Observations of Disks around Young Stellar Objects

Observations of Disks around Young Stellar Objects

G. Duchêne & F. Ménard (Obs. Grenoble)

G. Duchêne - Structure Formation in the Universe - May 2007

Goals of this talk

Consider as wide a range of datasets as possible in 30 minutes! Will skip some very exciting aspects

Discussion of selected physical aspects Leave out gas and chemistry

Persuade you that we can now constrain some physical processes Yet many open questions remain…


Outline

General motivation Observational methods Disks in the context of star formation Disks in the context of planet formation Debris disks: after planets formed Summary and perspectives


General MotivationGeneral Motivation


Why do we care about disks?

A natural outcome ofstar formation


Why do we care about disks?

A natural outcome ofstar formation

Planetary systemfactories


Expected physical processes (I)

Influence of central star/environment Disk lifetime Total mass reservoir Overall structure

Disk dispersal mechanism Viscous dissipation of angular momentum Photo-ionization Dynamical dispersal (companion)


Expected physical processes (II)

Substructure formation Spiral arms (instabilities, planets) Gap openings (planets)

Dust evolution Grain growth Radial migration Vertical sedimentation Change in grain structure


Observations of disksObservations of disks


Unresolved datasets: SEDs

The simplest approach: gather the energy and try to invert to disk structure Flared disks in most cases

Chiang & Goldreich (1997) Dullemond et al. (2007)

FlatFlared


Unresolved datasets: SEDs

Useful approach for statistical purposes Can be dangerous on an object-to-

object basis Need for resolved datasets!

Burrows et al. (1996)

D’Alessio et al. (2001)

All TaurusCTTS


Resolved datasets

A single image provides key parameters: Outer radius, position angle Inclination (sometimes) Optical depth (sometimes)

Bertout et al. (1998)

Guilloteau et al. (1999)

Not a normaldisk


Resolved datasets

VLT/VISIR

SpitzerInterf.

All probe different dust populationsmass

grain size

structure

composition


Resolved datasets

Need for complementary complex RT models

VLT/VISIR

SpitzerInterf.


Disks and Star FormationDisks and Star Formation


Disks and central object mass

How universal is star formation? Probe disk presence through IR excess

Overall fraction up to 90% (in Oph) Best studied population: the ONC

disks at all masses (0.1 - 5 M)

Hillenbrand et al. (1998)

Slight deficit atlow mass end?



Detection is harder around VLMS/BD because of cooler Teff

BDs: 40-75% up to ~5 Myr at least No substantial difference with stars

Liu et al. (2003)Jayawardhana et al. (2003)



Not only is disk frequency independent of mass, their structure is, too! Hydrostatic (flared) passive disks

Glauser et al. (2007) McCabe et al. (2007) Perrin et al. (2007)

~0.1 M ~0.5 M ~2 M

PDS 144HK TauIRAS 04158+2805



The special case of high-mass stars: Aligned (rotating) methanol masers, but not

so clear Norris et al. (1993), De Buizer et al. (2003)

Wide-angle outflows A huge ‘silhouette disk’ Difficult to conclude yet

Too far away Evolving too fast

Chini et al. (2004)

M17

20000 AU


Disks and orientation of stars

Taurus molecular cloud = series of filaments orthogonal to B field

So are individual pre-stellar cores

Hartmann (2002)CO map Prestellar cores


Disks and orientation of stars

Disks around T Tauri stars indicate the system’s symmetry axis Systems are randomly oriented w.r.t. local

magnetic field What happened?

Non-magnetic collapse?

Ménard & Duchêne (2004)


Disks and Planet FormationDisks and Planet Formation(overall disk properties) (overall disk properties)


Disks sizes and masses

Typical disk size ~ 200 AU Compares well with Solar System

Large scatter around median value!

Glauser et al. (2007)Stapelfeldt et al. (2003)

Kitamura et al. (2002)

HV TauIRAS 04158+2805

~ 40 AU

~ 1100 AU



Disk masses can be derived from thermal radio fluxes/maps Uncertain dust opacities Uncertain gas/dust ratio

Derived total masses: Consistent with MMSN Consistent with stability

Natta et al. (2000)



Radio interferometers (IRAM, OVRO) can resolve disks Typical surface density ~1 g.cm-3 @ 100AU Power law indices

– Temperature law– Surface density

‘Flat’ MMSN-like disks Good for planets! But interpolation…

Dutrey et al. (1996)


Disk asymmetries: large scales

Evidence for dynamical perturbation: Companion, planet, high-mass disk? What you see is NOT what you have…

Grady et al. (2001) Fukagawa et al. (2004)Piétu et al. (2005)

AB AurHD 100546

Optically thin!


Disk asymmetries: gaps

Planets embedded in disks open ‘gaps’ Can these be observed?

Gap size < 1AU High resolution + high contrast ALMA? New generation AO?

