structure and evolution of protoplanetary disks
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Structure and Evolution of Protoplanetary Disks. Carsten Dominik University 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 - PowerPoint PPT PresentationTRANSCRIPT
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
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Fig. from Dullemond et al, PPV
Submm allows us to look at the
whole disk
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V-band
24um
33um
Mulders et al in prep
The snowline, depending on accretion
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Min et al 2010
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Sources of Water in the disk:
+
photo desorption photo dissociation
gas phase formation route freeze-out/reformation
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shallowsample, ~2000 sec
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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?)
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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
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Weidenschilling 1977, Brauer et al 2008
1AU in 100 ys
Radial motion changes disk sizes
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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
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Brauer et al 2008
Coagulation only, different velocity sources
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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
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The inner disk
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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
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• Including backwarming
Optically thin dust inside the rim
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Moving the rim with refractory shields
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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
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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?
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Benisty et al 2010
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Dominik & Dullemond, in preparation
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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