High-angular resolution imaging of
protoplanetary disks…but, will cover statistical studies with low resolution

Misato Fukagawa(Osaka University)

Contents
1.

Introduction
 How to observe

2.

Statistics
Lifetime (frequency), disk massTransitional disks

3.

ALMA observations of the transitional disk around HD 142527

4.

Planets in disks

5.

Variability due to dynamical effectProtoplanetary disks Natural outcome of star formation

 Stellar age : 1~10 Myr
 Optically thick Gas-rich 100:1 in interstellar
 Formation of giant planets

 R ~ 100 AU ~1 arcsec in the nearest star-forming regions

Williams & Cieza (2011)

How disks evolve?
Dissipation mechanisms
 Viscous accretion

 Photoevaporation (e.g., Clarke et al. 2001)
 Grain growth, settling (e.g., Dullemond & Dominik 2005)
 MRI (e.g., Suzuki et al. 2010) Dynamical clearing by stellar companions
 Giant planet formation

What is happening there?
1.

Planet formation will occur, or is ongoing Initial condition

2.

Planet formation has recently finished interaction between planets and disks, evolution ofplanetary orbits, triggered formation of more planets?

 “evolution” + “diversity”
 Transformation of interstellar dust and gas
 Transport of material

What do we want to know?
Physical and chemical structure(r, z), T(r, z)…
 at various evolutionary stages

External radiation
X-ray, UV, optical for various stellar mass under various SF environments(stellar multiplicity, OB stars(external radiation),metallicity…)

mixingaccretion

radial drift

Observations of optically thick,protoplanetary disks(sub)mm

Thermal emission

FIRMIR

X-ray, UV, opticallog (F)NIRL*-1

LIR012log  (m)

3

1

10100 AU

Observations of optically thick,protoplanetary disks
Opt/NIR imaging(scat. light)

Dust scattered lightSensitive to (sub)micron-sizedin the upper surface↔ millimeter-sized grains
efficiently detected in mm
thermal obs.

1

10100 AU

Observations of optically thick,protoplanetary disks
Gas in IR & mm

Photondominated[CII] 157 um

X-ray, UV, optical[OI] 63 um
H2O

H2O
Ice (solid)

CO vib

CO high-J

10CO low-J100 AU

Warm mol.
layerMid-planeDisk propertyFrequency, excess in SED
Accretion onto the star

Observational diagnostic
Broad-band photometry
UV/optical/NIR spectroscopy

Dust mass
Gas mass
Rough geometry

mm/submm photometry
mm/submm/FIR/MIR spectroscopy
SED modeling

Detailed structureInterferometry, high-angular resolution
imaging
MIR spectroscopy

Dust mineralogy, size
distribution
Molecular abundance

mm/submm/FIR/MIR spectroscopyDirection
 Instrumentation, observing

technique High sensitivity
 High angular resolution
 High contrast

 Direction
 Obtaining more realisticpicture Case study + Statistics Generalization

 High dispersion
 High cadence Multi-wavelength High precision Multi-epoch

 Wide field
…

 Exploring a wider parameter

spaceStatistics:
Disk lifetime, mass evolution

Lifetime (of hot, inner disks) Disk frequency in clusters
 NIR excess (3.6 m) Spectral type: later than K5

 Frequency
 80% ( 100%) at ~1 Myr
 50% at ~2 Myr  0% at ~10 Myr

Hernandez et al. (2007)
Haisch et al. (2001), Hillenbrand et al. (2005),
Mamajek et al. (2004)

Lifetime (of hot, inner disks)
 Observations of youngclusters

 Dust – NIR excess (<10 m)
 Gas – H Timescale of disk dissipation

exp():  = 2—3 Myr

Fedele et al. (2010)IR excess vs. stellar-mass
 Clusters at >~ 2 Myr

2—3 Myr

Lada et al. (2006)
See also Hernandez et al. (2005), Dahm & Hillenbrand (2007),
Carpenter et al. (2009) For A/B-type stars(>~ 2 M), disk
dissipation is fasterthan K/M-type (0.1– 1 M)Disk (mm-dust) mass from dust continuum Uncertainty: opacity, T, gas/dust

