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Julia Venturini University of Zurich THE FORMATION OF GAS-RICH PLANETS Julia Venturini University of Zurich 10th RESCEU/Planet2 Symposium - Planet Formation around Snowline University of Tokyo, November 27-30, 2017 PLANETS IN THE SOLAR SYSTEM Credit: The International Astronomical Union/Martin Kornmesser GAS-RICH PLANETS IN THE SOLAR SYSTEM (credits: Wikiwand) Jupiter Saturn Uranus Neptune Mass [M⨁] 318 95 14.5 17 Radius [R⨁] 11 9.1 3.98 3.87 mass fraction of H-He ~87 - 95% ~ 68 - 91% ~10-25% ~10-25% EXOPLANETS: ~ 4000 CONFIRMED Credits: NASA planets with sizes between Earth and Neptune are VERY common DIVERSITY IN COMPOSITION FOR LOW-MASS PLANETS 8 80% H-He K-79d K-51b K-18d 7 K-89e HAT-P-26b K-87c 50% H-He 6 K-18c ] ⊕ K-223d 30% H-He 5 GJ 3470b K-223e K-89c WASP-47d 4 K-79c U N 10% H-He 5% H-He K-11d 3 100% H2O Planet radius [R 100 % MgSiO3 2 Earth-like Perovskite 1 (Venturini & Helled, 2017, ApJ) 0 5 10 15 20 25 Planet mass [M⊕] HOW DO PLANETS FORM? But disks don’t last forever… They are observed to have lifetimes of a few Myr Orion nebula. Credit: NASA, ESA, M. Robberto (STSI/ESA), the HST Orion Treasury Project Team and L. Ricci (ESO) TW Hydrae. Image credit: S. Andrews, Harvard-Smithsonian Center for Astrophysics / HL Tau. Credits: ALMA Credits: ALMA ALMA / ESO / NAOJ / NRAO. HOW ARE PLANETS FORMED INSIDE DISKS? Bottom-up: solids sediment and Top-down: disk formed around the star coagulate. When they grow large collapses under its own gravity. Planets form from the disk fragments (Disk Instability) — enough they can attract gas from the they are born gaseous. disk (Core accretion). Credit: Lucio Mayer & T. Quinn, ChaNGa code ThePLANET core-accretion FORMATION: model CORE ACCRETION MODEL Core formation Gas accretion by solid beyond critical accretion mass Protoplanetary disk 99% gas 1% solids Mcrit ~10 Earth masses — 7 Tdisk < 10 yr for H-He envelopes gas giant (Perri & Cameron 1974, Mizuno 1978, Bondeheimer & Pollack 1986, Pollack et al. 1996) PLANETPlanet FORMATION: formation: CORE core ACCRETION accretion model MODEL Core formation Core formation Gas accretion by planetesimals Slow core by solid beyond critical (“solids”) accretion accretion mass accretion Protoplanetary disk 99% gas 1% solids 7 Tdisk < 10 yr The core-accretion model gas giant Fine-tuning problem: disk must disappear when the planet acquired a non-negligible ice giant amount of H-He, but usually at this stage the planet is already in the runaway gas phase (Helled & Bodenheimer, 2014; Venturini & Helled, 2017). PLANET FORMATION: CORE ACCRETION MODEL Core formation Core formation notGas enoughaccretion by planetesimals Slow core by solid planetesimalsbeyond critical in (“solids”) accretion accretion feedingmass zone accretion Protoplanetary disk 99% gas 1% solids 7 Tdisk < 10 yr The core-accretion model gas giant ice giant terrestrial planet WHAT ARE THESE SOLIDS? Planetesimals-based model Pebbles-based model dust ⇒ ⇒ pebbles (~ cm size) ⇒ ⇒ high efficiency low efficiency planetesimals (~ km size) (courtesy of Y. Alibert) Planetesimals-based model Pebbles-based model 72 POLLACK ET AL. M. Lambrechts and A. Johansen: Forming the cores of giant planets 103 103 5AU 5AU 8AU 8AU 2 2 10 15AU 10 15AU 20AU 20AU 101 101 Z=0.02 E E /M /M c c M 100 M 100 10−1 (Lambrechts & 10−1 Z=0.005 (Pollack et al. 1996) Johansen, 2014) 10−2 10−2 5 6 7 5 6 7 ➤ Too long formation timescales unless planetesimals 10 10 10 10 10 10 t/yr t/yr are small (100 m - 1 km: ). Fortier et al. 2013 ➤ Fig.Formation 3. Core growth of asgas function giants of time. too The efficient di↵erent orbital, unless: separations Fig. 4. Core growth as function of time, for di↵erent values of the ini- ➤ Difficult to accrete planetesimals (Sho Shibata’s talk). ➡ (5high8 15disk20 viscosity AU) resemble (Ormel the compact 2017). orbital configuration expected tial dust-to-gas ratio Z0. Cores are placed on the same orbits as in Fig. 3, after− − disc− dissipation (Tsiganis et al. 2005). The red points identify the and similar labeling is used. Core growth is very sensitive to the initial ➤ Substantial H-He accretion onto a ~2 ME core requires ➡ criticallarge massorbital where distances a phase of + rapid high gas envelope accretion is triggeredopacities necessary metallicity: a twice as high value as the canonical dust-to-gas ratio of very special conditions ( Ikoma & Hori, 2012, for(Venturini the formation & Helled, of gas 2017 giants). (Lambrechts et al. 2014). Core growth Z0 = 0.01 leads to the formation of exclusively gas giants, while lower- ). comes to a halt after reaching the critical core mass. The dashed lines ing the metallicity by a