The atmosphere of Exoplanets AND Their evolutionary properties. I. Baraffe

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1 The atmosphere of Exoplanets AND Their evolutionary properties I. Baraffe

2 I) Properties of cool atmospheres: 1) Atmospheric chemistry 2) Main opacity sources 3) Non solar composition 4) Non equilibrium chemistry II) Irradiation effects III) Evaporation process IV) Evolutionary properties 1) The cooling of planets 2) Theory versus observations: transit planets

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4 1) Atmospheric chemistry Atmosphere: thiny region (in mass and radius) at the surface where photons escape -----> optically thin region where diffusion approximation is not valid anymore -----> modelling decoupled from inner structure calculation Model construction: - Solve the radiative transfer equation (in 1D i.e plane parallel geometry) - Equation of state: perfect gaz - requires a good knowledge of chemical composition and wavelength dependent opacities For a given chemical composition, an atmosphere model is characterised by - The effective temperature T eff: (temperature of the black body which yields the same surface flux of energy as the object) L = 4!R 2!T eff 4 - The surface gravity: g = GM/R 2

5 Atmospheric abundances (T eff < 2000 K, log g < 3.5) - atomic H - (T > 4000 K) - molecules (T < 4000 K): H 2, H 2 O,TiO, CO, CH 4, NH 3 abundances of molecules are calculated using chemical equilibrium codes - Condensed species (liquid/solid «grains» or «dust») (T < 2000 K) * condensation process sequesters most of the heavier elements (Si, Mg, Ca, Al,...) in complex compounds (MgSiO 3, CaTiO 3, Al 2 O 3,...) * condensation of water H 2 O and ammonia NH 3 (T < 300 K) > complex processes of gravitational settling and cloud layers formation Note : observations of brown dwarfs teach us a lot about these processes, in particular since the discoveries of - L-dwarfs (T eff ~ K) ---> «dusty» atmospheres - T- dwarfs (T eff < 1600 K) ----> presence of methane CH 4

6 Fraction of major gas phase species in an atmosphere model of Teff = 950 K, log g=3 Burrows et al. 1997

7 Condensation curves (solid lines): constituent condensate elements are removed from the gas above these curves Lodders & Fegley 2006

8 Formation and settling of «dust» in brown dwarf atmospheres from M ---> L -----> T dwarf Teff 2200K 1800K 1000K Baraffe et al. 1998,2003; Chabrier et al. 2000, Allard et al Marley et al. 2000, 2002; Burrows et al. 2003, 2006

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11 2) Main sources of opacity where - Gaseous molecular opacities: main absorbers in the optical (" ~ 0.5 #m) ----> TiO, VO main absorbers in the near-ir (" ~ 1-2 #m) ----> H 2, H 2 O, CH 4 - Dust opacities: $ " % =!! i & a n i (a)q ext (a,i, ")a 2 da n i (a): number of dust particles of species i of size a ---> depends upon chemical equilibrium of species i and size distribution of dust particles large uncertainty on dust sizes: requires assumptions for their distribution. Common assumptions: interstellar medium distribution (submicron size) or between ~ 0.1 #m to ~ 10 #m

12 Q ext (a,i, "): extinction efficiency > Mie Theory: describes the analytical solution of Maxwell s equations for the scattering of electromagnetic radiation by spherical solid particles with a given refractive index Allard et al. 2001

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15 3) Non solar composition Giant planet atmospheres are expected to be enriched in heavy elements, as observed in Jupiter and Saturn (inherited during planetesimal accretion as the planet formed): Jupiter: in situ measurement from Galileo enrichment by a factor 2-4 Tropospheric element abundances Saturn: spectroscopic determination C (CH4) and N (NH 3 ) significantly enriched Guillot 2005

16 Signatures of non solar metallicity not obvious to find Effect of an increase of metallicity (factor 5) on spectra Barman et al. 2006; Chabrier et al. 2006!Important topic: possible dinstinction between brown dwarfs (parent star metallicity) and exo-planets (because of different formation process)

