planet Jovian formation. Core-accretion or gravitational instability?
Ge/Ay133
PropertiesoftheJovianPlanetsintheSolarSystem
P 2 forH2He
I/MR2=0.4forauniformsphere I/MR2=0.26forP 2
Theradiusmass relationshipandM.o.I. areusedtoinferthe presenceofprimordial coresof1030Mearth.
1113
Caveat!CoremassestimatebasedonhighpressureEOS:
[preferred EOS]
OKforSaturn,but
[envelope]
Saumon & Guillot 2004 core mass constraints based on EOS
verylargeextrapolations&uncertaintiesforJupiter! (needbetterhighP,Tmeasurements,verydifficult)
dubious EOS? Previously favored. Currently preferred EOS (Boriskov et al. 2005)
[envelope]
Saumon & Guillot (2004) core mass constraints based on EOS
Theoryofnucleatedinstability:
CoresinJovianplanetsarealmostcertainlyprimordial,andthefactthatall suchobjectsinthesolarsystemradiatemoreenergythantheyreceivemeans theystartedhot.Thishasledtothedevelopmentofthecoreaccretionmodel inwhichgasaccretesontocoresbuiltalongthelinesdiscussedinLec.#12.
Densecore rH Ambient solarnebula
~Isothermal Photosphere
Adiabatic envelope
HillSphere
Theoryofnucleatedinstability:
Howdoweanalyzethissituation?Theextentoftheenvelopeisdetermined viahydrostaticequilibrium.Keyisthetemperatureprofile,whichis establishedbytheradiativetransferequationsbelow.Listheluminosity,Kis themassopacitycoefficient.
Densecore rH Ambient solarnebula
Adiabatic envelope
Photosphere HillSphere
Theminimumluminositythatneedstoberadiatedisthatwhichbalancesany ongoingaccretion(equationatleft).Fromthisthemass/densitypropertiesofthe envelopecanbeestimated:
Densecore rH Thusweneedtosolveforthe densitystructuretogetthe envelopemass,whichmeans weneedtoknowthe temperatureprofile.
Adiabatic envelope
Photosphere HillSphere
Solvingtheradiativetransferequationsyields(fortheenvelope):
Idealgas
Densecore rH ClearlythevalueofKis critical.Gascanonly contributeasmallfractionof theoverallopacity,andso thedustgrainoricecontent intheenvelopemustbe known,orassumed... HillSphere
Adiabatic envelope
Photosphere
Howmassivedoesthecoreneedtobefortheatmospheretocollapse?
f~Kincm2/g SettingdMc/dMt=0gives( 4 )
Densecore Photosphere
Adiabatic envelope
Stevenson1982,Pl.Sp.Sci.30,755
Thegas/dustratiointheenvelopeisalsocriticalforTIMESCALES! (determineshowrapidlytheenvelopecancool)
ISM/50ISMdust/gasSmaller core
Lissauer2001,Nature409,23
Themostrecentsimulationsincludedustgrain/heavyelementsettlinginthe enveloperatiointheenvelopetogive~23Myrtimes:
ISM/50ISMdust/gasSmaller core
Lissaueretal.2009,arXiv:0810.
Ifthegasinflowiscoherent,lotsofangularmomentumisinvolved:
2 /rd3~GMp/rd2 whereGisthespecificang.mom. andrdisthe(protoplanet)diskradius. Equatingthistotheorbitalspecific ang.momentumgives,roughly ~rH2/4orrd~20rplanet
Whatmightsuchaprotoplanetarydisktellusabouttheformationof satellites?Foradetailedrecentreview,see: Estrada,P.R.etal.2009,arXiv:0809.1419
PropertiesoftheinnermoonsofJupiter:
1000 T(K) 500
Hydratedsilicates
250 125
Waterice
AmmoniaWaterHydrate
Solarnebula (buffer)
102030RJ IoEuropaGanymedeCallisto 3.53.11.921.78gcm3 AnhydrousHydrated60/40rock/ice50/50rock/ice
silicatessilicates(initially)
Saturnpicturenotsoclear,butTitanslocationmayexplainlargevolatilecontent.
Comparisonofprotosolarversusprotoplanetarydisks: PropertyProtoplanetaryDiskProtosolarDisk
Size(centralbodyunits)~20~103104 Mass(centralbodyunits)0.050.050.1 TypicalTemperature~200K~200K (butupto2000K)(butupto>1000K) Verticalopticaldepth~100(gasalone)<<1(gasalone) ~10,000(withdust)~100(withdust) Masssurfacedensity(g/cm2)~105(gas)102103(gas) ~103(solids)110(solids) Gasdensity(g/cm3)10410610101012 Gaspressure(bars)~1~106 Viscousspreadingtime100yr~105106yr Coolingtime~104106yr~100104yr
vs.