Opacity vs. grain size𝐹 𝑑2𝑐 2 𝐹 𝑑2𝑀𝑑𝑖𝑠𝑘 ~~(in R-J limit)2𝐵(𝑇𝑑 ) 2𝑘   𝑇𝑑if    , 𝐹   ,  =  + 2

Williams & Cieza (2011)

Aikawa & Nomura (2006)
See also Draine (2006)

Disk mass evolution
Taurus
1—2 Myr Class II in Taurus,

Ophiuchus (1—2 Myr) Mdisk (median) ~ 5 MJup

Ophiuchus
1—2 Myr Mdisk/M* (median) ~0.5–1%

 Mass distribution is10 JupiterIC 3482—3 Myr

shifted to lower masses
by about a factor of 20for IC 348.

Lee et al. (2011)
Beckwith et al. (1990), Andrews & Williams (2005, 2007),
Williams et al. (2013)

Dust mass evolution

1—2 Myr
10 Jupiter(if G/D=100)

2—3 Myr~5 MyrMathews et al.
(2012)

Disk mass vs. stellar-mass
Mdisk  Mstar

 Taurus (1—2 Myr) Suggesting therelation: Mdisk  Mstar

 Scatter ~ 0.7 dex

Andrews et al. (2013)
See also Mohanty et al. (2013)

Dissipation timescale vs. stellar-mass
 U Sco (~5 Myr) Need to improvestatistics

 Disk dissipation isfaster for M latestars than K stars?
(recent talk by J.
Carpenter; ALMA
provided highersensitivity by ~2orders of mag.)

Mathews et al. (2012)

Caution: Are ages correct…?
IC 348~6 Myr

Bell et al. (2013)

IC 3482—3 Myr

Hernandez et al. (2007)

External radiation fieldOrion

Tau/OphMann & Williams (2010)

Multiplicity
 Binaries with ap = 30–300

AU are 5× fainter Although very few, closebinaries (ap < 30 AU) are
5× fainter Exceptions: a few brightcircumbinary disks

Flux at 0.9 or 1.3 mm for individualcomponents of 23 multiples in
Taurus (Harris et al. 2012)

Statistics:
Transitional disks

Disks in transition Classified based on SEDs (Strom et al. 1989)“THE DISKIONARY” by Evans et al. (arXiv:0901.1691)

Williams & Cieza (2011)Frequency of TD Rare population, ~10% (order)

Kim et al. (2013), Muzerolle et al. (2010)

 t ~ 0.3 Myr (Skrutskie et al. 1990)

Similar results have been obtainedin later studies (Luhman et al. 2010).
 ~100 disks in 1—3 Myr(62 in ONC/L1641; Kim et al. 2013)
 No difference between K and M
type stars (Koepferl, Ercolano et al.
2012) Complications
 Definition of classification,frequency Weak excess sources Young debris disks?

 No hole? (need longer )

See also Najita et al. (2007), Cieza et al. (2008,2010), Currie et al. (2009), Merin et al. (2010)

With or without a hole/gap
1.

Accreting objects, Ld/L* > 0.001

Settled(homologously depleted, anemic, grain-growthdominated)

2.

Disks with gaps or holes(classical transitional, cold disks, dynamically-cleared)

 The incidence of signatures of dynamical clearing

by giant planets is significantly lower than theoccurrence of giant planets within 20 AU ofmature solar-type stars (∼5% versus ∼20%).
 The fraction of grain-growth-dominated disks
decreases with time.
 Disks with signatures of dynamical clearing reveals
a lack of such objects among the youngest stars
in the sample.