17 4) Non equilibrium chemistry Common assumption of local chemical equilibrim (LCE). But, if some chemical reactions are very slow -----> vertical transport via convective motions can lead to departure from equilibrium Mechanism suggested to operate in Jupiter in 1997 (Prinn & Barshay) and expected as well in exoplanet atmospheres Non equilibrium carbon chemistry: main reaction CO + 3H 2 < > CH 4 + H 2 O below ~ 2000 K, CH 4 becomes the dominant form of C Transformation CO ----> CH 4 much slower than inverse reaction ' if ( mix << ( CO "CH4 ' abundance of CO much larger than LCE predictions

18 ) existence of this process confirmed by the detection of CO in the atmosphere of a cool brown dwarf GL 229b (T eff ~ 1000 K) Non equilibrium nitrogen chemistry: same process expected for N: N 2 + 3H 2 < > 2NH 3 reaction N > NH 3 much slower than inverse reaction

19 Effect of non chemical equilibrium 1600 K 800 K Saumon et al. 2003

20 II) Irradiation effects The increasing number of close-in planets (orbital separation a < 1 AU) ' important to account for irradiation effects from the parent star Model atmosphere: ----> Atmosphere models including the incident flux of the parent star in the solution of the radiative transfer equation (Barman et al. 2005; Sudarsky et al.; Hubeny et al. ; Marley et al.; Fortney et al.) R * : radius of the star F * : flux of the strar (=!T eff4 ) f: redistribution factor F inc = f/4 (R * /a) 2 F * f= > heat redistributed over the entire planet f= > heat redistributed only over the day-side of the planet

21 Common simplification: impinging radiation field is isotropic (attempts to take angle dependence by Barman et al. 2005) Evolution of the planet toward an equilibrium temperature A: the Bond Albedo A. F inc = F reflected (1-A). F inc = F absorbed T 4 eq = (1-A)/! F inc For Jupiter: A ~ 0.35 For typical close-in planets: A~0.1

22 Irradiation effects on atmosphere profiles irradiated non-irradiated Planet of intrinsic T eff = 100K, irradiated by a G-type star (a Sun) at a=0.046au irradiated non irradiated (Barman et al. 2001)

23 Irradiation effects on spectra: Spectrum of irradiated planets with intrinsic T eff =100K, at a=0.023 AU and a=0.046 AU of a sun $ T eq~2400k # T eq ~1700K F out =!T 4 eff +(1-A)F inc + AFinc Bond albedo A ~ 0.1: F refl =AF inc (Barman et al. 2001; Chabrier et al. 2004)

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26 IV) Evolutionary properties 1) The cooling of planets 2) Theory versus observations: transit planets

27 1) The cooling of planets Evolution characterised by contraction and cooling (no other energy source):. Rate of energy released: L(t) = db/dt. Variation of binding energy: d+,t-.-[d/(t) + du(t)] = & M Gmd(1/r)dm - & M du dm ----> EOS crucial for mechanical structure (radius R for a given mass and age) -----> Atmospheric properties are crucial for the evolution: R(t), L(t) ' Coupling between inner structure models (M, R, %(r), P(r), L(r)) and atmospheric models (T eff, g) which provide the outer boundary conditions to the inner structure model (%(R), P(R))

28 Irradiation effects slow down the cooling ' 10%-15% increase of the radius at a given age Evolution of a 1 M J gaseous planet -Irradiation by a Sun at 0.05 AU - No irradiation effects Baraffe et al. 2003

29 2) Theory versus observations: Transit planets Effect on evolution: radius larger at a given time (10%-15%) > success to explain the radius of some transit planets

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31 But still large uncertainty on the precise amount of heavy material and its composition because of uncertainties on: (i) EOS of heavy materials (water, rock, iron, etc) at conditions found in planetary interiors (P > Mbar, T > 5000 K) EOS of water can be probed by laboratory experiments up to P ~ 0.3 Mbar and T ~ 2000K. Above that, extrapolations are necessary. (ii) The distribution of heavy elements (everything in a heavy core or distributed over the entire planet?)

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