Origins of Gas Giant Planets - PDF Free Download

Origins of Gas Giant Planets Ruth Murray-Clay Harvard-Smithsonian Center for Astrophysics Image Credit: NASA

Graduate Students Piso Tripathi Dawson Undergraduates Wolff Lau Alpert Mukherjee Wolansky Jackson

HR 8799: A testbed for planet formation theories Neptune s orbital distance Marois et al. 2010

NASA, ESA, M. Robberto, HST Orion Treasury Project, L. Ricci

Alves, Lada, & Lada 2001

gas + dust + ice gas + dust

Test case HR 8799: Brown Dwarfs or Planets? M p /M * 10 0 10 1 10 2 ~stars ~brown dwarfs 10 3 Jupiter ~planets 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu Saturn 10 0 10 2 10 4 r p (AU)

Test case HR 8799: Brown Dwarfs or Planets? M p /M * 10 0 10 1 10 2 ~stars ~brown dwarfs Planetary companions 10 3 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu Jupiter Saturn 10 0 10 2 10 4 r p (AU) ~planets

Test case HR 8799: Brown Dwarfs or Planets? Brown dwarf companions M p /M * 10 0 10 1 10 2 ~stars ~brown dwarfs Planetary companions 10 3 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu Jupiter Saturn 10 0 10 2 10 4 r p (AU) ~planets

Test case HR 8799: Brown Dwarfs or Planets? M p /M * Planetary companions 10 0 10 1 10 2 10 3 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu 10 0 10 2 10 4 r p (AU) Brown dwarf companions HR 8799 Jupiter Saturn ~stars ~brown dwarfs ~planets Larger than most protoplanetary disks

Three ways to form companions turbulent fragmentation gravitational instability core accretion

Three ways to form companions HR 8799 turbulent fragmentation gravitational instability core accretion

Today s Goal: A framework for identifying giant planet formation processes. using HR 8799 as a test case Marois et al. 2010

Outline: ~70 AU Origin Stories ~ in situ formation by gravitational instability ~ in situ formation by core accretion Formation closer to star + scattering outward More complicated hybrid scenarios e.g., form a core close to the star, scatter outward, then accrete gas

Gravitational instability L self-gravity overcomes pressure and Keplerian shear most unstable scale: L H 3 M M frag (2 H) 2 4 H a

Gravitational Instability? Planets cannot grow after fragmentation and they must migrate in 100 minimum fragment distance Mass (MJup) 10 e minimum fragment mass 1 20 60 100 140 Radius (AU) Kratter, Murray-Clay, & Youdin, ApJ 2010

Gravitational Instability? Planets cannot grow after fragmentation and they must migrate in 100 minimum fragment distance Mass (MJup) 10 e minimum fragment mass Stability requires resonance, suggesting migration Fabrycky & Murray-Clay 2010 1 20 60 100 140 Radius (AU) Kratter, Murray-Clay, & Youdin, ApJ 2010

Collapse must occur at the end of infall or the fragment will grow into a binary star Ṁ Ṁ disk

Collapse must occur at the end of infall or the fragment will grow into a binary star Ṁ Ṁ disk isolation mass is stellar!

Gravitational instability planets can only be failed binary stars 10 0 ~stars M p /M * 10 1 10 2 10 3 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu HR 8799 expect higher mass companions from disk collapse 10 0 10 2 10 4 r p (AU) ~brown dwarfs ~planets Larger than most protoplanetary disks

Core accretion: Need to grow a solid core through collisions

Growth timescale is set by: A: cross-section for collisions v: relative velocities of colliding bodies n: density of bodies available to accrete v t Volume = Avt A number density = n v H ~ v torb Sun # collisions in time t ~ n Avt ~ (nh) A (t/torb)

mass of object trying to grow orbital time: longer farther from the star t grow = Ṁ M M A t orb disk surface density in solids cross-section for collisions

Cross-section regimes: physical cross-section gravitational focusing capture by gas drag into planetary atmosphere?

Let s make gas useful larger planetesimal or core Sun planetesimal

Gas Alters the Orbits of Single Planetesimals planetesimal wants to orbit The resulting F D m star at v Kep drag acceleration is small large radial pressure gradient but gas orbits more slowly v orb <v Kep v 2 orb r = GM r 2 + 1 dp dr

In the absence of gas, satellites can orbit within the Hill radius RHill Sun

In the absence of gas, satellites can orbit within the Hill radius RHill Sun planetary gravity

In the absence of gas, satellites can orbit within the Hill radius RHill tidal acceleration Sun planetary gravity

Wind Shearing (WISH) No gas: Gas: planetesimal RHill RHill core RWISH Perets & Murray-Clay (2011)

A core + planetesimal in gas can: shear apart orbit around star mutual orbit or inspiral and merge

A core + planetesimal in gas can: shear apart orbit around star mutual orbit or inspiral and merge F D m F D m

Wind-Shearing Limits Orbital Stability Hill radius WISH radius large body radius core radius R stab /r b 10 5 10 4 10 3 10 2 10 1 10 0 10 1 1 AU r b = 10km R H /r b Epstein Ram Pressure Drag 1 < Re < 800 Stokes R WS /r b r b /r b 10km r s r b = core radius Ram binary stability 10 4 10 2 10 0 10 2 10 4 10 6 10 8 r s (cm) small body radius (cm) Perets & Murray-Clay (2011)