Giant Planet Formation: Theory vs. Observations
AlanP.Boss TheFormationof CarnegieInstitutionofWashington PlanetarySystems
HereticsApproachto SolarSystem FormationFForm Molecules,MicrobesandtheInterstellarMedium
GeophysicalLaboratorysWesFest CarnegieInstitutionofWashington October26,2007
Outline:
ConventionalscenarioforSolarSystemformation: regionoflowmassstarformation(Taurus) collisionalaccumulationofterrestrialplanets formationofgiantplanetsbycoreaccretion HereticalscenarioforSolarSystemformation: regionofhighmassstarformation(Orion,Carina) collisionalaccumulationofterrestrialplanets formationofgiantplanetsbydiskinstability Observationalteststodiscriminatebetweenthesetwo formationmechanismsforgiantplanets?
3
Extrasolar Gas Giant Planet Census: Frequency
* Approximately 15% of nearby G dwarfs have gas giant
planets with relatively short orbital periods hot and warm Jupiters (Hatzes 2004) * Approximately 25% of nearby G dwarfs appear to have gas giant planets with even longer orbital periods Solar System analogues (Hatzes 2004) * Hence as many as 40% of nearby G dwarfs appear to have gas giant planets inside about 10 AU (Hatzes 2004) * Approximately 20% of FGK dwarfs have giant planets with orbital distances less than 20 AU (Marcy 2007) * More massive stars (up to 1.9 Msun) have more gas giant planets than lower mass dwarfs (Marcy 2007) * Using either set of statistics, gas giant planet formation mechanism must be relatively efficient and robust
Cieza et al. 2006 SST survey: ~65% of disks gone in < 1 Myr
Gravitational Instabilities (GIs) in disks, can rapid planet formation result?
Compact, massive disks are susceptible to clumping: 1.0 Msun protostar with a 20 AU radius disk of mass 0.09 Msun
Boss (2003) disk instability model after 429 yrs, 30 AU radius GIclumpsformrapidly, thekeyquestionsabout planetformationare whethersuchclumps cancoolefficiently enoughtocontinuetheir contractionorwhether theybounceandthus dissipateMuchlike envelopecollapsein coreaccretionmodels. ThisapproachisFAST, however,butneeds compact&massive disks.
Inaba, Wetherill, & Ikoma (2003) core accretion Critical model
mass for onset of gas accretion
* first model which included effects of planetesimal fragmentation and loss by orbital migration as well as capture by protoplanets gas envelope * 21 Earth-mass core forms at 5.2 AU in 3.8 Myr * no Saturn formed * disk mass = 0.08 solar masses
Helled et al. 2006
accreted mass log radius/RJupiter ~36 MEarth
Time in units of 105 yrs
GImodelscangeneratesubstantial heavyelementcores,ifthereis substantialdustsettlingbeforethe instabilitiesleadtocollapse.
Time in units of 105 yrs
Anewparadigmforformingthegiantplanetsrapidly: Marginallygravitationallyunstableprotoplanetarydiskforms fourormoregiantgaseousprotoplanetswithinabout1000 years,eachwithmassesofabout1/3to1Jupitermasses Dustgrainscoagulateandsedimenttocentersofthe protoplanets,formingsolidcoresonsimilartimescale,with coremassesofnomorethanabout6Earthmassesper Jupitermassofgasanddust(Z=0.02) DiskgasbeyondSaturnsorbitisremovedinamillionyears byultravioletradiationfromanearbymassivestar(Orion, Carina,)
Continued
OutermostprotoplanetsareexposedtoFUV/EUV radiation,whichphotoevaporatesmostoftheirenvelope gasinaboutamillionyearsorless Outermostplanetsgasremovalleadstoroughly15 Earthmasssolidcoreswiththingasenvelopes:Uranus, Neptune InnermostprotoplanetisshelteredbydiskHgas gravitationallyboundtosolarmassprotosunandso doesnotloseanygas:Jupiter Protoplanetattransitionalgaslossradiuslosesonlya portionofitsgasenvelope:Saturn TerrestrialplanetregionlargelyunaffectedbyUVflux [TPF/Darwintargets]
Discovery space around M discoveries added Discovery space with planets with latest (K?) dwarf stars highlighted
GJ 876 GJ 317 GJ 849 GJ 876
OGLE-2003-BLG-235 OGLE-2005-BLG-071
OGLE-2006-BLG-109b,c
GJ 436 GJ 581 GJ 176 OGLE-2005-BLG-169
GJ 876
OGLE-2005-BLG-390
Laughlin et al. 2004 core accretion models
1.0 Msun total core
* gas giants rarely form by core accretion around M dwarfs: process too slow
0.4 Msun
total core
SufficientlymassivedisksaroundlowmassstarsdoshowGIs:
0.5 solar mass star with a 20 AU radius disk after 215 yrs (Boss 2006) Jupiterand/orsuperEarth formationaroundKstars?
