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Structure and Evolution of Protoplanetary Disks Carsten Dominik University of Amsterdam Radboud University Nijmegen.

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Presentation on theme: "Structure and Evolution of Protoplanetary Disks Carsten Dominik University of Amsterdam Radboud University Nijmegen."— Presentation transcript:

1 Structure and Evolution of Protoplanetary Disks Carsten Dominik University of Amsterdam Radboud University Nijmegen

2 What would we like to know? Formation and Evolution Spectral Energy Distributions –and what they do and dont tell us Grains –Sizes –Composition –Distribution as function of r,z,t Gas –Mass –Composition –Distribution as function of r,z,t Dynamics –Rotation and inflow –Disk winds –Viscosity, turbulence, accretion, instabilities

3 Figure from Greene 2001)

4 Disks in a nutshell Infalling matter has non-zero angular momentum, lands on rotation plane away from star Star mass dominates, matter on largely Keplerian orbits Some kind of viscosity couples different annuli of the disk, matter spreads, most falls onto star, some mass moves outward and carries angular momentum As infall stops, disk mass decreases, eventually disappears into star, planets, or space!

5 Formation & viscous spreading of disk Fig. from C. Dullemond

6 Formation & viscous spreading of disk Hueso & Guillot (2005)

7 Evolution of disks with time Disks live a few million years Near-IR disk fraction J. AlvesL. Hillenbrand

8 How large are disks? Hundreds, up to a thousand AUs –In scattered light –Dust millimeter emission –Images in CO mm lines Different techniques will give different sizes –The mm continuum probes the dust in the disk midplane –Scattered light and CO probe a layer higher –CO lines are brighter than the dust continuum, so disks are larger in the CO lines than in the continuum –Disk size will depend on sensitivity unless sharp outer edge Surface density -Little direct information in the inner disk -Measured in the outer disk (R>30AU) with continuum maps r -1

9 Pre-MS disks are big DM Tau0.5 Msun850 AU GM Aur0.8 Msun500 AU LkCa151 Msun500-600 AU MWC 4802 Msun450 AU HD1632962.4 Msun550 AU AB Aur2.3 Msun1000 AU HD 342822 Msun800 AU

10 Mdisk ~ 0.001- 0.1 Msun if k(1mm)~1 cm 2 /g Disk masses Dust mass from submm flux, assume k(1mm), gas-to-dust ratio = 100

11 Dynamics in viscous disk Keplerian rotation: v φ =(GM * /R * ) 1/2 Radial drift toward the star: v R ~ c s H/R No vertical motions: v z =0 Turbulence –v t < c s << v φ

12 The deadzone in accretion disks x M =10 -15 10 -12 10 -9

13 HD163296 : 12CO J=2-1 Rotation from CO mm lines: a velocity gradient across the major axis [Isella et al. 2006] N E Mstar = 2.0 0.5 Msun incl = 45°

14 Deviations from Keplerian: Hogerheijde 2001 –infall in TMC1 Pietu et al 2005 –V R 0.41 +/- 0.01 in AB Aur Turbulence in the outer disk is very hard to measure

15 Gas and dust are initially well mixed Dust dominates the opacity at almost any wavelength Disk is thick because of hydrostatic equilibrium (pressure against gravity). –Density decreases exponentially with height –When small grains exist and are well mixed, stellar radiation is absorbed at about 4 pressure scale heights. Disk shape and composition

16 1.Viscous dissipation (~(M 1/2 /r 3 * dM/dt) 2.Stellar radiation (~L * /r 2 ) Disks contain warm dust around a star - what it heating the dust? HAeBe T Tauri

17 Disk emission Star Disk

18 Dust, Gas, Radiation small(?),large grains small grains, PAH PDR: atoms, ions, small molecules CI, NeII... Hot gas CO, H 2 O Molecules: CO,HCO +... Ice mantles, H 3 + PAH UV CR,X dead?

19 Gas temperature gets very high in upper layers Woitke, Kamp, Thi 2009

20 General structure of the disk Fig. from Dullemond et al, PPV

21 Submm allows us to look at the whole disk

22 V-band 24um 33um Mulders et al in prep

23 Spectral Energy Distribution(SEDs)

24 The snowline, depending on accretion Min et al 2010



27 Sources of Water in the disk: + photo desorption photo dissociation gas phase formation route freeze-out/reformation

28 shallowsample, ~2000 sec DM Tau Integrated for 198 min at 557 GHz and 328 m at 1113 GHz No significant detection of either ortho or para H 2 O weak 6σ detection of 557 GHz line (1 10 - 1 01 ) Models indicate ice depletion (due to settling?) Bergin et al 2010

29 Grain sizes and spatial distribution

30 Main grain size processes Settling Radial Drift Turbulent mixing and concentration Gravitational instabilities?

31 Effects of dust settling Dullemond & Dominik (2004)

32 SED differences in FIR As before, but replacing mass by large grains at the equator instead of removing it

33 Evidence for grain growth v Boekel et al 2003 Small grain Large grain

34 MostT Tauri disks shows evidence for grain growth Kessler-Silacci et al. 2006,2007 10 m band20 m band Obs Model

35 Radial drift of particles Weidenschilling 1977, Brauer et al 2008 1AU in 100 ys

36 Radial drift times Brauer et al 2007 Collective effects help only if the disk is very weakly turbulent Reducing the gas mass does help

37 Radial motion changes disk sizes Takeuchi, Clarke, Lin 2005 Mm-sized grains move to below 100 AU in 10 5 years Porosity increases life time

38 Sources of relative velocities Brownian motion Settling Radial drift Coupling and decoupling to turbulent eddies –Complex expression depending on details of turbulence and dust properties (e.g. Ormel & Cuzzi 2007)

39 Relative velocities: Total Brauer et al 2008

40 Coagulation only, different velocity sources Brauer et al 2008

41 Effects of radial motion Brauer et al 2008 1 cm 1 m

42 With fragmentation at 10m/s Brauer et al 2008 1 cm 1 m

43 With higher fragmentation speed 30m/s Brauer et al 2008 1 cm 1 m

44 Testi et al 2001 Optically thin disk: Observed: Large grains in outer disk

45 Birnstil et al 2010

46 The inner disk

47 Isella and Natta 2005

48 The location of the inner rim Equilibrium temperature of a dust grain in free space Including backwarming

49 Optically thin dust inside the rim

50 Moving the rim with refractory shields Kama et al 2009

51 A selection of inner rim structures

52 Inner holes: transition objects Calvet et al. 2005 Rin=0.03,1,10,30 AU

53 Disk evaporation Photoevaporation by EUV, FUV and X-ray photons, @ 30 AU Life times enough for planet formation Disk survives for 10 6 years after gap formation Short-lived disks for M * >3M o Gorti et al 2009

54 LkCa 15 and disk geometry See also talk by Nuria Calvet Espalliat et al 2007-2010 Mulders et al 2010

55 Optically thin matter in the inner disk? Benisty et al 2010

56 Dominik & Dullemond, in preparation Chiang & Murray-Clay 2007 Clearing out the disk when a gap is already present

57 Summary Disks are everywhere, with a wide variety of properties Planet formation by just coagulation seems to be too hard The presence of mm-cm grains in the outer disk is not fully understood Inner gaps and optically thin material in these gaps are a hot topic right now

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