A core + planetesimal in gas can: shear apart orbit around star mutual orbit or inspiral and merge F D m F D m

τmerge (yr) 10 merge (yr) time toτ merge (yr) 2 100 Epstein 1010 ~Pluto radius 40 AU 1 AU 8 typical disk lifetime Stokes: 2 τmerge rs 104 10 108 106 ~km 8 Stokes: 2 τmerge rs 104 Porb 102 102 10 40 AU 106 Ram: τmerge rs Porb 0 Epstein 10 10 106 Stokes Porb Porb Porb Epstein: τmerge rs 100 Stokes 10 (yr) 104 Stokes 100 1010 106 For Small Planetesimals, Merger Times <10Disk Lifetimes 4 τmerge (yr) τmerge (yr) 10 102 Ram Ram 6 10 5 4 AU 10 2 100 102 rs (cm) 104 106 108 Epstein: τmerge rs 10 4 10 2 100 102 rs (cm) 104 106 108 small body radius (cm) Ram Perets & Murray-Clay (2011)

Binary Capture RHill dissipation due to interaction with gas RWISH Murray-Clay & Perets (in prep) (see also Ormel & Klahr 2010)

Integration of forces confirms that capture can happen with a cross-section as large as the Hill radius 0.02 y (AU) 0.01 0.00 0.01 RHill Ratm 0.02 0.02 0.01 0.00 0.01 0.02 x (AU)

RHill Ratm Capture by the atmosphere is only possible if the small particle RWISH can decouple from the exterior gas.

Integrations confirm that well-entrained particles can be swept around the core, preventing accretion 0.02 y (AU) 0.01 0.00 0.01 RHill Ratm RWISH 0.02 0.02 0.01 0.00 0.01 0.02 x (AU)

Growth times at 70 AU can be short enough to nucleate an atmosphere Mcore (MEarth) 10 3 10 2 10 1 10 0 10 1 10 2 10 3 10 4 10 4 10 5 10 6 time (yr) M crit (Rafikov 2006,2010) 50% of MMSN solids in equal mass per log bin from mm to 10cm Murray-Clay et al. (in prep.)

Turbulence doesn t prevent the final stage of core growth by capture of planetesimals 0.02 y (AU) 0.01 0.00 0.01 RHill Ratm 0.02 0.02 0.01 0.00 0.01 0.02 x (AU) =0.01

HR 8799: Brown Dwarfs or Planets? 10 0 ~stars M p /M * 10 1 10 2 10 3 HR 8799 ~brown dwarfs ~planets 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu 10 0 10 2 10 4 r p (AU) Larger than most protoplanetary disks

10 0 Gemini Planet Imager (PI: Bruce Macintosh) ~stars M p /M * 10 1 10 2 10 3 HR 8799 ~brown dwarfs ~planets 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu 10 0 10 2 10 4 r p (AU) Larger than most protoplanetary disks

Overlapping populations will have different mass distributions 10 0 10 1 M p /M * 10 2 10 3 10 4 10 0 10 2 10 4 r p (AU)

Overlapping populations will have different mass distributions Excluding radial orbits: 10 0 10 1 Number M p /M * 10 2 Mass 10 3 10 4 10 0 10 2 10 4 r p (AU) Number Mass Stellar metallicity and planetary atmospheric composition provide additional discriminants

CO snowline ~40 AU C and O frozen out together H2O snowline ~2 AU C rich gas O rich solids C Oberg, Murray-Clay, & Bergin, ApJL (2011) O

Metal-rich stars host more moved planets: A signature of planet-planet interactions planets{ moved hot Jupiters highly eccentric planets ~ in situ formation? Dawson & Murray-Clay 2013

Test case HR 8799: Brown Dwarfs or Planets? 10 0 GPI ~stars M p /M * 10 1 10 2 HR 8799 ~brown dwarfs 10 3 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu Jupiter Saturn 10 0 10 2 10 4 r p (AU) ~planets Larger than most protoplanetary disks

How do giant planets and brown dwarfs form? M p /M * 10 1 10 2 10 3 core accretion disk gravitational instability? 10 0 10 4 Kratter, Murray-Clay, & Youdin, ApJ (2010) Data: Zuckerman & Song 2009; exoplanet.eu HR 8799: signpost Jupiter Saturn 10 0 10 2 10 4 r p (AU) turbulent fragmentation ~stars ~brown dwarfs ~planets Larger than most protoplanetary disks

Formation and evolution of planetary systems around stars with different masses, at different stages of evolution, and in different environments. Infall and disk accretion sets the initial conditions for growth of planetesimals and planets Piso, Youdin, & Murray-Clay (final stages) Wolansky (senior thesis) Mukherjee Graduate Students Undergraduates which then migrate as they interact with other planets, residual gas and planetesimals Wolff, Dawson, & Murray-Clay (2012) Dawson & Murray-Clay (2012) Dawson, Murray-Clay, & Fabrycky (2011) Lau & Murray-Clay (final stages) Jackson (junior project) and then evolve dynamically over long timescales all the while developing atmospheres that depend on this history. Tripathi Dawson, Murray-Clay, & Johnson (submitted) Dawson et al. (2012) Dawson & Murray-Clay (2013) Alpert