Inga Kamp The role of the central star for the structure of protoplanetary disks Peter Woitke, Wing-Fai Thi, Ian Tilling (all Edinburgh)

Protoplanetary Disks and Planet Formation Dust dynamics are controlled by gas, even at late times ! dust substructure dust coagulation dust settling [HR4796: NICMOS 1.6  m Schneider et al. 1999] 8 Myr  Jy/pixel [Reverse ∂p/∂r at 70 AU Klahr & Lin 2001] [Barge & Sommeria 1995, Johansen et al. 2006, 2007]

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

Protoplanetary Disk Models Key astrophysical questions: What is the environment in which planets form ? disk structure - boundary conditions for planet formation Can planets form around stars different from our Sun ? star-disk interaction - impact on disk structure and dispersal

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

Disk properties [2 2 and 10 Myr: Sterzik et al. 2004] [8  m 5 Myr: Carpenter et al. 2006] [Pascucci et al. 2003, Allers et al. 2006, Bouy et al. 2008] Dust in BD disks evolves on same time scale as T Tauri disks Inner disk clearing occurs much faster in T Tauri disks Dust models fit BD, T Tauri and Herbig SEDs in the same way indicating the first steps of planetesimal formation in all cases some BD SEDs require flaring disks, others flat… low # statistics

Disk properties [Pascucci et al. 2009] [Kessler-Silacci et al. 2007] lower HCN/H 2 C 2 ratios in the disks around brown dwarfs weaker and more processed silicate features

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

comparison with observation radiative transfer Protoplanetary Disk Models physical structure chemical composition

comparison with observation radiative transfer Protoplanetary Disk Models physical structure chemical composition surface density  =  0 (r/r 0 ) -p dust temperature T = T 0 (r/r 0 ) -q disk scale height H = H 0 (r/r 0 ) -e [e.g. Menard et al.; Pinte et al. 2006]

comparison with observation radiative transfer Protoplanetary Disk Models physical structure chemical composition surface density  =  0 (r/r 0 ) -p dust temperature 2D continuum RT disk scale height H = (c s 2 r 3 /GM * ) 0.5 [e.g. D’Alessio et al. 1998, Dullemond et al. 2002] Gas chemical modeling on top of fixed density structure [e.g. Aikawa et al. 2002, Semenov et al. 2008]

comparison with observation line radiative transfer chemical composition physical structure Protoplanetary Disk Models atomic/ionized rich molecular chemistry molecule freeze-out hot flaring surface cold midplane [e.g. Aikawa et al. 2002, Kamp & Dullemond 2004] [PPV chapters: Dullemond et al. 2007, Bergin et al. 2007]

comparison with observation Protoplanetary Disk Models radiative transfer physical structure chemical composition [Woitke, Kamp, Thi 2009] ProDiMo Similar approaches have been followed by other groups [Nomura & Millar 2005, Gorti & Hollenbach 2004, 2008]

Chemistry stationary solution with modified Newton-Raphson algorithm (switch to time-dependant) 65 species: H, H +, H -, H 2, H 2 +, H 3 + He, He + C, C +, CO, CO +, CO 2, CO 2 +, HCO, HCO +, H 2 CO CH, CH +, CH 2, CH 2 +, CH 3, CH 3 +, CH 4, CH 4 +, CH 5 + O, O +, O 2, O 2 +, OH, OH + H 2 O, H 2 O +, H 3 O + N, N +, NH, NH +, NH 2, NH 2 +, NH 3, NH 3 +, N 2, HN 2 + CN, CN +, HCN, HCN +, NO, NO + Si, Si +, SiH, SiH +, SiO, SiO +, SiH 2 +, SiOH + S, S+, Mg, Mg+, Fe, Fe+ --> plus CO, H 2 O, CO 2, CH 4, NH 3 ice

Heating and cooling photo-electric heating, PAH heating, viscous (  ) heating, cosmic ray, C photo-ionisation, coll. de- excitation of H ⋆ 2, H 2 on grains, H 2 photodissociation, IR background line heating thermal accomodation on grains, OI, CII, CI fine- structure cooling, CO ro-vibrational (110 levels, 243 lines), o/p H 2 O rotational (45/45 levels, 258/257 lines), o/p H 2 quadrupole (80/80 levels, 803/736 lines), SiII, SII, FeII semi-forbidden (80 levels, 477 lines), Ly , MgII h&k, OI 6300Å Heating Cooling

Heating and cooling photo-electric heating, PAH heating, viscous (  ) heating, cosmic ray, C photo-ionisation, coll. de- excitation of H ⋆ 2, H 2 on grains, H 2 photodissociation, IR background line heating, (X-ray heating) thermal accomodation on grains, OI, CII, CI fine- structure cooling, CO ro-vibrational (110 levels, 243 lines), o/p H 2 O rotational (45/45 levels, 258/257 lines), o/p H 2 quadrupole (80/80 levels, 803/736 lines), SiII, SII, FeII semi-forbidden (80 levels, 477 lines), Ly , MgII h&k, OI 6300Å Heating Cooling

Protoplanetary Disk Models soft inner and outer edges [Hartmann et al. 1989, Hughes et al. 2008, Woitke, Kamp & Thi 2009] stellar UV interstellar UV stellar X rays

The Thermal Balance stellar photons e-e- e-e- e-e- e-e- small grains: a <<  (e.g. PAHs) yield ~ ?? large grains: a >> (micron sized) yield ?? Photoelectric heating of the gas

