Primordial Disks: From Protostar to Protoplanet Jon E. Bjorkman Ritter Observatory

Cloud Cores

Bok Globule: Isolated Cloud Core

Theorist’s Cloud Core

Star Formation Within Cloud Cores –gravity overcomes gas pressure gas must be cold –cores collapse Free-fall Inside out (Shu 1977) Form protostars –rotation Cloud flattens into disk material falls on disk –protostar accretes material from disk

Rotating Infall Streamlines follow ballistic trajectories –Ulrich (1976); Cassen & Moosman; Terebey, Shu, & Cassen (1984) Keto

Accretion with Rotation Accretion termination shock above/below disk surface Material added at centrifugal radius (orbital periastron) –Centrifugal radius grows with time

Young Stellar Objects

Circumstellar Disks

Disk Winds Matt 2005 Magneto-Centrifugal –Blandford & Payne (1982) –Pudritz & Norman (1983) Magnetospheric –X Wind (Shu et al. 1994)

T Tauri SED Adams, Lada, & Shu 1987 IR Excess –Starlight reprocessed by disk (passively irradiated disk) –L disk ~ 1/4 L star –Shape determined by temperature vs radius UV excess –Disk-Star boundary layer / accretion shock –Causes “veiling” of spectral lines

SED Classification Class 0-III –Adams, Lada & Shu 1987 Class 0: –Mostly sub-mm emission –Deeply embedded protostars Class I: –Rising SEDs from 2 to 100  m –Protostars still accreting from infalling envelope Class II (Classical T Tauri): –Falling IR SEDs –Stars surrounded by disks Class III (Weak-lined T Tauri): –Little IR excess –Almost no circumstellar material

Star/Disk Formation Sequence Class 0Class IClass II Class IIIDebris Disks

Keplerian (Orbiting) Disks Fluid Equations Vertical scale height (Keplerian orbit) (Scale height) (Hydrostatic)

Disk Temperature Kenyon & Hartman 87 Flared Reprocessing Disk Adams, Lada, & Shu 88 Flat Reprocessing Disk

Flaring Effects: Disk Temperature & SED Kenyon & Hartmann 87 log wavelength (micron) Near IR Far IR

Viscous Accretion Disk Sources of Viscosity –Eddy Viscosity (Shakura & Sunyaev 1977) –Magneto-Rotational Instability (Balbus & Hawley 1991) requires slight ionization Possible dead zones in disk interior Lee, Saio, Osaki 1991

Viscosity in Keplerian Disks Viscosity Diffusion Timescale (eddy viscosity) Lynden-Bell & Pringle 1974

Steady State Accretion  -Disks (surface density) (scale height) (Keplerian orbit) (hydrostatic) (continuity eq.)

Power Law Approximations Keplerian Accretion Disk Flaring

3-D Monte Carlo Radiation Transfer Divide stellar luminosity into equal energy packets Pick random starting location and direction Transport packet to random interaction location Randomly scatter or absorb photon packet When photon escapes, place in observation bin (frequency and direction) REPEAT times

T Tauri Model SED

MC Radiative Equilibrium Sum energy absorbed by each cell Radiative equilibrium gives temperature When photon is absorbed, reemit at new frequency, depending on T

T Tauri Envelope Absorption

Monte Carlo Disk Temperature Whitney, Indebetouw, Bjorkman, & Wood 2004

Radial Temperature Structure Snow Line: Water Ice Methane Ice Optically thin T ~ r -0.4 Midplane Surface

Vertical Temperature Structure Dullemond

3-D Temperature Effects At large radii –outer disk is shielded by inner disk –temperatures lowered at disk mid-plane Surface layers –Heat up to optically thin dust temperature (Chiang & Goldreich 97) –Upper layers “puff up” Inner edge of disk –Heats up to optically thin dust temperature –Inner edge “puffed up” (relative to flat disk) –Shadows disk behind inner wall

Effect of Inner Wall Dullemond, Dominik, & Nata 01

Disk Self-Shadowing Dullemond, Dominik, & Nata 01 Dullemond 02

Protostar Evolutionary Sequence i =80i =30 Mid IR Image DensitySpectrum Whitney, Wood, Bjorkman, & Cohen 2003

Protostar Evolutionary Sequence Mid IR Image DensitySpectrum i =80i =30 Whitney, Wood, Bjorkman, & Cohen 2003

Disk Evolution: Decreasing Mass Wood, Lada, Bjorkman, Whitney & Wolff 2001

Forming Planets: Standard Model Dust grains stick together –form rocks Grow into planetesimals –some still survive today Asteroids & comets Larger planetesimals attract smaller ones (gravity) Planetesimals accrete –form planet cores

Dust Processing in Disks Gravity causes dust settling toward mid-plane –~10 4 yr Grain Growth –Grain size increases with disk age? Ice Condensation –dust may be coated with ice Dust Removal –Radiation Pressure Poynting Robertson Effect –Gas Drag Accretion onto star (or planets) Blown away by stellar / disk wind –Evaporation (when dust gets too hot)

Dust Opacity Mie Scattering Opacity Dust has a particle size distribution

Dust Opacity Wood, Wolff, Bjorkman, & Whitney 2001

Evidence for Grain Growth Wood, Wolff, Bjorkman, & Whitney 2001 Bjorkman, Wood, & Whitney ISM Dust GrainsLarge Dust Grains (1mm)

Evidence for Grain Growth Wood et al Cotera et al Small Grain ModelLarge Grain Models HH30 Observations

Evidence for Dust Settling Observed scale height < thermal value Self-Shadowed Disks? –Dust settling reduces opacity in disk surface layers –Reduced absorption in surface layers reduces disk heating –Causes outer disk collapse, producing fully self- shadowed disk

Holes in Protoplanetary Disks

Transition Disks: GM AUR SED Inner Disk Hole Size = Jupiter’s Orbit Rice et al. 2003

Planet Hole-Clearing Model Rice et al. 2003

Planetary Gaps Kley 1999

Gap Structure Bjorkman et al. 05

Predicted Gap Images Bjorkman et al. 05

Predicted Gap SED Gap + Inner Hole Gap Only Varniere et al. 2004