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43. “ASSESSING MOLECULAR OUTFLOWS AND TURBULENCE IN THE
PROTOSTELLAR CLUSTER SERPENS SOUTH”!
A. L. Plunkett et al.!
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ApJ!

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42. “The formation of a quadruple star system with wide separation”!
Jaime E. Pineda et al.!
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http://ads.nao.ac.jp/abs/2015Natur.518..213P!
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Nature Letter!

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41. “Another deep dimming of the classical T Tauri star RW Aur A”!
Peter P. Petrov et al.!
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A&A accepted!

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44. “SMA OBSERVATIONS OF W3(OH) COMPLEX: PHYSICAL AND CHEMICALDIFFERENTIATION BETWEEN W3(H2O) AND W3(OH)”!
Sheng-Li Qin et al.!
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ApJ accepted!

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45. ”Protoplanetary disk lifetimes vs. stellar mass and possible implications for giant planet
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Álvaro Ribas et al.!
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The Astrophysical Journal, 803:22 (24pp), 2015 April 10

doi:10.1088/0004-637X/803/1/22© 2015. The American Astronomical Society. All rights reserved.

The Astrophysical Journal, 803:22 (24pp), 2015 April 10

Plunkett et al.

ASSESSING MOLECULAR OUTFLOWS AND TURBULENCE IN THE PROTOSTELLAR CLUSTER SERPENS
SOUTH
1
3
43.
“ASSESSING
AND 2TURBULENCE
IN THE
Adele
L. Plunkett1,5MOLECULAR, Héctor G. ArceOUTFLOWS, Stuartt A. Corder, Michael M. Dunham, Guido Garay4, and
PROTOSTELLAR CLUSTER SERPENS
SOUTH”!
Diego
Mardones4
1
of Astronomy, Yale University, P.O. Box 208101, New Haven CT 06520, USA
A. L. Plunkett etDepartment
al.!
2
Joint ALMA Observatory, Av. Alonso de Córdova 3107, Vitacura, Santiago, Chile3!
http://ads.nao.ac.jp/abs/2015ApJ...803...22P!
Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, MS 78, Cambridge, MA 02138, USA
4
Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile!
ApJ!
Received 2014 August 27; accepted 2015 February 10; published 2015 April 9!

ABSTRACT
Molecular outﬂows driven by protostellar cluster members likely impact their surroundings and contribute toturbulence, affecting subsequent star formation. The very young Serpens South cluster consists of a particularlyhigh density and fraction of protostars, yielding a relevant case study for protostellar outﬂows and their impact on
the cluster environment. We combined CO J  1 0 observations of this region using the Combined Array for
Research in Millimeter-wave Astronomy and the Institut de Radioastronomie Millimétrique 30 m single-dishtelescope. The combined map allows us to probe CO outﬂows within the central, most active region at size scalesof 0.01–0.8 pc. We account for effects of line opacity and excitation temperature variations by incorporating 12CO
and 13CO data for the J = 1 − 0 and J = 3 − 2 transitions (using Atacama Pathﬁnder Experiment and Caltech
Submillimeter Observatory observations for the higher CO transitions), and we calculate mass, momentum, andenergy of the molecular outﬂows in this region. The outﬂow mass-loss rate, force, and luminosity, compared withdiagnostics of turbulence and gravity, suggest that outﬂows drive a sufﬁcient amount of energy to sustain
turbulence, but not enough energy to substantially counter the gravitational potential energy and disrupt the clump.
Further, we compare Serpens South with the slightly more evolved cluster NGC 1333, and we propose anempirical scenario for outﬂow-cluster interaction at different evolutionary stages.
Key words: ISM: individual objects (Serpens South) – ISM: jets and outﬂows – ISM: molecules – stars: formation– stars: protostars – techniques: interferometricProtostellar1.clusterのOutflowは周辺環境に影響を与えることが期待され、また乱流に寄与
components within the protostellar cluster Serpens South in
INTRODUCTION

