1. Chemistry and line emission of outer protoplanetary disks Inga Kamp Introduction to protoplanetary disks and their modeling Introduction to protoplanetary disks and their modeling Chemistry in the outer disks: - the influence of the central star, Chemistry in the outer disks: - the influence of the central star, PAHs, and X-rays on the disk PAHs, and X-rays on the disk - Deuterium Chemistry - Deuterium Chemistry Pushing the limits of future observing facilities Pushing the limits of future observing facilities Collaborators: Kees Dullemond, Jesus Emilio Enriquez, Bastiaan Jonkheid, Ewine van Dishoeck, Michiel Hogerheijde, et al. Michiel Hogerheijde, et al.

2. Why are the outer disks important?

3. Models of Protoplanetary Disks A quiet protoplanetary disk: - stationary 2D disk models - irradiation by the star (+ accretion) determines the disk structure A more dynamical picture of a protoplanetary accretion disk: - matter is mixed and transportet by turbulence - matter accretes onto the central star dM/dt~10 -7 M Sun /yr - matter continuously falls in from the envelope causing an accretion shock at the disk surface Infalling gas and dust Protostar Protoplanet Chemically active zone Transport of matter and angular momentum V~100 km/s V~10 km/s Visible and UV radiation IR radiation Accretion shock [Aikawa et al. 1999, Gail 2001, Ilgner et al. 2004] [Chiang & Goldreich 1997, Willacy & Langer 2000, Aikawa et al 2002, Jonkheid et al. 2004, Kamp & Dullemond 2004] Posters: Semenov et al. I.63 Willacy et al. III.73

4. Free parameters: Stellar properties, L *, R *, M * Dust properties, opacities, sizes Elemental abundances Disk dimensions, R i, R o Surface density (disk mass) turbulence/diffusion constants

5. The Chemical Network Example: CO formation and destruction C + OH  CO + H k ijk ~ 10 -10 … 10 -9 s -1 cm -3 CO +  C + O  ij ~ 10 -10 … 10 -8 s -1 [ Wedemeyer-Böhm, Kamp, Freytag, Bruls 2004] stationary solution with modified Newton-Raphson algorithm time dependent solution using the Backward Differentiation Formula (BDF) e.g. VODE [ Hindmarsh 1980 ] artificial neural networks [Asensio Ramos et al. 2005] - S(T)  a 2 n g v i n i + n i i e (-E(ads)/kT)

6. What do we know about disk chemistry? The disks are layered: surface layer --> photochemistry intermediate layer --> neutral & ion molecule gas chemistry disk midplane --> gas-grain chemistry [Aikawa & Herbst 1999, Willacy & Langer 2000, van Zadelhoff et al. 2003, Semenov et al. 2004]

7. no PAHs The surface layers can get very hot (UV irradiation) Gas and dust temperatures are not coupled in the surface layers Photoelectric heating on PAHs set the gas temperature in the surface layer [Jonkheid et al. 2004, Kamp & Dullemond 2004, Nomura & Millar 2005] Poster: Geers et al. I.27

8. Chemical destruction of H 2 : H 2 + O --> H + OH C/CO transition at lower/same optical depth as H/H 2 transition Higher UV fluxes lead to lower molecule abundances in the disk atmosphere Very confined OH layer in all T Tauri and Herbig models log n(H 2 )/n(H)log n(CO)/n tot log n(OH)/n tot log n(HCO + )/n tot [Kamp et al. 2004, Nomura & Millar 2005] H/H 2

9. X-rays affect the chemistry and the disk temperature: X-rays enhance the ionization fraction of the disk surface Some molecules have higher abundances due to efficient ion-molecule chemistry (HCN) X-rays can efficiently heat the disk in the absence of strong UV irradiation [Aikawa & Herbst 1999, Kamp et al. 2005] AU Micno X-rays T gas in a 0.01 M  disk around an M star R=700 AU, Z=220 AU no X-rays Poster: Aikawa & Nomura III02

10. Deuterium chemistry: H 3 + is formed by cosmic rays throughout the disk (UV and X-rays do not H 3 + + HD --> H 2 D + + H 2 penetrate that deep) H 2 D + + HD --> HD 2 + + H 2 HD 2 + + HD --> D 3 + + H 2 Posters: Ceccarelli et al. III.13 Ceccarelli & Dominik III.14 D/H in molecules is higher than the elemental D/H ratio in the ISM Destruction via grain surface recombination and reactions with CO,N 2 [Aikawa & Herbst 1999, 2001, Ceccarelli & Dominik 2005]

11. [OI] 6300 Å OH layer above the disk photosphere: OH +  O* + H O* is in the 1 D excited level; it decays to the ground state by emitting a 6300 Å photon + collisional excitation for T gas > 3000 K OI 6300 Å emission in Orions proplyds is restricted to the skin of the disk [Bally et al. 2000, Störzer & Hollenbach 1998; Orion proplyds] [Acke et al. 2005 (Herbig stars)] T Tauri star 0.01 M  dM/dt = 10 -9 M  /yr hot gas log n(OH)/n tot chromosphere

12. Pushing the limits of future observations The mass of small dust grains decreases with stellar age (ISO, Spitzer) [Habing et al. (1999), Meyer et al. (2000), Habing et al. (2001), Spangler et al. 2001] J=2-1 J=1-0 J=3-2 J=4-3 ALMA detection limit CO rotational lines How and when does the gas disappear from the disks? Boundary conditions for planet formation How many failed planetary systems are out there ? Optically thin models (late stages of Herbig Ae star) 1.5 x 10 -4 - 1.5 x 10 -7 M 

13. Conclusions: Need for self-consistent disk models: disk structure + gas chemistry Posters: Semenov et al. I.62, Jonkheid et al. III.35, Nomura et al. III.51 Proper inclusion of ALL radiation sources: stellar UV, X-rays and external Chemistry of the outer protoplanetary disks is driven by irradiation --> Importance of photochemistry Future instrumentation will allow the detection of transition disks down to 0.5 M Earth of gas Outlook: Photochemistry, X-ray chemistry, three-body reactions Gas-grain chemistry: desorption processes, molecule reactions on grains next generation models: 2D hydrodynamical disks with a realistic energy equation (gas temperature), radiative transfer, and full chemistry (gas, gas-grain and grain surface)