Heretical Explanation for Long-Period Super-Earths
Most stars form in regions of high-mass star formation (e.g., Orion, Carina) where their protoplanetary disks can be photoevaporated away by nearby O stars. Photoevaporation converts gas giant protoplanets into ice giants if the protoplanet orbits outside a critical radius, which depends on the mass of the host star. For solar-mass stars, the critical radius is > 5 AU, while for a 0.3 MSun M dwarf star, the critical radius is > 1.5 AU. If M dwarfs have disks massive enough to undergo disk instability, then their gas giant protoplanets orbiting outside ~1.5 AU will be photoevaporated down to super-Earth mass, for M dwarfs in regions of high-mass star formation. In low-mass star formation regions (e.g., Taurus), their gas giant protoplanets will survive to become gas giant planets.
Giant Planet Census: Host Star Metallicity
Correlation of short-period Jupiters with stellar metallicity is usually attributed to formation by core accretion RV searches are beginning to find planets around low [Fe/H] dwarfs (HD 155358: [Fe/H] = -0.68 has two planets with masses of 0.5 and 0.9 MJup, Cochran et al. 2007; HD 171028: [Fe/H] = -0.49 has one with 1.8 MJup, Santos et al. 2007) Most M dwarfs with known planets (GJ 176, GJ 876, GJ 317, GJ 436, GJ 581) have metallicities less than solar: [Fe/H] = -0.1, -0.12, -0.23, -0.32, and 0.33, while only GJ 849 has [Fe/H] = +0.16 (Butler et al. 2006) Short-period SuperEarths do not correlate with the host stars [Fe/H] (Mayor 2007) Low [Fe/H] giant stars have more (long-period) gas giants than high [Fe/ H] giant stars (Hatzes 2007) M4 globular cluster has [Fe/H] ~ -1.5, yet pulsar B1620-26 has a giant planet with a mass ~ 2.5 MJup (Sigurdsson et al. 2003) Core accretion cannot work as [Fe/H] drops to low values
Mayer et al. (2007) 3D SPH with radiative transfer, convection, fragmentation
= 2.4 3000
= 3.0 6000
Fragments for higher mean molecular weight and larger radiating surface area
= 2.4 4000
= 2.7 4000
Core Accretion Mechanism
Pro: Leads to large core mass, as in Saturn Higher metallicity may speed growth of core Based on process of collisional accumulation, the same as for the terrestrial planets Does not require external UV flux to make ice giants, so works in Taurus HD 149026: 70 Earth-mass core plus 40 Earth-mass gaseous envelope? Formed by collision between two giant planets (Ikoma et al. 2006)? Failed cores naturally result
Con: Jupiters core mass is too small? If gas disks dissipate before critical core mass reached failed Jupiters result Difficult to form gas giant planets for M dwarfs, low metallicity stars (e.g., M4), or rapidly (CoKu Tau/4?) Loss of growing cores by Type I migration? Needs disk mass high enough to be ~ gravitationally unstable No in situ ice giant formation?
Disk Instability Mechanism
Pro: Can explain core masses, bulk compositions, and radial ordering of gas and ice giant planets in Solar System Requires disk mass no more than that assumed by core accretion Forms gas giants in either metal-rich or metal-poor disks (M4) Clumps form quickly (CoKu Tau/4?) even in short-lived disks Works for M dwarf primaries Sidesteps Type I (and III) orbital migration danger Works in Taurus or Orion, implying Solar System analogues are common
Con: Requires efficient cooling of midplane (e.g., convection), coupled with efficient cooling from the surface of the disk: subject of work in progress Clump survival uncertain: need for models with detailed disk thermodynamics and higher spatial resolution (e.g., AMR) Requires large UV dose to make ice giant planets in Taurus would make only gas giant planets