order to determine the impact of outﬂows on the surrounding
しその後の星形成に影響することが考えられる。!
1.1.
Molecular Outﬂows and Protostellar Clusters
cloud environment.
Outﬂows
and
jets
occur
duringthe
earlystagesofOutﬂows may reach distances up to several parsecs, affecting
対象；Serpens South の中心部0.01–0.8 pc!
protostellar evolution, and they facilitate the accretion of
the surrounding protostellar environment corresponding to
観測；CARMA,
30m,away
APEX,
CSO!
material
onto the protostar IRAM
by carrying
excessangularthese scales, including the core and cloud. Since outﬂows
momentum (Bachiller 1996). A molecular outﬂow consists ofsweep up ambient gas, they inﬂuence the intimate relationLine;
12CO(1-0),
13CO(1-0),
12CO(3-2),
13CO(3-2)!
ambient gas that is entrained by a protostellar jet, expected to
between a protostar and its surroundings. Considering that the
これらの輝線を用いて、光学的厚みと励起温度を考慮して、outflouの質量、momentum,
occur during the Class 0 and I stages of star formation. Class IImajority of clustered star formation occurs within clusters ofsources also power jets and winds and may drive outﬂows;
100 or more members, and the size scales of clusters areenergyを求めた。!
however,molecular outﬂows are less prominent as a source
typically less than a few parsecs (Lada & Lada 2003), there isevolvesand the surrounding
molecular gas becomes depleted.
ample opportunity for outﬂows to encounter other nearby
Outflowenergyは乱流を維持するのに十分であるが、重力ポテンシャルに対してクランプ
Observationally, outﬂows serve as signposts that emanate fromoutﬂows and their respective protostars, as well as potentiallydisrupt the surrounding cloud environment.
their を壊すには及ばないことが示唆される。!
driving protostellar sources and indicate that star
It remains a challenge to quantify the cumulative impact offormation is ongoing in a particular region. Further, the
また、NGC1333と今回のSerpens
Southの結果を比較し異なる進化段階でのOutflowprotostellar outﬂows on molecular clouds and subsequent starcharacteristics and morphology of outﬂows may indicate theformation,but evidence of outﬂow-cloud interaction has been
evolution of the forming protostar and constrain the gas
Cluster interactionの観測に基づいたシナリオを提案する。!
observed in a variety of star-forming regions (e.g., Arce
entrainment mechanism (Arce et al. 2007).
et al. 2010; Nakamura et al. 2011a, 2011b; Duarte-Cabral
Carbon monoxide is the most common probe !of molecular
The Astrophysical10
Plunkett
et al.
et al. 2012).
Speciﬁcally,outﬂows interact with the ambient
it is 2015
theAprilmostprevalent molecule inoutﬂows,
because Journal,
after 803:22
H2 (24pp),
h Aprils
AstrophysicalJournal, 803:22(24pp),
2015
Plunkett
et al.
30m(Fuller0010, decl. &= Ladd 2002; Arce & Sargent 2006),
OurThemosaicwas centeredat R.A.
= 18
clustergas
injectmolecular clouds and its low-level rotational statesareeasily2n2a0´ and consisted of 66 pointings in a hexagonal-packed
momentum
and withenergy into the surrounding cloud (Arce
excited in these environments (e.g., Wilson et al.pattern;2009).
The
the map
achieved better
than Nyquist samplinghorizontal
spacingofandverticalspacingof.
In
total,30´25.
´
8
et al. 2007; Swift & Welch 2008), and they may feed turbulent
CO molecular outﬂow components are generally characterized
we observed 33.8 hr, with 24.6 hr on source. Integrations werethatas standard high-velocity (SHV) with velocities 30ofs pera mosaic
few position,
to and 3motions
minutes on thephaseaffectcalibrator ongoing star formation (see also Frank
1743-038, with 33 pointings observed in between each phaseet al. 2014). Models of clustered star formation must
about 20 km s−1, and extremely high velocity (EHV)
with
even
calibrator observation. Uranus was the primary ﬂux calibrator.
We simultaneously
the J = 1 − 0 outﬂows
transitions of in order to drive turbulence and produce
incorporate
Bachil- observed
higher velocities, typically up to ∼100 km s−1 12(seeCO (115.27 GHz) and 13CO (110.20 GHz) in the upper side12
to ofwhat is observed in star-forming regions,ler 1996). Typically the SHV outﬂows are more massive
thanwindows ofoutcomesband. The spectral
CO and 13COclose
had widths−1
31 MHz (_80 km s−1) and 8 MHz
(
km
s
),
_20
including starrespectively.
formation rates (SFRs) and efﬁciencies (e.g., Li
EHV, and in this paper we conduct a census of With
mostlySHV
383 channels in each CO window, the velocity−1resolutions were 0.2 and 0.06 &km s Nakamura, respectively. In addition,
2006; Nakamura & Li 2007, 2011; Carroll
we simultaneously observed the continuum emission using
5
etMHz
al.each,2009;
Wang
NSF Graduate Research Fellow.
eight spectral windows of 495
resultingin a mapet al. 2010; Krumholz et al. 2014). The