Cieza et al. (2012)

Evolutionary paths
flared, primordial

truncated disk?
planet formation finished?
star
1100

Settled

t

1100

1

λ (µm)λ (µm)disk
100

λ (µm)?

?

Gap/hole1100

λ (µm)

Evolutionary paths
flared, primordial

truncated disk?
planet formation finished?
star
1100

1λ (µm)t

100

λ (µm)

100?

Settled1disk?

λ (µm)

Gap/hole1100

λ (µm)

Two populations
Important quantities:
Disk mass, accretion rate, Rwall

mm-bright and mm-faint mm-bright show;

mm-faintprimordial

mm-bright
Owen &Clarke(2012) higher accretion rates

 larger Rwall

 Some mm-bright are very

massive; Mdisk ~10% of M* mm-bright are preferably seenaround early type stars
 mm-faint are consistent withphotoevaporation

Owen etal. (2011)

Identifying gaps/holes is sometimes noteasy from SEDs

Maaskant et al. (2013)

Case study (imaging):
Transitional disks

Planet footprint?
Disks with holes or radial gaps
 Hole size > 20 AU
 Carved by planets, or by othermechanisms?

MIR dip

submm
dust cont.

Williams & Cieza (2011)

Disks in high-angular resolution(NIR scattered light, ~0.06”)
Radial gap
Dark in scat. light• Devoid ofmaterial• Vertical disk
structure(cannotintercept stellarphotons)
In some cases,
submm cavities
are filled withsmall grains

Spiral0.5 = 70 AU

Spiral wake by aplanet?

(Muto+ 2012)(Mayama+ 2012)

0.50.5(Hashimoto+ 2012)(Armitage 2005)(Grady+ 2013)

Disks in high-angular-resolution(submm, ~0.3”—0.5”)
Case study: HD 142527 with ALMA
Discussion with:
T. Tsukagoshi, M. Momose, K. Saigo, N. Ohashi, Y. Kitamura, S.-I. Inutsuka,T. Muto, H. Nomura, T. Takeuchi, H. Kobayashi, T. Hanawa, E. Akiyama, M.
Honda, H. Fujiwara, A. Kataoka, S. Z. Takahashi, H. Shibai

HD 142527
 Herbig Fe star (F6, Teff = 6250 K), d ~ 140 pc
 L = 20 L, M = 2.2 M, Age ~5 Myr(Verhoeff et al. 2011)

 Abundant IR excess (LIR ~ L*) Imaging – one of the peculiar disks

far-side

1.6 μm, scattered light(Fukagawa et al. 2006)

24.5 μm, thermal(Fujiwara et al. 2006)

near-sideThe disk(Verhoeff et al. 2011)

1. Inner disk (r < 30 AU)
 High crystallinity, large grains(Bouwman et al. 2001, van Boekel 2004)

2. Gap (r ~ 100 AU)
 Gap seems too wide to beaffected by the stellar companion
inferred at ~12 AU (Biller+ 2012)
3. Outer disk (r > 100 AU)
 Multiple spirals + an outer arm

Outerdisk

h= 60 AU at
r = 130 AU

Another motivation toward ALMAStrong non-axisymmetry:“horseshoe” disk
(1) Very wide (~1”) cavity(2) Enhancement in the northernarea→ One of the best protoplanetarydisks for ALMA early science
Contours: SMA 880 μm continuum(from 2, spacing of 2)(Ohashi 2008)

ALMA’s results:
Gap-crossingflows?
 HCO+, dust streamers(Casassus et al. 2013)
for press release

Our observationsBand 7 (336 GHz), covering 13CO(J=3-2), C18O(J=3-2)
 Optically thinner lines than 12CO20–26 12-m antennasBeam size = 0.5” x 0.4” (70 x 60 AU)

Effective velocity resolution ~0.2 km/s

Beam size

Continuum

13CO (3-2)

C18O (3-2)

0.50” x 0.41”

0.52” x 0.43”

0.53” x 0.43”10(per channel)

11(per channel)

Rms noise
0.2(mJy/beam)

Continuum at 340 GHz
 Inner disk + gap (non-detection) + outer diskOuter “Horseshoe” : N-Sbrightness contrast ~ 25

Peak = 280 mJy/beam,brightness temp:
Tb ~ 22 K1”
140 AU

Inner disk and radial gap in dust Inner disk

 detected at 10 
 spatially unresolved disk mass = ~4  105 M(lower limit since estimated under
the optically-thin assumption)

Similar compact emission wasdetected toward AB Aur.

