Ecole Normale Supérieure, 22 March 2011, LA-UR Images: Cassini; Marois et al. (2008) D. Saumon Los Alamos National Laboratory Clouds in brown dwarfs and hot young planets

Modeling Mark Marley (NASA Ames) Katharina Lodders (Washington U.) Richard Freedman (NASA Ames) Observations/Analysis Michael Cushing (IPAC) Sandy Leggett (Gemini Observatory) Tom Geballe (Gemini Observatory) Adam Burgasser (UCSD) Denise Stephens (Brigham Young) Spitzer/IRS Dim Suns Team: Tom Roellig, Amy Mainzer, John Wilson, Greg Sloan, Davy Kirkpatrick, Jeff Van Cleve

Brown dwarfs and hot young planets Clouds and the L  T transition Back to hot young planets

The basics of brown dwarfs Main sequence stars Brown dwarfs Planets Substellar in mass ~ 12 to 77 M Jupiter Compact! R ~ 0.09 to 0.16 R sun (all ~ size of Jupiter) No stable source of nuclear energy Evolution = cooling: L bol and T eff decrease with time

The basics of brown dwarfs Two new spectral classes cooler than dM: L (T eff ~ 2400  1400K) T (T eff ~ 1400  600K) … and soon Y Very strong molecular bands H 2 O, CO, CH 4, NH 3, FeH, TiO Condensates form in the atmosphere for T eff  2000K

M6.5 L5 T4.5 From Cushing et al. (2006) Spectral energy distributions: M, L and T dwarfs

NIR color-magnitude diagrams: Field brown dwarfs Data from Knapp et al (2004), Burgasser et al. (2006), Liu & Leggett (2005) L dwarfs naturally extend the dM sequence T dwarfs are “blue” in near-IR colours! Rapid shift in IR colours at the L  T transition. Early T’s are brighter in J than late L’s!? Sample of local disk brown dwarfs

Hot young planets HR 8799 b,c,d,e Star: A5 V 30  160Myr old Inner & outer dust disks (R ~ 8AU and R ~ 66AU) bcde a (AU) Mass (M J )* T eff (K)*~900~1100 ? Marois et al. (2008, 2010), Currie et al. (2011) * Masses and T eff are approximate

Near IR color-magnitude diagrams: Young planets Planet data from Marois et al. (2008), Lafrenière et al. (2008) and Neuhäuser (2008) In the near IR, hot young planets ~ consistent with the L dwarf sequence and its extension

Hot young planets compared to field brown dwarfs Similarities Young planets far from their star have effectively evolved in isolation Their NIR colours extend the brown dwarf sequence T eff range (L~ K, T~1400 – 600K) Differences Lower gravity (lower mass, younger) Possibly metal-rich if formed in a protoplanetary disk Hot young planets are more hip Atmospheric physics & chemistry should be very similar

Brown dwarfs and hot young planets Clouds and the L  T transition Back to hot young planets

Color-magnitude diagrams: Limiting cloud models Cloudless models can explain late T dwarfs (>T4) A diffuse cloud model works best for optically thin clouds (<L4)  A more sophisticated cloud model is required (cloud deck) condensation level diffuse cloud cloud deck no cloud

Condensation and qualitative cloud behaviour Condensation: Fe (T eff ~ 2000K) Silicates (T eff ~ 1600K) H 2 O (T eff ~ 300K) Main chemical transitions: CO  CH 4 (T eff ~ 1200K) N 2  NH 3 (T eff ~ 800K) photosphere T eff 2000K 1600K 1200K 800K 400K } } L T A cloud layer will gradually disappear below the photosphere as T eff decreases  cloudless atmosphere!

A simple cloud model A minimalist 1-D cloud model results from a balance between: 1) Vertical transport of particles (e.g. turbulent diffusion) 2) Gravitational settling of the particles K zz : coefficient of diffusive mixing f sed : dimensionless sedimentation parameter ( as f sed increases, cloud deck becomes thinner) The location of the bottom of the cloud is determined by the condensation curve Ackerman & Marley (2001)

Color-magnitude diagrams: Models with clouds For L dwarfs, f sed ~ 1-2 The L  T transition occurs over T eff ~ 1400  1200K Transition can be accounted for by an increase of f sed At constant f sed, the cloud layer gradually disappears below the photosphere

