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Circumstellar disks - a primer
Ast622
The Interstellar Medium
Partially based on Les Houches lecture by Michiel Hogeheijde
(http://www-laog.obs.ujf-grenoble.fr/heberges/Houches08/index.htm)
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Motivation
• The last step in the transport of the ISM to stellar scales
• The first step in the formation of planetary systems
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Disks are an inevitable ( ubiquitous?) consequence of angular momentum conservation
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Indirect evidence for disks
• Emission line (H) stars above the main sequence accretion
• Infrared-millimeter excess emission reprocessing of starlight by a non-spherical geometry
• Ultraviolet excess and X-ray emission accretion hot spots and star-disk interface
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Direct evidence for disks(i.e. imaging)
Smith & Terrile 1984
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Direct evidence for disks(i.e. imaging)
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SED classification
αIR = -dlog(νFν)/dlog(ν)
= log(25F25/2F2)/log(2/25)
(Lada 1987)
Fig from Andre
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SED theory
• Chiang & Goldreich (1997) following pioneering work by Adams, Lada & Shu (1987), Kenyon & Hartmann (1987)
• Also see reviews by Beckwith (1999) and Dullemond et al. (2006)
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Flat blackbody disk
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Flat blackbody disk
Observe
d
Fig. 1.— SED for the flat blackbody disk, with contributions from star and disk identified. The n = 4/3 law is evident between 30 μm and 1 mm. The turnover near 1 mm is due to our truncation of the disk at ao ≈ 270 AU. Chiang & Goldreich 1997
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Flared blackbody diskThe vertical component of gravity will decrease with radius along with the surface density. Hydrostatic equilibrium then implies the disk scale height increases with radius: the disk is flared.
The outer regions of the disk of a flared disk intercept more starlight than a flat disk and the mid-to-far infrared emission is stronger.
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Flared blackbody disk
Fig. 2.— SED for the flared blackbody disk. At mid-IR wavelengths, Lν ∝ ν−2/3. At longer wavelengths, Lν ∝ ν3.
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Radiative equilibrium disk
Fig. 3.— Radiative transfer in the passive disk. Stellar radiation strikes the surface at an angle α and is absorbed within visible optical depth unity. Dust particles in this first absorption layer are superheated to a temperature Tds. About half of the emission from the superheated layer emerges as dilute blackbody radiation. The remaining half heats the interior to a temperature Ti.
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Radiative equilibrium disk
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Radiative equilibrium disk
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Radiative equilibrium flared disk
Fig. 6.— SED for the hydrostatic, radiative equilibrium disk. At mid-IR wavelengths, the superheated surface radiates approximately 2–3 times more power than the interior. Longward of 300 μm, n gradually steepens from about 3 to 3 + β as the disk becomes increasingly optically thin.
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Radiative equilibrium flared disk
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Adding in solid state features
Fig. 10.— SED for the hydrostatic, radiative equilibrium disk using a grain emissivity profile motivated by data from Mathis (1990). For wavelengths shorter than 0.3 μm, our assumed emissivity is unity; longward of 0.3 μm, it obeys a (single) power-law relation ∊λ = (0.3 μm/λ)1.4, on which are superposed two Gaussians centered on 10 and 20 μm, having amplitudes that are 3 times their local continuum emissivity and FWHM equal to 3 and 9 μm, respectively.
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Flaring + hot inner rim
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Dullemond et al. 2006, PPV review
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Dependence of SED on disk geometry
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Dependence of SED on disk geometry
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Dependence of SED on disk geometry
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Dependence of SED on disk geometry
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Dependence of SED on disk geometry
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SED + spatial modeling
disk mass, radius and temperature and surface density profiles, T ~ R-q, ~ R-p
Andrews & Williams 2007
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Accretion disk theory
Same temperature profile (and hence SED) as as passive flat blackbody disk, T R-3/4
Flared disk SEDs dominated by stellar irradiation.
Accretion critical for understanding disk evolution
L = GMMdot/R
Annual Reviews 1981
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Viscous evolution
accretion shocks
magnetospheric accretion
spreading
Muzerolle et al. (1998, 2001)
Gullbring et al. (1997)
As disk accretes to star, conservation of momentum implies disk spreads out; mass, accretion, decrease with time, radius increases with time.