 Macc <~ 7 x 10-8 Myr-1(Garcia Lopez et al. 2006; SNR seems low) Gap
 non-detection; 1 = 0.3 mJy/beam No clear emission corresponding to
the HCO+ flow

AB Aur(Tang et al. 2012)

Structure Radial brightness distribution is well

represented by Gaussian
 ring, torus…

Sharp outer/inner edges intransitional disks (e.g., Isella et al.
2010, Andrews et al. 2011, Tang etal. 2012) FWHM in ris narrower at the northernhorseshoe peakRadial cut

StructurePA

 Anti-correlation between the(13CO out to ~450 AU
 ~300 AU for dust )

200

20190

15180

10170

5160

0150

peak

(peak) (g cm-2)

 “Ring + floor”

25 Horseshoe

060 120 180 240 300

PA from major-axis(G/D = 100 is assumed in )(Ataiee et al. 2013)

Peak Position (AU)peak intensity versus peak position
for the fitted GaussianNot an eccentric disk (higherdensity expected at apastron)

Molecular lines
 Integrated intensity:

ring-like distribution

 Velocity field follows

13CO

Keplerian rotation, with (J=3-2)
vel. resolution of 0.2km/s
 major-axis PA = 19
 i = 27  2 for M* = 2.2

 0.3 M

C18O
(J=3-2)integrated I

velocity

1”1”

Tb

13
for CO(J=3-2)

 13CO emission should be
optically thick 13CO/C18O << canonicalvalue (~7; ISM, HD
163296; Qi et al. 2011)
 Tb(13CO)  Tb(C18O) = 5 K;in N, 130 < r < 300 AU

 Tb reflects temperature, notcolumn density

Surface density of the horseshoe
How massive does the disk at the horseshoe peak?
 Tb = 36 K of 13CO provides the upper bound for temperature ofdust grains emitting in submm(13CO must be located in the upper layer than the dust) (dust) = 0.22 g cm2 at the peak (dust) = 0.087 g cm2 at 161 AU,
azimuthally-averaged Mdisk = 0.08 M
 T = const. = 36 K opacity = 0.034 cm2/g ( = 1.0)
 optical depth estimated for 36 K

Enhancement toward North: Implications
Discussion using the lower limit of 
1. Self-gravitating?
 Gas-to-dust = 100
 Toomre Q for a Keplerian disk
 Q ~ 3 for  = 8.7 g cm2 azimuthally-averaged (160AU)
 Q ~ 1 for  = 22 g cm2 at the density peak
… Gravitationally unstable (!)
Looking at the disk in a transient phase? (dynamical 103 yr,age 106 yr)Enhancement toward North: Implications
2. Dust accumulation that lowers gas-to-dust
 Gas-to-dust < 100, at least ~10 makes Q >> 1
 Particle trapping at pressure maximae.g., large-scale vortices (e.g., Birnstiel et al.
2013, Ataiee et al. 2013)

In either case, the enormous pile-up of disk
material beyond 100 AU from the star planet forming site?