factor 2 leads to systems of small ice giants. Bodenheimer & Lissauer 2014 represent an extrapolation ignoring this halt in accretion. The yellow 3 shaded area marks the dissipation of the gas disc after ⌧dis = 3 Myr. 10 5AU 8AU 2 Figure 3 illustrates how the growth of the core depends little 10 15AU on the separation from the host star. The initial embryo masses 3 5 20AU were taken to be 10− ME and inserted at a time ti = 10 yr. Embryo growth depends little on these assumptions4, as can be 101 seen in Eq. (35). The model parameters are the metallicity Z = E β0=1000 /M 2 c 0.01 and the initial gas surface density of 500 g cm− at 1 AU. M FIG. 1. (a) Planet’s mass as a function of time for our baseline model, case J1. InPlanets this case, within the approximatelyplanet is located 10 at AU 5.2 AU, reach the the initial critical surface core mass 100 density of the protoplanetary disk is 10 g/cm2, and planetesimals that dissolve duringand their trigger journey rapid through gas accretion the planet’s (the envelope red dot are marks allowed the to pebble sink to the planet’s core; other parameters are listed in Table III. The solid line representsisolation accumulated mass), while solid planets mass, at the wider dotted orbits, line which accumulated do not reach 5 −1 gas mass, and the dot–dashed line the planet’s total mass. The planet’s growth occursM iniso three, are stranded fairly well-defined as ice giants. stages: During the first p5 3 10 10 β0= 250 years, the planet accumulates solids by rapid runaway accretion; this ‘‘phase 1’’ ends whenCore the growth planet is has highly severely sensitive depleted to its the feeding metallicity, zone of as can ˙ ˙ 6 planetesimals. The accretion rates of gas and solids are nearly constant with MXY P be2–3 seenMZ during from most Eq. ( of35 the). Figurep7 3 104 showsyears’ the duration di↵erence of phase between −2 2. The planet’s growth accelerates toward the end of phase 2, and runaway accumulationan initial of gas dust-to-gas (and, to a lesser ratio extent, of Z solids)= 0.005 characterizes and Z = phase0.02 (while 10 3. The simulation is stopped when accretion becomes so rapid that our model breakskeeping down. The other endpoint parameters is thus fixed). an artifact The evolution of our technique of the and core mass 105 106 107 should not be interpreted as an estimate of the planet’s final mass. (b) Logarithm ofis the similarly mass accretion sensitive rates to of the planetesimals choice of (solid the (initial) line) and gas gas surface t/yr (dotted line) for case J1. Note that the initial accretion rate of gas is extremely slow,density. but that Figure its value5 illustrates increases rapidly the growth during of phase planetary 1 and early cores for a Fig. 5. Core growth as function of time, for two values of the initial phase 2. The small-scale structure which is particularly prominent during phase 2 isgas an surface artifact densityproduced half by ourand method double of that computation of our standard of the choice added gas mass from the solar nebula. (c) Luminosity of the protoplanet as a function of time for case2 J1. Note the strong correlation between gas surface density β0, which has been altered by a factor of two from (β = 500 g cm− ), both for the cases with exponential gas dissi- 2 0 the standard value used here (β0 = 500 g cm− ). Cores are placed on the luminosity and accretion rate (cf. b). (d) Surface density of planetesimals in the feedingpation zone over as a function time and of without. time for case Wediscuss J1. Planetesimals the sensitivity become of plan- substantially depleted within the planet’s accretion zone during the latter part of phase 1, and the local surface density of planetesimals remains same orbits as in Fig. 3, and similar labeling is used. The grey lines give etary growth to the metallicity and gas surface density in more small throughout phase 2. (e) Four measures of the radius of the growing planetary embryo in case J1. The solid curve shows the radius of the the evolution in a disc with a temporally constant gas surface density detail in Sect. 5.4. profile, corresponding to Eq. (35). planet’s core, Rcore , assuming all accreted planetesimals settle down to this core. The dashed curve represents the effective capture radius for planetesimals 100 km in radius, Rc . The dotted line shows the outer boundary of the gaseous envelope at the ‘‘end’’ of a timestep, Rp . The long- Here, c is a parameter that depends on the radial pressure and and short-dashed curve represents the planet’s accretion radius, Ra . 4. Planetary migration temperature structure of the protoplanetary disc.
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