The L  T transition as cloud clearing The L  T transition can also be explained as a gradual clearing of the cloud Transition from f sed =2 to cloudless models. Marley, Saumon & Goldblatt (2010)

Clouds in brown dwarfs: Summary Condensates (Fe & silicates) are present throughout the L spectral class (T eff > 1400K) The condensates gravitationally settle into cloud decks Mid- to late-T dwarfs appear to be free of clouds Clearing of clouds at the L  T transition indicated by CMDs Physical mechanism: unknown An increase of sedimentation efficiency? A break up of the cloud layer? There are hints that the L  T transition is gravity sensitive The L  T transition must also occur in hot young planets

Brown dwarfs and hot young planets Clouds and the L  T transition Back to hot young planets

HR 8799b, c, d photometry: Clouds All 3 planets: zJHKL’M photometry can only be fitted with cloudy atmospheres models Marois et al. (2008)

HR 8799b 2.2  m spectrum: Detection of CH 4 Bowler et al. (2010) When fitted with field BD spectra, this spectrum matches that of T1-T3 dwarfs. CH 4 is present in this spectrum

planet b Trouble with HR 8799bcd 9 photometric channels spanning 1.03 to 5  m d=39.4  1.0 pc Luminosities of log L/L  = - 5.1, and (  0.1) Age of the star: Myr Evolution models then give M, R, T eff, gravity Using a variety of brown dwarf atmosphere models, fits to the photometry give T eff for all 3 planets that are too high (or R too small). Bowler et al. (2010) planet b

Trouble with HR 8799bcd Burrows, Sudarsky & Hubeny (2006) Brown dwarf models used by Bowler et al (2010) & Currie et al (2011) The model E sequences do not enter the mag-color space of the L8-L9 dwarfs  Can fit the color but not the magnitude (radius) of the planet Model A cloud extends to the top of the atmosphere (much thicker than E)  more promising (as found by Currie et al 2011) Model E MJMJ J-K MJMJ

A new class of objects? Madhusudhan, Burrows & Currie (2011) `T eff log gM/M J Age (Myr) b c d J-K s MJMJ model E model AE A new, intermediate, “AE” cloud model Good fits for all 3 planets Consistent with evolution and age of the system (30-160Myr) Claim: The hot young planets occupy a different NIR color space than BDs Require models with very thick clouds to explain their colors

Our fits of HR 8799 bcd (work in progress) These best fitting model spectra for c & d are too old and too massive! BUT: within 1.5 , models agree with the upper age limit of 160Myr T eff log gf sed M/M J Age b Myr c Myr d Gyr Data from Marois et al. (2008), Hinz et al. (2010), Currie et al. (2011)

A different view of clouds in hot young planets These planets look like very late L dwarfs only colder Need only to delay the cloud clearing to lower T eff (< 1400K) The T eff of the L  T cloud clearing appears to be rather sensitive to gravity (HD B, 25 M j, has T eff ~1150K, but is cloudy)

Hot young planets have much in common with field brown dwarfs and can be modelled with the same tools and physics The study of clouds in field brown dwarf is highly relevant to hot young planets: Cloud decks of iron and silicate particles Clouds present for T eff > 1400K (all L), absent for T eff <1200K (late T) in field brown dwarfs Clouds disappear over only  200K of cooling in T eff Apparently gravity dependent (e.g. HD B, 2MASS 1207B?) Mechanism presently unknown Exoplanets: What have we learned from brown dwarfs?

HR 8799 planets: NIR colors similar to those of the latest L dwarfs but b & d are fainter Their NIR colors extend the cloudy L dwarf sequence Good fits of the 1 – 4.7  m photometry can be obtained with cloudy models Claims that they have much thicker clouds than field brown dwarfs are not justified The same cloud model (and same f sed range) as for BDs is adequate Understanding these planets only requires that the cloud clearing of the L  T transition occurs at lower T eff for lower gravities Independent evidence of the latter has been accumulating Exoplanets: What have we learned from brown dwarfs?