ALMA simulated image(Birnstiel et al. 2013)

Uniqueness?
 Azimuthal asymmetry; x 10 compared to other ring-like disks

 Relatively massive disk (>~5% of the stellar mass)

J1604-2130(Mathews et al. 2012)

AB Aur(Tang et al. 2012)

etc…

Another one with strong asymmetry
Azimuthal asymmetry (x100 in flux density)
Forming site of rocky objects? (9 earth mass for dust)(van der Marel+ 2013)

Multi-wavelength data
Larger grains in the mm-brighter
region.
Segregation of dust particle sizeis also confirmed radially forsome disks using VLA at ~1 cm(Perez et al. 2012)(van der Marel+ 2013)
submm thermal

MIR thermal

NIR scat. light

Better constraint on disk
spatial structure, dependence of opacity

Azimuthal asymmetry

Isella et al. (2013)
For gap-making by multiple planets, see also
Zhu et al. (2011), Dodson-Robinson & Salyk (2011)

Azimuthal
asymmetry:
Why?

Gas-to-dust
mass ratio, andits spatial
variation

Grain growth:
leads to planet
formation?

Inner disk:
needs supply fromthe outer disk?
Gas amountin radial gap

Modeling
Bruderer et al. (2013)

Mass determination
 Dust Temperature CO snowline (Qi et al. 2013, Mathews et al. 2013)

 Opacity (absolute value) Reasonable in a statistical sense? (The upper envelope of Mdisk is ~0.1 M*)
 How about for individual disks?

 Gas
 CO

temperature, abundance, freeze-out effect less uncertainty: HD, H2…

 Indirect methods
 Dynamics, modeling, …

Does a planet(s) really exist in
the large hole?
(observationally)

Possible protoplanets in disks
A newly formed planetwith ~6—15 MJup ,surrounded by dusty
material.
De-projected R~20 AU.

850 micron, dust cont.
Rgap = 55 AU

Blue: K’, red: L’

Proper motion of the emission structure in 2012 (compared to2010), consistent with Keplerian rotation at 20 AU (Ireland & Kraus)

Possible protoplanets in disks
Resolution (~15–20 mas; 2–3 AU)

Kraus et al. (2011)

Possible protoplanets in disks
Quanz et al. (2013)

L’ band (3.8 micron)Separation = 68 AU
Mass ~15—20 MJup

Need deeper search?

Casassus et al. (2013)

5 σ detection limits in Ks and K’
15–20 MJ at 0.25” , 8–10 MJ at 0.5”
3–4 MJ at 1”,
Grady et al. (2013)

Detecting forming planets Pros Primordial: less affected by orbital

migration
 Probably brighter Cons PMS are far away (>100 pc)

 Disk becomes noise source Luminosity evolution is notcalibrated
 Effect from accretion is unknownMarley et al. (2007)

Transitional disks are variable:
Temporal variation of disk structureIR variability in transitional disks

Muzerolle et al. (2009)

 LRLL 31 in IC 348
 Macc ~ 2 x 10-8 M /yr, “Gapped” disk
 “seesaw” behavior
 Timescale ~1 week “Seesaw” may be common only among
disks with large radial gaps/holes Change in the scale height by ~20%

variable shadowing(e.g., Juhász et al. 2007) Espaillat et al. (2010)

See also Flaherty et al. (2011, 2012)

Variability in IR Most YSOs are variable in NIR(typically ~0.2 mag).

(Carpeter et al. 2001, Sitko et al. 2008, MoralesCalderon et al. 2009,2011, Kospal et al. 2012,
Megeath et al. 2012) Dynamical effect?
 Periodic (25 day), pulsed

accretion for a protostar:
Variable accretion?

 Emission from dust streamers fromcircumbinary to circumstellar?
(Nagel et al. 2012)

Muzerolle et al. (2013)

Summary Direction
 Obtaining more realistic picture Case study + Statistics  Generalization

Multi-wavelength, Multi-epoch, Exploring a wider parameter space Some problems
 Age (clock) Dust, gas mass (Gas-to-dust mass ratio) Luminosity evolution of forming planets (detection feasibility) High-resolution, multi-wavelength data

Interpretation by modeling (both for dust and gas) is important