALMA, no gaps) in very young, massive disks  In the worst-case scenario of unmitigated Type I migration:  protoplanets in a young, massive gas disk fall onto central star long before they can reach gas giant core size (~10 M Earth )... ...but as the gas disk dissipates, a window may open for cores to form and survive  endgame: gas envelope accretion plays large role in cleaning up rest of disk (cf. Lecar & Sasselov 2003)  Predictions  Favourable disk properties: high M(0), high [Fe/H], and  ~  no giant planets (i.e. for ALMA, no gaps) in very young, massive disks

“Dead zones” in disks  Magnetorotational instability (MRI) (Balbus & Hawley 1991) leading candidate for disk viscosity  MRI requires ionized disk, to couple it to magnetic field  cosmic rays, stellar X-rays (near star)  When the full vertical column not ionized, dead zone forms (Gammie 1996, Matsumura & Pudritz 2003)  Magnetorotational instability (MRI) (Balbus & Hawley 1991) leading candidate for disk viscosity  MRI requires ionized disk, to couple it to magnetic field  cosmic rays, stellar X-rays (near star)  When the full vertical column not ionized, dead zone forms (Gammie 1996, Matsumura & Pudritz 2003) Gammie 1996

Disk evolution with a dead zone  Dead zone: lower viscosity  slower accretion  pile-up of gas  Steep jumps in surface density can result  How does this affect migration...?  Dead zone: lower viscosity  slower accretion  pile-up of gas  Steep jumps in surface density can result  How does this affect migration...? ?

Disk torques at a surface density jump  Type I migration:  inner <  outer,  gas  Introducing jump in  gas can reverse the torque imbalance  outer edge of a dead zone can completely stop Type I migration!  Type I migration:  inner <  outer,  gas  Introducing jump in  gas can reverse the torque imbalance  outer edge of a dead zone can completely stop Type I migration! Matsumura, Thommes & Pudritz, in prep.

Thommes

A “hybrid” code: N-body+gas disk  The N-body part: SyMBA (Duncan, Levison & Lee 1998)  uses Wisdom-Holman (1991) symplectic method  fast for near-Keplerian systems  bounded energy error  resolves close encounters  The disk-evolution part: 1-D (azimuthally, vertically averaged) Keplerian disk, Σ evolves according to  The N-body part: SyMBA (Duncan, Levison & Lee 1998)  uses Wisdom-Holman (1991) symplectic method  fast for near-Keplerian systems  bounded energy error  resolves close encounters  The disk-evolution part: 1-D (azimuthally, vertically averaged) Keplerian disk, Σ evolves according to (Goldreich & Tremaine 1980, Ward 1997)  -∫(dT/dr)dr applied to planet ...Fast! Can simulate 10 7 yrs in ~2 days -

Thommes 2005

Resonant exoplanets Marcy et al. 2005

The “standard model” of core accretion  Pollack et al (1996): 3 stages: 1.solid core accretion 2.slow gas accretion until M gas ~ M core 3.runaway gas accretion  Corrections to the standard model:  Stage 1 simplified, actually takes longer (e.g. Thommes et al. 2003)  Stage 2 HAS to be a lot shorter (can be done by lowering envelope opacity)  Pollack et al (1996): 3 stages: 1.solid core accretion 2.slow gas accretion until M gas ~ M core 3.runaway gas accretion  Corrections to the standard model:  Stage 1 simplified, actually takes longer (e.g. Thommes et al. 2003)  Stage 2 HAS to be a lot shorter (can be done by lowering envelope opacity) Pollack et al. 1996

Outline  Background  giant planet formation by core accretion  migration by planet-disk interaction  The timescale problem  Calculations of concurrent core accretion and migration in an evolving disk  A way around the timescale problem  Disk properties and the prospects for planet formation  Summary  Background  giant planet formation by core accretion  migration by planet-disk interaction  The timescale problem  Calculations of concurrent core accretion and migration in an evolving disk  A way around the timescale problem  Disk properties and the prospects for planet formation  Summary