Andrews & Williams 2007
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Dust mineralogy
van Boekel et al. (2004)
observed
olivine
pyroxene
hydrosilicate
ISM silicate
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Grain growth
submillimeter emission “efficiency” ~ 2 ~ 0
ISM grains pebbles/snowballs
related to size of largest solids in diske.g. Pollack et al. (1994), Draine (2006)
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Grain growth
Isella et al. 2007
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disk ~ 1ISM ~ 2
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Grain growth
Andrews PhD thesis 2007
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Dust settling
Dullemond et al. 2004
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The gaseous disk
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Molecular Hydrogen
H2 is difficult to detect• no permanent dipole -> no dipole rotational transitions; only weak quadrupole transition in mid-IR that require hundred K or more to excite• conflicting reports about detection• fluorescent H2 emission in UV (electronic transitions) and near-infrared (vibrational) has been detected but is difficult to analyze quantitatively
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Molecular Hydrogen
Lahuis et al. (2007)
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Near infrared disk ro-vibrational lines
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Boogert et al. 2002
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Recent Spitzer IRS results
Watson et al. 2007 Carr & Najita 2008
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Atomic fine structure lines in disks:probes of the giant planet forming region
Herschel GASPS Key Program
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Atomic fine structure lines in disks:probes of the giant planet forming region
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Herschel GASPS Key Program
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Millimeter observations:the cold outer reservoir
Simon et al. (2000)
• <1% by mass of gas consists of CO, and smaller quantities of other molecules and atoms
• CO easily detected in mm rotational transitions• shows rotation patterns• inferred masses 10-100 times smaller than from dust:
depletion• CO freezes out on dust grains for T<20 K
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Qi PhD thesis 2000
Millimeter observations:the cold outer reservoir
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Disk chemistry
• most molecules now understood to be present only in a warm layer at intermediate height and close to the star• frozen out in mid-plane• photo-dissociated in the disk surface
Semenov et al. (2008)
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Disk chemistry:resolving the D/H ratio
Qi et al. (2008)
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Disk lifetimes
Haisch et al. 2001 Hillenbrand 2005fdisk > 80% at ~1 Myrfdisk ~ 50% at ~3 Myr
fdisk ~ few% at >10 Myr
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Disk lifetimes
Andrews & Williams 2005
NIR excess outer diskInner and outer disks have
similar dissipation timescales
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sub-mm emission(disk masses)decreases with IR SED evolution
sub-mm SEDchanges with
IR SED evolution(particle growth)
Class I disks
Class II disks
Class III disks
Disk evolution (at mm)
Sean Andrews PhD 2007
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Transitional disksViscous evolution is expected to be quicker at small radii but transitional disks, with mid-infrared dips in their SED and cold outer rings of dust and gas are rare (and possibly only seen around binaries?)
Brown et al. 2008
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Disk clearing through photoevaporation
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Alexander et al. 2006
Alexander “UV-switch” model where stellar wind very rapidly erodes disk (from inside out but in only ~105yr) as accretion rate drops below photoevaporation rate
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External photoevaporation
O’Dell, McCaughrean, Bally Williams et al. 2005
Rapid mass loss, 10-5 M☉/yr, at center, but massive disks survive at large distances (Rita Mann PhD)
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Debris disks
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astro-ph/0511083
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Debris disks
Williams et al. 2004
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Isella et al. 2007
Debris disks have double peaked SEDs with a stellar photosphere plus generally a single temperature dust component. They have very low (if any) gas and have a much simpler geometry than protostellar disks.
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Debris disks
See http://astro.berkeley.edu/~kalas/disksite/
As for protostellar disks, images are rare (but critically important); many properties inferred from infrared excesses and SED studies alone.
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Summary• Disks are ubiquitous
– but generally only indirectly inferred from infrared excesses
• Masses range from 0.001–0.3 M☉
• Radii range from tens to hundreds of AU• Grains in disks grow to cm sizes• Gas shows Keplerian motion
– Many molecules (but not H2) frozen out in cold interior
• The fraction of stars with disks decreases with time– from >80% at <1 Myr to <10% at 10 Myr– ‘half-life’ of disks ~3 Myr– inner and outer disk dissappear almost simultaneously
• Debris disks from planetesimal collisions may be visible for >>100 Myr after star formation