Figure 1. Coverage of observations, overlaid on a map of Spitzer 8 μm
emission (Gutermuth et al. 2008). Blue (solid) inner contour marks the
CARMA mosaic map where Nyquist sampling and best sensitivity wereachieved, and blue (solid) outer contour marks degraded sensitivity at the edgeof the map. The inner contour of this map comprises the majority of molecular
outﬂow emission. Green (dashed) contour marks the IRAM 30 m mapcoverage, which is composed of nine 2a q 2a and one 2a q 1a OTF maps.
Yellow (dash–dot) contour marks the APEX 12COJ = 3 − 2 map. Orange(dashed) circles mark four CSO pointings of 13CO J = 3 − 2 observations, anddiameters correspond to the CSO beam.

ページ
3/12
Caltech Submillimeter Observatory (CSO), as well as theprocess of combining interferometer and single-dish data of thesame isotopologue. To put the observations into context, inFigure 1 we show the map coverages that we describe here;
Table 1 gives a summary of the observations.
We also observed the J = 3 − 2 energy level transition of13
CO with the CSO, and 12CO with APEX. Observations of
this higher energy transition are utilized to determine
velocity-dependent opacity described in Appendix A and
excitation temperature according to the method described inAppendix B.

with a total continuum bandwidth of 4 GHz that covers thesame region as that of CO. The beam size of _5´ corresponds
10.01 pc, or 2000 AU (at a distance of
to a physical resolution of
429 pc).
Data were reduced using MIRIAD (Sault et al. 1995).
Anomalous amplitudes and phases were ﬂagged, as were any
observations
systemhigher
K.8No
Figure 2. (a)with
Serpens
Southtemperatures
region as observedwith thanSpitzer500
IRAC
μm (Gutermuth et al. 2008), overlaid with sources from the literature according to theirclassiﬁcations
Maury used,
et al. (2011,
star symbols) whenand in the
shadowed
databywere
nor shown
were withobservations
theSpitzer Gould Belt survey (Evans et al. 2007; Gutermuth et al. 2008, shown with circles).
The white
box marksbetween
the region85°
where
the 90°.
majorityof COsourcewas dashedat elevations
andFluxandoutﬂows
gain are found, shown in Figure 7. The red dashed box shows the region where we detect
continuum sources, shown in more detail in the right panel. The 2.7 mm continuum sources (see Table 2) are given by crosses and diamonds, such that crosses
calibrations
wereapplied,
and
the
data
were
Fourier
indicate previously identiﬁed YSO candidates (Evans et al. 2007; Maury et al. 2011), and orange diamonds mark CARMA-2 and CARMA-4, which previously weretransformed.
not identiﬁed as YSO candidates. (b) 2.7 mm continuum (contours) observed with CARMA, overlaid on the Spitzer 24 μm map (Gutermuth et al. 2008). Noise rms ofthe continuum
mapdatais 1 mJy/beam,and the beam
is shown
in the invert,
lower right corner. Contours begin at a level of 4T and increment by 8T . Overlaid are the positions of
Imagesof thewere obtained
using
the task
2.7 mm continuum
sources,
described
for thedeconvolved
left panel.
specifying
a cell sizeof as2″,and thenwith themaximum entropy deconvolution mosmem. The clean images
2
were then restored with a Gaussian beam. We subtracted Continuum
the Table Sources
continuum image from the molecular line maps using
Position A summary ofSize
avmaths
the CARMA
Other Names
Source in the image domain.
P.A.
Peak Intensity
Total Flux
Mass
c2da
observations is given
in Table 1.
α(J2000)
δ(J2000)
Maj (″) Min (″)(°)(mJy/beam)(mJy)(Me) Classiﬁcation
4b, Y1c

38.9
−61.6
18.6

3.6 ± 0.8
3.6 ± 1.3
17.5 ± 2.4

7.1
7.5
39.1

0.07
0.07
0.36

YSOc
L
L

We
mappedregion−02:02:48.3
with the IRAM
telescope
CARMA-3
SWthe same
18:30:02.5
4.0 30 m 3.0
−87.4
h
s,
CARMA-3
NEdThe18:30:02.6
4.0 = 183.1
in 2012April.
center of −02:02:41.4
the map is R.A.
30 m0380.1
decl. = 02n02a to match the CARMA mapf coverage. We
CARMA-4
18:30:02.8−02:04:08.8
L
L
L
mapped the region using individual 2a q 2a on-the-ﬂy (OTF)

