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From Pebbles to Planetesimals and Beyond Overview of Topics

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From Pebbles to Planetesimals and Beyond Overview of Topics From pebbles to planetesimals ... and beyond Anders Johansen (Lund University) \Origins of stars and their planetary systems" Hamilton, June 2012 1 / 16 Anders Johansen From pebbles to planetesimals and beyond Overview of topics Size and time Dust Pebbles Planetesimals Planets µ m cm 10−1,000 km 10,000 km 1 Forming large planetesimals by streaming instabilities Johansen, Youdin, & Lithwick (2012) 2 Rapid growth of gas-giant cores by pebble accretion Lambrechts & Johansen (A&A, in press) 3 Using Kepler data to learn about planet formation Johansen, Davies, Church, & Holmelin (ApJ, submitted) Buchhave, Latham, Johansen, Bizzarro, Torres, Rowe, et al. (Nature, in press) 2 / 16 Anders Johansen From pebbles to planetesimals and beyond Early stages of planet formation Dust and ice particles in the protoplanetary disc coagulate to cm-sized pebbles and rocks Pebbles and rocks sediment to the mid-plane of the disc Further growth frustrated by high-speed collisions (>1{10 m/s) which lead to erosion and bouncing ) Need some way for particle layer to get dense enough to initiate gravitational collapse 3 / 16 Anders Johansen From pebbles to planetesimals and beyond How turbulence aids planet formation 1 Passive concentration as particles pile up in long-lived high-pressure regions excited in the turbulent gas flow (e.g. Whipple 1972; Klahr & Bodenheimer 2003; Johansen et al. 2009; Lyra et al. 2009; Simon et al. 2012) 2 Active concentration as particles make dense filaments and clumps to protect themselves from gas friction (e.g. Youdin & Goodman 2005; Johansen & Youdin 2007; Johansen et al. 2009; Bai & Stone 2010) 4 / 16 Anders Johansen From pebbles to planetesimals and beyond Streaming instability 1 Gas orbits slightly slower than Keplerian 2 Particles lose angular momentum due to headwind and spiral in towards the star 3 Particle clumps locally reduce headwind and are fed by isolated particles ) Run-away pile-up of particles v Kep (1− η ) F F GP Youdin & Goodman (2005): \Streaming instability" (confirmed numerically by Youdin & Johansen 2007; Johansen & Youdin 2007; Bai & Stone 2010) 5 / 16 Anders Johansen From pebbles to planetesimals and beyond Streaming instability 1 Gas orbits slightly slower than Keplerian 2 Particles lose angular momentum due to headwind and spiral in towards the star 3 Particle clumps locally reduce headwind and are fed by isolated particles ) Run-away pile-up of particles v Kep (1− η ) F F GP Youdin & Goodman (2005): \Streaming instability" (confirmed numerically by Youdin & Johansen 2007; Johansen & Youdin 2007; Bai & Stone 2010) 5 / 16 Anders Johansen From pebbles to planetesimals and beyond Planetesimal formation t=120.4 t=130.4 t=140.4 0.10 Clumping and 1283 0.05 No collisions planetesimal formation 20.0 H N=13 N=7 / Σ Σ 0.00 p/< p> triggered at metallicity y Mbound= 1.34 Mbound= 1.35 M M 0.0 1= 0.44 1= 0.95 −0.05 M = 0.19 M = 0.16 slightly above solar 2 2 M13= 0.0061 M7= 0.029 −0.10 0.10 Largest planetesimals are Collisions with shear 0.05 ε typically 500 km in radius =0.3 H N=7 N=3 / 0.00 y Mbound= 1.35 Mbound= 1.47 M1= 0.59 M1= 0.75 −0.05 M2= 0.49 M2= 0.70 Comparable to dwarf M = 0.020 M = 0.014 7 3 −0.10 0.10 planet Ceres Collisions without shear 0.05 ε =0.3 H N=9 N=6 / 0.00 ) Asteroids born big? y M = 1.42 M = 1.46 (Morbidelli et al. 2009) bound bound M1= 0.45 M1= 0.63 −0.05 M2= 0.33 M2= 0.59 M9= 0.012 M6= 0.021 −0.10 Scaling to Kuiper belt 0.10 Collisions with shear gives twice as large 0.05 ε =1.0 H N=7 N=3 / 0.00 planetesimals y Mbound= 1.37 Mbound= 1.44 M1= 0.54 M1= 0.63 −0.05 M2= 0.31 M2= 0.59 M7= 0.0052 M3= 0.22 ) May explain why Kuiper −0.10 −0.10 −0.05 0.00 0.05 −0.10 −0.05 0.00 0.05 −0.10 −0.05 0.00 0.05 0.10 belt objects are larger x/H x/H x/H than asteroids (Johansen, Youdin, & Lithwick 2012) 6 / 16 Anders Johansen From pebbles to planetesimals and beyond Classical core accretion scenario 1 Dust grains and ice particles collide to form km-scale planetesimals 2 Large protoplanet grows by run-away accretion of planetesimals 3 Protoplanet attracts hydrostatic gas envelope 4 Run-away gas accretion as Menv ≈ Mcore 5 Form gas giant with Mcore ≈ 10M⊕ and Matm ∼ MJup (e.g. Safronov 1969, Mizuno 1980, Pollack et al. 1996, Rafikov 2004) 7 / 16 Anders Johansen From pebbles to planetesimals and beyond Core formation time-scales The size of the protoplanet relative to the Hill sphere: R r −1 p ≡ α ≈ 0:001 RH 5 AU Maximal growth rate _ 2 M = αRHFH ) Only 0.1% (0.01%) of planet- esimals entering the Hill sphere are accreted at 5 AU (50 AU) Planet Gravitational cross section ) Time to grow to 10 M⊕ is ∼10 Myr at 5 AU Hill sphere ∼50 Myr at 10 AU ∼5,000 Myr at 50 AU 8 / 16 Anders Johansen From pebbles to planetesimalsx=0 and beyond x Directly imaged exoplanets (Marois et al. 2008; 2010) (Kalas et al. 2008) HR 8799 (4 planets at 14.5, 24, 38, 68 AU) Fomalhaut (1 controversial planet at 113 AU) ) No way to form the cores of these planets within the life-time of the protoplanetary gas disc by standard core accretion 9 / 16 Anders Johansen From pebbles to planetesimals and beyond Pebble accretion Most planetesimals are simply scattered by the protoplanet Planetesimal is scattered by protoplanet Pebbles spiral in towards the protoplanet due to gas friction ) Pebbles are accreted from the entire Hill sphere Pebble spirals towards protoplanet due to gas friction Growth rate by planetesimal accretion is _ 2 M = αRHFH Growth rate by pebble accretion is _ 2 M = RHFH 10 / 16 Anders Johansen From pebbles to planetesimals and beyond Time-scale of pebble accretion 0.10 0.10 Ω−1 Ω−1 5.0 0.05 t=0.0 t=120 > p Σ 0.00 0.05 z/H /< p −0.05 Σ −0.10 0.0 −0.10 −0.05 0 0.05 0.10 0.00 x/H Core growth to 10 M⊕ y/H 108 Planetesimals −0.05 107 −0.10 0.10 6 t=131 Ω−1 t=134 Ω−1 10 /yr t 0.05 5 ∆ 10 Fragments Pebbles 0.00 4 y/H 10 3 −0.05 10 −1 0 1 2 −0.10 10 10 10 10 −0.10 −0.05 0.00 0.05 0.10 −0.05 0.00 0.05 0.10 x/H x/H r/AU ) Pebble accretion speeds up core formation by a factor 1,000 at 5 AU and a factor 10,000 at 50 AU (Lambrechts & Johansen, in press; also Ormel & Klahr 2010, Murray-Clay, Kratter, & Youdin, in prep) ) Cores form well within the life-time of the protoplanetary gas disc, even at large orbital distances Requires large planetesimal seeds, consistent with turbulence-aided planetesimal formation 11 / 16 Anders Johansen From pebbles to planetesimals and beyond Kepler planets J 10 N E R / R 1p 2p 3p E 1 1 10 100 P/days In the first 16 months the Kepler satellite detected 2321 planet candidates 253 Earth-sized (R≤1:25R⊕) 245 double systems 712 Super-Earth-sized (1:25R⊕<R ≤ 2R⊕) 84 triple systems 1078 Neptune-sized (2R⊕<R≤6R⊕) 27 quadruple systems 207 Jupiter-sized (6R⊕<R≤15R⊕) 8 quintuple systems 71 Super-Jupiter-sized (15R⊕<R) 1 sextuple system 12 / 16 Anders Johansen From pebbles to planetesimals and beyond Mutual inclination between planetary orbits β = 0◦ β = 5◦ Ratio of single and double transits to triple transits increases with increasing mutual inclination parameter β ) Relative numbers of multiple transit systems constrain underlying planetary system architectures Johansen, Davies, Church, & Holmelin (ApJ, submitted): Use Kepler triple systems as templates for physical planet systems and perform synthetic transit observations to match Kepler data 13 / 16 Anders Johansen From pebbles to planetesimals and beyond Planetary archictectures J 10 N E R / R 1p 2p 3p E 1 1 10 100 P/days Number of triple-planet systems increases with β Number of double-planet systems decreases with β For β > 5◦ triple-planet systems make too many double transits ) planetary systems are flat (Lissauer et al. 2011, but see also Tremaine & Dong 2012) Around 20% of solar type stars have planets of Earth size or larger within 0.5 AU Gas-giant planets appear almost exclusively in single-planet systems (The Kepler Dichotomy) 14 / 16 Anders Johansen From pebbles to planetesimals and beyond The Kepler dichotomy and metallicity Instability hypothesis First determination of metallicities of Kepler planet hosts: Suppressed formation hypothesis (Buchhave, Latham, Johansen, et al., Nature, in press) (Johansen et al., ApJ, submitted) ) Terrestrial planets form around stars with a wide range of metallicities ) Metallicity determines whether planetary systems have gas giants within 0.5 AU 15 / 16 Anders Johansen From pebbles to planetesimals and beyond The Kepler dichotomy and metallicity Instability hypothesis First determination of metallicities of Kepler planet hosts: 100 10 p N Suppressed formation hypothesis 1 Large (R>4R ) R RE Small ( <4 E) −0.6 −0.4 −0.2 −0.0 0.2 0.4 [m/H] (Johansen et al., ApJ, submitted) (Buchhave, Latham, Johansen, et al., Nature, in press) ) Terrestrial planets form around stars with a wide range of metallicities ) Metallicity determines whether planetary systems have gas giants within 0.5 AU 15 / 16 Anders Johansen From pebbles to planetesimals and beyond Take home messages Pebbles undergo spontaneous clumping by streaming instabilities and form large planetesimals Cores of gas giants and ice giants grow rapidly by pebble accretion, even in 50{100 AU orbits Planetary systems come in two flavours within 0.5 AU, dominated either by one or more small planets or by a single gas-giant planet This dichotomy is likely triggered by metallicity as gas-giant planets disrupt terrestrial planet formation Metallicity is less important for terrestrial planet formation than previously thought 16 / 16 Anders Johansen From pebbles to planetesimals and beyond.
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