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

Structure and Evolution ofProtoplanetary Disks

Carsten DominikUniversity of Amsterdam

Radboud University Nijmegen

What would we like to know?• Formation and Evolution• Spectral Energy Distributions

– and what they do and don’t 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

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Figure from Greene 2001)

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!

Formation & viscous spreading of disk

Fig. from C. Dullemond

Formation & viscous spreading of disk

Hueso & Guillot (2005)

Evolution of disks with time

• Disks live a few million years

Near-IR disk fraction

J. Alves L. Hillenbrand

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

Pre-MS disks are big

DM Tau 0.5 Msun 850 AU

GM Aur 0.8 Msun 500 AU

LkCa15 1 Msun 500-600 AU

MWC 480 2 Msun 450 AU

HD163296 2.4 Msun 550 AU

AB Aur 2.3 Msun 1000 AU

HD 34282 2 Msun 800 AU

Mdisk ~ 0.001- 0.1 Msun if k(1mm)~1 cm2/g

Disk masses• Dust mass from submm flux, assume k(1mm), gas-to-dust ratio = 100

Dynamics in viscous disk

• Keplerian rotation: vφ=(GM*/R*)1/2

• Radial drift toward the star: vR~cs H/R

• No vertical motions: vz=0

• Turbulence– vt < cs << vφ

HD163296 : 12CO J=2-1HD163296 : 12CO J=2-1

Rotation from CO mm lines: a velocity gradient across the major

axis [Isella et al. 2006]

N

E

Mstar = 2.00.5 Msun

incl = 45°

Deviations from Keplerian:• Hogerheijde 2001

– infall in TMC1

• Pietu et al 2005–V R 0.41 +/- 0.01

in AB AurTurbulence in the outer disk is very hard to measure

• 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

1. Viscous dissipation (~(M1/2/r3

* dM/dt)

2. Stellar radiation (~L*/r2)

Disks contain warm dust around a star - what it heating the dust?

HAeBe

T Tauri

Disk emissionStar

Disk

Dust, Gas, Radiation

small(?),large grainssma

ll gra

ins, P

AH

PDR: atoms, ions, small molecules

CI, NeII...

Hot gas

CO, H2O Molecules: CO,HCO +...

Ice mantles, H3+

PAH

UV

CR,X

dead?

Gas temperature gets very high in upper layers

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Woitke, Kamp, Thi 2009

General structure of the disk



Fig. from Dullemond et al, PPV

Submm allows us to look at the

whole disk



V-band

24um

33um

Mulders et al in prep

The snowline, depending on accretion



Min et al 2010





Sources of Water in the disk:

+

photo desorption photo dissociation

gas phase formation route freeze-out/reformation



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 H2O

• weak 6σ detection of 557 GHz line (110 - 101)

• Models indicate ice depletion (due to settling?)



Bergin et al 2010

Grain sizes and spatial distribution

Main grain size processes Settling

Radial Drift Turbulent mixing and

concentration Gravitational instabilities?

Effects of dust settling

Dullemond & Dominik (2004)

SED differences in FIRAs before, but replacing mass by large grainsat the equator instead of removing it

Evidence for grain growth

v Boekel et al 2003v Boekel et al 2003

Small grain

Large grain

MostT Tauri disks shows evidence for grain growthKessler-Silacci et al. 2006,2007

10 m band 20 m band

Obs

Model

Radial drift of particles



Weidenschilling 1977, Brauer et al 2008

1AU in 100 ys

Radial motion changes disk sizes



Takeuchi, Clarke, Lin 2005

• Mm-sized grains move to below 100 AU in 105 years

• Porosity increases life time

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)

€

ΔvB =8kT

π×m1 + m2

m1m2

€

Δvsett =3ΩK

2 z

4ρcs

m1

σ 1

−m2

σ 2

€

Δvdrift = vd,1 − vd,2

Relative velocities: Total



Brauer et al 2008

Coagulation only, different velocity sources



Brauer et al 2008

Effects of radial motion

Brauer et al 2008

1 cm

1 m

With fragmentation

at 10m/s

Brauer et al 2008

1 cm

1 m

With higher fragmentation speed 30m/s

Brauer et al 2008

1 cm

1 m

€

a = β + 2

β =α −1

Testi et al 2001

€

Fν ∝να

€

κν ∝ν β

Optically thin disk:

Observed: Large grains in outer disk

Birnstil et al 2010



The inner disk



Isella and Natta 2005

€

Td =Td (ρ )

The location of the inner rim

• Equilibrium temperature of a dust grain in free space

€

rrim ∝ cBW€

L*

4πr2cabs (T*) = 4cabs (Td )σTd

4

€

L*

4πr2cabs (T*) =

4

cBWcabs (Td )σTd

4



• Including backwarming

Optically thin dust inside the rim



Moving the rim with refractory shields



Kama et al 2009

A selection of inner rim structures

Inner holes: transition objects

Calvet et al. 2005

Rin=0.03,1,10,30 AU

Disk evaporation• Photoevaporation by EUV, FUV and X-ray photons, @ <10 and >30 AU

• Life times enough for planet formation

• Disk survives for 106 years after gap formation

• Short-lived disks for M*>3Mo



Gorti et al 2009

LkCa 15 and disk

geometry

See also talk by Nuria CalvetEspalliat et al 2007-2010Mulders et al 2010

Optically thin matter in the inner disk?



Benisty et al 2010



Dominik & Dullemond, in preparation



Chiang & Murray-Clay 2007

Clearing out the disk when a gap

is already present

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

Documents

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