How asteroids grow Anders Johansen (Lund University) Star and - - PowerPoint PPT Presentation

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How asteroids grow Anders Johansen (Lund University) Star and Planet Formation For All, Lund, February 2014 1 / 14 Planets and exoplanets First exoplanet was discovered in 1995 ( Mayor & Queloz , 1995) The Kepler satellite

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How asteroids grow

Anders Johansen (Lund University)

“Star and Planet Formation For All”, Lund, February 2014

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Planets and exoplanets

◮ First exoplanet was discovered in 1995 (Mayor & Queloz, 1995) ◮ The Kepler satellite identified over 2300 exoplanet candidates

in the 16-months data (Batalha et al., 2013) ⇒ 50% of stars have planets within 0.4 AU (Fressin et al., 2013) ⇒ Most exoplanets in close orbits are super-Earths or small Neptunes ⇒ Nature is very efficient at turning dust and ice into planets

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Classical picture of planet formation

Planetesimal hypothesis of Viktor Safronov 1969: Planets form in protoplanetary discs around young stars as planetesimals collide to form ever larger bodies

1. Dust to planetesimals

µm → km: contact forces

2. Planetesimals to protoplanets

km → 1,000 km: gravity (run-away accretion)

3. Protoplanets to planets

Terrestrial planets: protoplanets collide (107–108 years) Gas and ice giants: 10 M⊕ core accretes gas (< 106...7 years) Severe problems with classical model: 1 Growth of macroscopic particles is frustrated by erosion and bouncing 2 Planetesimals colliding at high speeds will shatter each other 3 Core formation takes much longer time than the life-time of the nebula

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Planet formation with pebbles

Pebble hypothesis: Planetesimals form by gravitational collapse of dense clumps of pebbles and planets form mainly by pebble accretion onto planetesimals

1. Dust to pebbles

µm → cm: coagulation and condensation

2. Pebbles to planetesimals

km → 100–1,000 km: particle concentration and gravitational collapse

3. Planetesimals to planets

Terrestrial planets: pebble accretion, giant impacts (106–108 years ?) Gas and ice giants: pebble accretion to 10 M⊕ (≪ 106 years)

See Protostars and Planets VI reviews by Johansen et al. (2014) and Chabrier, Johansen, et al. (2014)

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Dust to pebbles

◮ Collisions between dust

aggregates can lead to sticking, bouncing or fragmentation

(G¨ uttler et al., 2010)

◮ Sticking for low collision speeds

and small aggregates

◮ Bouncing prevents growth beyond

mm sizes (bouncing barrier)

◮ Further growth may be possible

by mass transfer in high-speed collisions (Windmark et al., 2012) or by ice condensation (Ros & Johansen, 2012)

→ SPFFA talk by Katrin Ros (Zsom et al., 2010)

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Pebbles to planetesimals

v η (1− )

Kep

F FG

t=40.0 Ω−1

−20.0 −10.0 +0.0 +10.0 +20.0 x/(ηr) −20.0 −10.0 +0.0 +10.0 +20.0 z/(ηr)

t=80.0 Ω−1

−20.0 −10.0 +0.0 +10.0 +20.0 x/(ηr) −20.0 −10.0 +0.0 +10.0 +20.0 z/(ηr)

t=120.0 Ω−1

−20.0 −10.0 +0.0 +10.0 +20.0 x/(ηr) −20.0 −10.0 +0.0 +10.0 +20.0 z/(ηr)

t=160.0 Ω−1

−20.0 −10.0 +0.0 +10.0 +20.0 x/(ηr) −20.0 −10.0 +0.0 +10.0 +20.0 z/(ηr)

◮ The radial drift flow of particles is linearly

unstable to streaming instability

(Youdin & Goodman, 2005; Youdin & Johansen, 2007)

◮ Particles pile up in dense filaments

(Johansen & Youdin, 2007; Bai & Stone, 2010a)

◮ Particle concentration triggered above a

threshold metallicity around Z ≈ 0.015

(Johansen et al., 2009, 2012; Bai & Stone, 2010b,c)

◮ Possible to concentrate particles down to mm

sizes at 2.5 AU (Carrera, Johansen, & Davies, in prep)

→ SPFFA talk by Daniel Carrera

−0.10 −0.05 0.00 0.05 0.10 x/H −0.10 −0.05 0.00 0.05 0.10 y/H −0.10 −0.05 0.00 0.05 0.10 x/H −0.10 −0.05 0.00 0.05 0.10 y/H −0.10 −0.05 0.00 0.05 0.10 x/H −0.10 −0.05 0.00 0.05 0.10 z/H

t = 337.5 Ω−1 Z = 0.020

325 330 335 t/Ω−1 10−8 10−7 10−6 10−8 10−7 10−6 µ 58 97 162 269 R/km

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Planetesimals to planets

−0.10 −0.05 0.00 0.05 0.10 y/H 0.0 5.0 Σp/<Σp> −0.10 −0.05 0.00 0.05 0.10 x/H −0.10 −0.05 0.00 0.05 0.10 y/H −0.05 0.00 0.05 0.10 x/H x/H −0.10 −0.05 0.00 0.05 0.10 z/H −0.10 −0.05 0.05 0.10 t=0.0 Ω−1 t=120 Ω−1 t=131 Ω−1 t=134 Ω−1