Remember, however: Spatial resolution remains an issue Gaps may be partly filled in


Disk dissipation

Using disk counts in independent SFRs provides survival time of inner disk Essentially nothing left after 10 Myr

No environment effect OB vs T associations,

clusters Large bodies may still

be present and hiddenMeyer et al. (2000)

Disk lifetime


Disk dissipation

Does disk dissipation depend on central object mass? Spitzer surveys of UpSco (~5Myr)

– G-B: 5 +/- 2 %– K0-M5: 19 +/- 3%– BDs: 37 +/- 9 %

Disk lifetime is longer for lower mass objects Because of slower viscous timescale?

Carpenter et al. (2005)

Scholtz et al. (2007)

}


Inner disk dispersal

Disks disappear after inner hole clearing Evidence shows that disks dissipate

inside-out in <105 yrs (viscous timescale)

McCabe et al. (2006)D’Alessio et al. (2005)

0.2-

0.5

AU

mat

eria

l

0.5-2 AU material

Very fewtransition objectsCoKu Tau 4


Inner disk dispersal

How long does the outer disk remain? Spitzer searches for disk with only outer

disk material (>5-10 AU) Only a few percent of such objects

Outer disk falls belowdetection threshold in <~ 105 yrs

Too fast for viscosity?

Padgett et al. (2006)


Disks and Planet FormationDisks and Planet Formation(dust properties) (dust properties)


Grain growth: mm view

First approach: SED slope (mm regime) Typically, amax ~ few mm to few cm

Natta et al. (2007)D’Alessio et al. (2001)

Smallgrains

Largegrains

Observeddistributionof spectral

indices


Grain growth: silicates view

Silicate feature is size-dependent Small (< 0.1 m) vs large grains (~1 m) Larger grains do not contribute

Crystallinity produces sharp features

Kessler-Silacci et al. (2006)


Grain growth: silicates view

Clear evolutionary sequence Larger grains come together with higher

grain crystallinity (above a threshold)

Kessler-Silacci et al. (2006) Van Boekel et al. (2005)

Hig

her

crys

talli

nity

Smaller grains


Grain growth: scattered light

Stellar photons can scatter off dust grains at the disk surface Phenomenon depends on /a Larger grains scatter preferentially forward,

with a lower polarization rate Images and polarization maps can be

used to infer grain sizes Up to amax ~ few m typically

Advantage: longer probes deeper!!


Grain growth: scattered light

Single power law size distribution Increasingly more isotropic scattering

HK Tau images (increasingly ‘peakier’) reveal larger grains inside (up to 3-5 m)

McCabe et al. (2003)McCabe et al. (in prep)

2.2 m 3.8 m 4.7 m 11.3 m

VLT/AO Keck/AO Keck/AO Keck


Grain growth: the big picture

Each aspect probes A different region of disks Different grain sizes/populations

In each case, analysis requires knowledge of additional information (radius, inclination, …)

Ideally, comparison all datasets to a single (complex) radiative transfer model


Vertical sedimentation

If large grains disappear from the surface, thermal equilibrium is changed Change in disk SED

Difficult to ascertain, however

Dullemond & Dominik et al. (2004)

Sedimentation mimics a flat disk



Confront mm regime and silicates Can be convincing (if composition is well

distributed throughout the disk)

Pinte et al. (in prep)

IM Lup

Small grainsonly

Small andlarge grains



Confront mm regime and silicates Can be convincing (if composition is well

distributed throughout the disk)

Pinte et al. (in prep)

IM Lup

Small grainsonly



sedimentation


Radial migration

Interferometry + spectroscopy (MIDI) Silicate features a few AU from the star Higher crystallinity!

Grain processing?

van Boekel et al. (2004)

Shegerer et al. (subm.)

RY Tau (K1)

small

crystalline


Radial migration

Difficult to quantify differentiation Many assumptions in analysis

Nonetheless, there is evidence that grain properties depend on radial distance to the star

However, we cannot prove that grains have migrated! Crystallinization may be a local processing


Further in time: debris disksFurther in time: debris disks


Debris disks: basics

Debris disks are the final stage in planet formation before zodiacal disks Formed through collisions of solid bodies

They are optically thin Easier to interpret Harder to observe

SED is usually limited Rough constraints only

Beichman et al. (2006)


Debris disks: porosity, aggregates With many independent observables,

finer models can be tested The AU Mic debris disk is made of (small)

porous grains

Fitzgerald et al. (2007)

porous

compact


Debris disks: porosity, aggregates Another debris disk: HD 181327 All observables cannot be explained

simultaneously with spherical grains Aggregates?

Schneider et al. (2006)

vs?

SED

Phasefunction


Summary and PerspectivesSummary and Perspectives


Summary

We have access to many types of complementary observations

Several physical processes can be (somewhat) constrained Core collapse/fragmentation Disk dissipation and inner hole clearing Grain growth Dust settling Presence of planetesimals


Perspectives

More observations will come with future instrumentation (e.g., ALMA)

At this stage, we still need Complex modeling/analysis of datasets More multi-technique analysis Tests of the basic processes in models

Wait for next talks, to get the theorists’ point of view!

Observations of Disks around Young Stellar Objects

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