The Thermal Balance stellar photons e-e- e-e- e-e- e-e- small grains: a <<  (e.g. PAHs) yield ~ ?? large grains: a >> (micron sized) yield ?? Photoelectric heating of the gas [Abbas et al. 2006]

Impact of UV Irradiation FUSE: Apr 2000 STIS: Oct 2001 May 2001 GHRS: Nov 1994 OVI H2 HD H2, CO photodissociationC ionization

Impact of UV Irradiation time UV excess 5000 K

Impact of UV Irradiation time UV excess 0 log n(OH)/n(tot)

Protoplanetary Disk Models soft inner and outer edges [Hartmann et al. 1989, Hughes et al. 2008, Woitke, Kamp & Thi 2009] stellar UV interstellar UV stellar X rays

The Thermal Balance stellar X-rays e-e- e-e- X-ray heating of the gas [Glassgold et al. 2004, Meijerink et al. 2008] TW Hya L X ~ erg/s [Kastner et al. 1997] [Nomura et al. 2007] e-e- * Auger electrons Secondary ionization Primary ionization Å1.24 Å *

Protoplanetary Disk Models Flaring disk structure “early Sun” M * = 1 M  L * = 1 L  T eff = 5800 K M disk = M  AU dust: 90% Mg 2 SiO 4 10% Fe  m (2.5) [Woitke, Kamp, Thi 2009] X-ray 20 R  to irradiate disk midplane

Protoplanetary Disk Models Flaring disk structure “Herbig star” M * = 2.2 M  L * = 32 L  T eff = 8500 K M disk = M  AU dust: 90% Mg 2 SiO 4 10% Fe  m (2.5) T TauriHerbig Ae

Protoplanetary Disk Models Flaring disk structure “Herbig star” M * = 2.2 M  dust: L * = 32 L  90% Mg 2 SiO 4 T eff = 8500 K 10% Fe M disk = 10 -2, 10 -3, M   m (3.5) AU self-similarity decreasing disk mass Herbig AeT Tauri

Outline Various planet forming environments (BD - T Tauri - Herbig) Second generation disk structure modeling How do observations inform us about the star-disk interaction ?

Probes of UV Irradiation: [OI], OH [Bally et al. 2000, Störzer & Hollenbach 1998, Acke et al. 2005] WFPC2 OH photodissociation layer: OH +  O* + H O* denotes the 1 D 2 excited level (decay emits a 6300 Å photon) 460 AU [Mandell et al. 2008, Salyk et al. 2008]

Probes of UV Irradiation: [OI], OH [Bally et al. 2000, Störzer & Hollenbach 1998, Acke et al. 2005] OH photodissociation layer: OH +  O* + H O* denotes the 1 D 2 excited level (decay emits a 6300 Å photon) [Mandell et al. 2008, Salyk et al. 2008] normalized flux velocity [km/s] VLT/UVES: HD97048

Probes of UV, X-ray Irradiation: CO, H 2 [Brittain et al. 2007] Fluorescence excitation vs. thermal emission: CO UV pumping of electronic states => cascade into high v levels H 2 UV (Lyman & Werner bands) or X-rays (collisions with fast e - ) [v=1-0 S(1), S(0), v=2-1 S(1) NIR: Carmona et al. 2008] see talk by Gerrit van der Plas

Probes of UV, X-ray Irradiation: CO, H 2 [Brittain et al. 2007] Fluorescence excitation vs. thermal emission: CO UV pumping of electronic states => cascade into high v levels H 2 UV (Lyman & Werner bands) or X-rays (collisions with fast e - ) [S(1), S(2), S(4) pure rotational MIR lines: Bitner et al. 2008] see talk by Gerrit van der Plas

Probes of X-ray Irradiation: NeII [Glassgold et a. 2007, Pascucci et al. 2007, Lahuis et al. 2007, Meijerink et al. 2008] see talk by Richard Alexander

R (AU) D: 13 CO 3-2 E: 13 CO 1-0 C: 12 CO 2-1 H: HCO H [van Zadelhoff et al. 2001] Probes of gas-dust decoupling [Pietu et al. 2007] temperatures derived from gas lines at the surface are higher than dust temperatures => support for decoupling of gas and dust in the surface layers C E D

[Fedele et al. 2008] Probes of gas-dust decoupling dust is flat (self-shadowed) and gas is flaring => support for decoupling of gas and dust in the surface layers MIR dust emission [OI] emission

GASPS: A Herschel Key program GAS evolution in Protoplanetary Systems (PI: Dent) ~200 disks distributed over spectral type, age, and disk mass [CII], [OI], CO and H 2 O lines Aim: probe the gas mass evolution throughout the planet forming phase

line radiative transfer physical structure chemical composition comparison with observation GASPS: A Herschel Key program Monte Carlo line radiative transfer understand where the line originates and then put resolution there !! Kamp, Tilling, Woitke, Thi, Hogerheijde [Hogerheijde, van der Tak 2000]

Key points to take home Models predict that stellar UV and X-ray radiation shape the disk structure and chemistry and we actually observe that ! Individual gas lines DO NOT probe total disk mass ! but we keep trying with multiple tracers … Gas lines DO probe the physical and chemical conditions in the region where they arise (interplay between radiative transfer and chemical structure) and we try hard with GASPS to find the desperately wanted M  disks…

Thank You !