Planet Formation Phil Armitage
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Planet Formation Phil Armitage Colorado è Stony Brook / Simons CCA overview • stone age planet formation – formation of planetary systems from a smooth radial distribution of km-scale planetesimals embedded in gas, with no migration • planetesimal • planetesimals formation or pebbles? • migration • gas accretion How do we form observed systems? Is there a timing problem? classical planet formation Adopt effective initial conditions: radially smooth distribution of small (100m – km) planetesimals embedded in gas, which acts to damp {e,i} but does not cause migration • growth phases: runaway, oligarchic final assembly / giant impacts • if fast enough final outcome is ~determined by stability M 1/3 r = p a Hill 3M ✓ ⇤ ◆ 1/4 tinstability = f(∆a/Mp ) • collisions make a small-N system more stable classical planet formation t / Myr 100 Ratio of collisions to scatterings is 10 f(vK / vesc), favors collisions for MSun at a ~ AU and less 1 “in situ” formation e 1 AU 2 AU 3 AU • disk with a few ME / AU at 1 AU will work • requirement of stability can hide a multitude of sins • many important but lesser constraints (small mass of Mars…) classical planet formation quasi-static envelope extends to rHill, rBondi core grows from accretion of planetesimals cooling is limited by grain opacity in radiative zone in the envelope “success” requires reaching Menv ~ Mcore ~ 5-20 ME within gas disk lifetime possible at few AU if the gravitational focusing factor 2 [1 + (vesc / σ) ] is large (small planetesimals) and κ low what’s wrong with this story? • major physical problem in omission of migration torque for Jupiter’s core • BUT internally consistent model for Solar System terrestrial planets and ~1 ME planets at ~1 AU around Solar-type stars… final assembly >> gas disk lifetime • DOES NOT scale to ~1 ME planets in habitable zone of low mass stars, or to 5-10 ME super-Earths • Solar System occupies an unusual corner of parameter space? getting to planetesimals • experimental / modeling evidence that we can reach mm-cm scales in inner disk • water ice physics promotes higher vf and possibly larger sizes if drift is ignored (Gundlach & Blum 2015; Wada+ 2009) Birnstiel+ 2010 Very high confidence that we attain conditions linearly unstable to streaming instability (Youdin & Goodman 2005) getting to planetesimals For gravitational collapse particle density needs to reach Roche density: M ⇢ ⇤ p ⇠ a3 Extremely strong, non- linear clustering Yang+ 2017 • best determinations require dust / gas ratios > 10-2 • sweet spot at dimensionless stopping time τ ~ 0.1 getting to planetesimals data: Zhaohuan Zhu, Jake Simon When collapse occurs, moderately high confidence: • top-heavy mass function dN 1.6 Mp− (Simon+ 16; Schafer+ 17) dMp / • universal: no measured dependence on the aerodynamic properties of the participating particles (Simon+ 17) • implies large planetesimals dominate (~102 km) how to attain conditions for planetesimal formation, pick your poison especially Z > 10-2 • radial drift leading to particle pile-up in inner disk (Youdin & Chiang 2004) • local concentration in pressure maxima (Pinilla+ 2012) • “unique” processes at ice lines (Stevenson & Lunine 1988; Ros & Johansen 2013) observational constraints on the size and radial distribution of solids in disks of different ages are key (e.g. Tazzari+ 2016) next steps… 2D slice of a 3 or 4D parameter space: key physics input for model building {τ,Z,η,α} p = -1.6 is not p = -2 in current simulations, but are there conditions where mass function is not top-heavy? role for other flavors of gas-dust instability (Squire & Hopkins 18) pebble accretion aerodynamically assisted accretion of radially drifting small solids can be efficient growth channel (Ormel & Klahr 2010; Lambrechts & Johansen 2012) Hill limited Hill radius r accretion Drift limited H accretion Keplerian shear Particle radial drift -1/2 Accretion radius (tB/tf) rB Bondi radius rB pebble accretion physically entirely independent of streaming instability / planetesimal formation n BUT for a disk with a power-law pressure profile: P r− “optimal” sized pebbles have, / 6 8 Mp h − ⌧ = n3 M r ✓ ⇤ ◆✓ ◆ …very roughly τ ~ 1 solids can do double duty in first forming planetesimals and then accreting rapidly onto proto-planets MOREOVER if planetesimals form with a top-heavy mass function gravitational focusing is reduced, relative efficiency of pebble vs planetesimal accretion is higher migration • Lindblad torque, whose asymmetry is a weak function of disk structure (Ward 1997) circulating librating streamlines streamlines • co-orbital torque, low low entropy density dependent on disk gradients in entropy, high high horseshoe region density entropy surface density, and disk r net torque φ diffusion (Paardekooper+ on planet 2011) • further modification due to thermal effects for luminous planets (Lega+ 14; Benitez-Llambay+ 15; Masset 17) migration • possible but not very plausible to construct disk models where Type I migration torques are generically zero Bitsch+ 2013 • more likely a discrete set of null points for migration, where giant planet cores would accumulate • where these are depends on disk structure, second Zhaohuan’s polite request that you measure it for us J envelope accretion quasi-static envelope extends to rHill, rBondi Simulations: disk gas can circulate into and out of Hill sphere (D’Angelo & Bodenheimer 13; Ormel+ 15; Fung+ Lambrechts 15; Lambrechts & Lega 17) & Lega 2017 Key physics relevant to the distinction between mini-Neptunes, Neptunes, and gas giants forming Kepler planetary systems growth timescales for ~10 ME very fast at sub-AU scales migration timescales for > ME very short at sub-AU scales no physical analog of “Solar System-like” in situ formation – these planets must have had significant gravitational interactions with the gas disk “in situ” formation from migration of ~fully-formed planetesimals formed planets from significantly from radially drifting larger orbital radii (e.g. solids (Chatterjee & Izidoro+ 17; Hands & Tan 2014) Alexander 18…) is there a timing problem? • Time = 0 for growth determined by collision velocities in young disks and (possibly) particle concentration in spiral arms and gravitational excitation of σ (Rice+ 06; Booth & Clarke 16) • rapid sequestration of large masses of solids into planetesimals is not just possible but expected è planetesimal formation action could be largely Ansdell+ 2017 complete during Class I is there a timing problem? • can massive enough planets form early enough and far enough out to explain observed disk structure? Levison+ 2010 • rapid planetesimal formation • growth in inner disk • planetesimal migration to observed scales modeling of HD 142527 Price+ 2018 Avenhaus+ 2017 • WYSIWYG disk – observed components (including a companion) generate observed structure • unquestionably some disks with large-scale structure contain planets • remains theoretical prior against most disks having planet-generated structures on > 10 AU scales questions? .