Backup slides

New spectral class: L dwarfs K I TiO VO FeH H2OH2O Cs I Rb I CrH FeH K I Adapted from Kirkpatrick et al. (1999) TiO and VO bands peak at dM9 then weaken through L sequence Strong CrH and FeH FeH peaks at mid-L Lines of alkali metals appear CO at 2.3  m Strong H 2 O J-K increases steadily T eff ~ 2400 – 1400K L dwarfs are nearly all brown dwarfs

New spectral class: T dwarfs Spectra from Geballe et al. (2002) CH 4 appears in H and K bands CO (2.3  m) disappears by T2 Very strong H 2 O No TiO or VO, weakening FeH and CrH Lines of alkali metals J-K decreases steadily T eff ~ 1400 – 600K T dwarfs are all brown dwarfs 1 2.5

Hot young planets Fomalhaut b Star: A3 V 300  500Myr old dust disk (R= AU) Planet: a ~ 115 AU Mass* ~ 1  3 M Jupiter T eff * ~ 400K Kalas et al., Science, 322, 1345 (2008) * Mass and T eff are early estimates

Other hot young planets? GQ Lup b 2MASS 1207 b AB Pic b  Pic b 1XRS J1609 b star K7 (TTau)M8 (BD)K2 VA5 VK7 V age (Myr) a (AU) >100>54> >330 Mass (M J ) T eff (K) ~1400~1800 notes Too young for models Not bound to 2M1207A Brown dwarf? L’ data only Lafrenière et al (2008), Lagrange et al. (2010) and data compiled in Neuhäuser (2008)

T eff =800 K  cloud = 2/3 T eff =1200 K  cloud = 2/3 Clouds and photospheres All models: log g=5 f sed =3 T eff =1600 K  cloud = 2/3 T eff =2000 K

Fits of entire SED of brown dwarfs Stephens et al.(2009)

Fits of entire SED of brown dwarfs Stephens et al.(2009)

Brown dwarfs and hot young planets Clouds and the L/T transition Non-equilibrium chemistry Back to hot young planets

Unexpected species in the atmosphere of Jupiter AsH 3 Noll, Larson & Geballe (1990) Bjoraker, Larson & Kunde (1986) CO GeH 4 Bézard et al. (2002 )

Chemistry and vertical transport in brown dwarf atmospheres Chemistry of carbon & nitrogen: slow  CO + 3H 2  CH 4 + H 2 O  fast slow  N 2 + 3H 2  2NH 3  fast Transport: Convection (  mix ~ minutes) Radiative zone (  mix ~ hours to years?) e.g. meridional circulation Two characteristic mixing time scales in the atmosphere  CO &  N 2 : ms to > Hubble time! sun

Effect of mixing on abundances Net effect: Excess of CO in the spectrum Depletion of CH 4, NH 3 and H 2 O For CO, CH 4 and H 2 O, the faster the mixing in the radiative zone, the larger the effect (up to saturation) A way to measure the mixing time scale in the atmosphere! The most important opacity sources! }

A case study: Gl 570D (T8) T eff =820K log g=5.23 [M/H]=0 Equilibrium model Model with mixing Model fluxes are computed at Earth Saumon et al. (2006)

Gl 570D: 3-5  m spectrum and photometry T eff = 820K log g=5.23 [M/H]=0 Equilibrium (no CO) With mixing (excess CO) Absolute fluxes CH 4 CO Geballe et al. (2009)

Departure from chemical equilibrium: Summary The consequence of very basic considerations (“universality”) Important at low T eff : late L dwarfs and T dwarfs Affects CO (  ), CH 4 (  ), H 2 O (  ) and NH 3 (  ), all important sources of opacity CO excess observed in all 6 T dwarfs subjected to detailed analysis (e.g. Gl 570D) NH 3 depleted in the IRS spectra of 3 T dwarfs analyzed so far Additional evidence shows that departures from equilibrium chemistry (i.e. vertical mixing) are common in brown dwarfs An opportunity to measure the mixing time scale in the atmosphere Mixing mechanism: turbulent break up of gravity waves?