Core growth to 10 M⊕

10−1 100 101 102 r/AU 103 104 105 106 107 108 ∆t/yr Pebbles F r a g m e n t s Planetesimals

⇒ Pebble accretion speeds up core formation by a factor 1,000 at 5 AU and a factor 10,000 at 50 AU

(Lambrechts & Johansen, 2012; also Ormel & Klahr, 2010; Morbidelli & Nesvorny, 2012)

⇒ 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 → SPFFA talk by Michiel Lambrechts

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Evidence for giant impact stage

(Wetherill, 1985)

◮ The Moon’s mean density is very low, with uncompressed density

ρ = 3.3 g cm−3 [Earth’s uncompressed density: ρ = 4.4 g cm−3]

◮ The Moon is highly differentiated – with a dense core, a mantle, and a

crust – but must be lacking iron and volatiles ⇒ Moon formed from the impact debris after Mars-sized protoplanet impacted the young, differentiated Earth ⇒ Taken as evidence for giant impact stage of classical planet formation ? Any evidence for pebble accretion? YES – encoded in the asteroid sizes

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Asteroid birth sizes

Asteroid size distribution

100 1000 D [km] 10−3 10−2 10−1 100 101 102 103 dN/dR [km−1]

◮ Size distribution of asteroids shows distinct bumps at diameters D = 120

km and D = 350 km

◮ Forming asteroids from km-sized planetesimals does not reproduce the

first bump – bump is primordial (Bottke et al., 2005)

◮ Asteroids must be born BIG (100 – 1000 km) in order to not overproduce

asteroids with diameters less than 100 km (Morbidelli et al., 2010)

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Planetesimal formation

102 103 104 105 106 107 108 109 dN/dM [M22

−1]

25 50 100 200 400 R [km] 1020 1021 1022 1023 1024 M [g]

2563, 1.0×MMSN 2563, 2.5×MMSN 5123, 5.0×MMSN 2563, 5.0×MMSN 1283, 5.0×MMSN

qM = 1.31 +/− 0.07

◮ Streaming instability leads to concentration of pebbles and to

planetesimal formation (Johansen et al., 2014, Protostars and Planets VI, arXiv:1402.1344 )

◮ Higher resolution gives smaller planetesimals (PRACE grant “PLANETESIM”) ◮ Birth sizes of planetesimals show no sign of a bump – most of the

planetesimals are small but most mass is in the largest bodies

◮ Powerlaw in dN/dM ∝ M−q is approximately q = 1 . . . 1.5 ◮ Gravitational collapse of clumps → SPFFA talk by Kalle Jansson

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Chondrules

◮ Meteorites recovered on Earth are fragments of asteroids ◮ Oldest condensates in the Solar System are CAIs with a

narrow age range of 4567.30 ± 0.16 Myr (Connelly et al., 2012)

◮ Primitive meteorites (chondrites) contain a large fraction of

0.1-1-mm-sized chondrules (formed over the first 3 Myr)

◮ Chondrites contain up to 80% of their mass in chondrules ◮ What role did chondrules play in asteroid formation?

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Chondrule accretion

asteroid due to gas friction by protoplanet Large chondrule is scattered Chondrule spirals towards

∆v ≈ 50 m/s Bondi radius: RB =

GM (∆v)2

˙ M = πR2

Bρc∆v ∝ R6

(Lambrechts & Johansen, 2012)

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Planetesimal size distribution

10−1 100 101 102 103 104 105 106 dN/dR [km−1] Nominal model 10 100 1000 R [km] Lower pressure support 10 100 1000 R [km] 10−2 10−1 100 101 102 103 104 105 dN/dR [km−1] 10 100 1000 R [km] 10−2 10−1 100 101 102 103 104 105 dN/dR [km−1] Steeper chondrule size distribution 10 100 1000 R [km] Larger chondrules 10−2 10−1 100 101 102 103 104 105 dN/dR [km−1]

◮ The nominal model reproduces four features of the asteroid size

distribution: the bump at R = 60 km, the steep size distribution up to R = 200 km, the bump at R = 200 km and the shallow size distribution for the largest sizes (Johansen, Mac Low, Lacerda, & Bizzarro, in prep)

◮ Variation in the parameters gives different realisations of the asteroid belt

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Implications

(Elkins-Tanton et al., 2011)

◮ Asteroids grew primarily by chondrule accretion ◮ Size distribution of asteroids shows evidence of this chondrule accretion ◮ General validation that pebble accretion occurred in the Solar System ◮ Pebble accretion likely driven by icy pebbles beyond the ice line ◮ Planetesimals in the terrestrial planet formation region grew by accreting

chondrules – could this explain rapid formation of Mars?