Chemistry and vertical transport Temperature surface  CO <  mix  CH4 <  mix CO > CH 4 (equilibrium) CH 4 > CO (excess CO)  CO >  mix  CH4 <  mix slow CO  CH 4 fast CH 4  CO

Cloudless model T eff =1000K log g=5 equilibrium Non-equilibrium abundances in brown dwarfs

Cloudless model T eff =1000K log g=5 log K zz (cm 2 /s) =2

Non-equilibrium abundances in brown dwarfs Cloudless model T eff =1000K log g=5 log K zz (cm 2 /s) =4

Non-equilibrium abundances in brown dwarfs Cloudless model T eff =1000K log g=5 log K zz (cm 2 /s) =6

Ground-based photometry: K, L’ & M’ The K L’ M’ color-color diagram provides strong evidence that non-equilibrium chemistry (i.e. enhanced CO) is a common feature of T dwarfs T dwarfs L dwarfs M dwarfs Adapted from Leggett et al. (2007)

Effect on 4.6  m CO band T eff =1000K log g=5 K zz =0, 10 2, 10 4, 10 6 cm 2 /s K zz =0, 10 4 cm 2 /s M’ 4-5  m M’ flux Can affect searches for exoplanets

Effect on CO bands T eff =1000K log g=5 K zz =0, 10 2, 10 4, 10 6 cm 2 /s K zz =0, 10 4 cm 2 /s M’ K band 4-5  m M’ flux Can affect searches for exoplanets

Effect on NH 3 bands T eff =1000K log g=5 K zz =0, 10 2, 10 6 cm 2 /s K zz =0, 10 4 cm 2 /s & no NH 3 log g=5 H bandK band 10  m

Gl 570D: M band spectrum T eff =820K log g=5.23 [M/H]=0 Equilibrium model (no CO) With mixing:  mix =14d With mixing:  mix =6.7h Model fluxes normalized to data 6.5h of integration on Gemini… CO is present in the spectrum (4.4  detection) Geballe et al. (2009)

Other examples of non-equilibrium abundances in brown dwarfs Object Spectral Type T eff (K) log g (cm/s 2 ) [M/H] log K zz (cm 2 /s) Under abund of NH 3 ? Excess of CO? 2MASS 0415 T > 4 Gl 570DT  0.7 2MASS 1217 T > 2 Data too noisy Gl 229BT7p to Marginal data 2MASS 0937 T6p  0.3  Ind Ba T ~ ? No data

HR 8799c spectroscopy: Non-equilibrium? Models in chemical equilibrium (with or without clouds) cannot reproduce the shape of the 4  m spectrum Janson et al. (2010) Cloudy model T eff =1100K, log g=4

HR 8799b, c, d photometry: Excess CO? Models in chemical equilibrium are inconsistent with the M band upper limits Hinz et al. (2010) ( K) ( K)

Burrows, Sudarsky & Hubeny (2006) clouds Model E

Brown dwarfs and hot young planets Clouds and the L  T transition Evolution of brown dwarfs Back to hot young planets

Evolution of brown dwarfs end of the main sequence time BD evolution is primarily: Cooling, contraction and a phase of deuterium fusion Controlled by surface boundary condition Opacity of the atmosphere! (clear vs cloudy)

Brown dwarf evolution: a hybrid model The L-T transition can be reproduced by increasing f sed Toy model of the evolution: Boundary Condition: cloudy (T eff > 1400K) to clear (T eff < 1200K) Interpolate atmospheric B.C. (not the evolution) in between Synthetic disk population IMF: dN/dM ~ M -1 SFR: Constant (0-10Gyr) Monte Carlo sampling Let evolve and… cloudyclear evolution 10 Gyr envelope Reproduces CMD fairly well

Brown dwarf evolution: a hybrid model clear cloudy hybrid Observations: An apparent dearth of brown dwarfs in the transition! A more complete sample should reveal a pile up at the transition Pile up of brown dwarfs in the L-T transition A direct consequence of changes in the atmosphere

Synthetic galactic clusters Synthetic clusters Hybrid evolution IMF: dN/dM ~ M -0.6 Constant SFR over age  1Myr (~ isochrone) A distinctive feature due to deuterium fusion (50-100Myr) In principle, detectable with deep enough surveys

Evolution of brown dwarfs: Summary Brown dwarf evolution is relatively simple Two new features: Cloudy to clear atmosphere transition will cause an accumulation of brown dwarfs in the L-T transition region Deuterium burning phase should be detectable in the CMD of galactic clusters (50-100Myr) Both are potentially observable predictions (e.g. Pan-STARRS)

Synthetic disk population of brown dwarfs Assume IMF: dN/dM ~ M -1 Constant SFR (0-10Gyr) Single BDs only Monte Carlo sampling Let evolve and… evolution 10 Gyr envelope

Brown dwarf evolution: variations on hybrid model