19.1 ± 2.4
6.4 ± 2.1

29.5
10.0

0.27
0.09

L
L

L
L

6.2 ± 1.7

7.1

0.07g

Galc_redL

CARMA-1
CARMA-2
CARMA-3d

18:30:01.2−02:03:43.3
6.6
18:30:01.4−02:01:58.7
6.5
2.2. IRAM −02:02:47.5
Observations 8.0
18:30:02.5
d

2.9
3.3
3.2

L
SerpS-MM16e

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Figure 3. Integrated intensity maps of the blue- and redshifted outﬂow emission in Serpens South traced by (a) 12CO J = 1 − 0 emission observed with CARMA+IRAM, and (b) 12CO J = 3 − 2 emission observed with APEX, with beam sizes for each given in lower left corners. The scale bar applies for both panels. Grayscaleshows integrated intensity with 3  ∣ vout ∣  31 km s−1, and blue and red contours show the blue- and redshifted moment zero maps within this velocity range.
Contours begin at 5T for each respective map and increase in intervals of 10T . Green contours show 2.7 mm continuum observed with CARMA, beginning with 4T
and increasing by 8T . Crosses mark the 45 protostars in the region from the Spitzer Gould Belt survey (Evans et al. 2007; Gutermuth et al. 2008), and these are also
shown in Figure 2 according to their classiﬁcations.
The Astrophysical Journal, 803:22 (24pp), 2015 April 10

Plunkett et al.

3.2.1. Blueshifted Outﬂows∼10–20 protostars that reside in the central region but remain
undetected in continuum observations.
Prevalent blueshifted CO features appear northeast and south
The northern wall of the blue northeast outﬂow structure isof the continuum sources CARMA-3/CARMA-5 andcomprised of clumpy emission (features A and B in Figures
CARMA-6/CARMA-7, shown in Figures 3–4 and the upper
7(b)–(c)) along an axis with position angle of _20n, appearing
row of Figures 7–8. In the lower blue velocity channels (i.e.,
to originate from near the sources CARMA-3 and CARMA-5
especially with approximately vout  6 to −4 km s−1), theseand extend _120´ (0.25 pc) northeast. At blue velocities withfeatures appear as wide-angle outﬂows with opening angles of∣ vout ∣ 2 17 km s−1, the outﬂow is comprised of delineated∼70° and 45° and position angles of ∼45° and 180°,
clumpy features that form part of the structure’s wall. Anotherrespectively. However, they may in fact be several moreclumpy feature (C in Figure 7(b)) is seen up to ∣ vout ∣  15collimated outﬂows superimposed on one another (seekm s−1 at a location 45″ (0.1 pc) northeast of CARMA-3 andFigures 4, 7, and 8). These outﬂowing structures correspond
CARMA-5 (along an axis that intersects CARMA-5 withapproximately to those identiﬁed as B1, B2, B3, and B4 by
position angle _60n). Outﬂow emission in the northeast is seen
with ∣ vout ∣ up to ∼30 km s−1 (Figures 7(a)–(c)), whereas the
Nakamura et al. (2011a), which we have also labeled insouthern blue outﬂow shows emission up to ∣ vout ∣ _ 10
Figure 7(a) for reference. The outﬂow features in low-velocitykms−1 (Figure 7(a)).
channels suggest that B1 and B2 comprise the cavity walls of a
These clumpy features may be evidence of episodic ejection
single northeastern outﬂow lobe, and similarly B3 and B4
from one or more sources in the central region, but we cannot
comprise the walls of a southern outﬂow lobe. Wide-angle, farbe certain of whether these come from one or more driving
reaching SHV outﬂows with cavity walls that are nearlysources in that region. We also note that the clumpy and bulletparallel at larger distances from the driving sources have been
like blueshifted emission in high-velocity channels is sufﬁobserved in other regions, including IRAS2-south in NGC 1333
ciently strong and localized that negative sidelobes appear at(Plunkett et al. 2013), HH 46/47 (Arce et al. 2013), and RNO
their edges due to the synthesized beam, as seen in the upper
91 (Lee & Ho 2005). In the same direction as the outﬂows B1
left panels of Figure 4. These negative sidelobes have absoluteand B2 to the northeast, we see that the peaks of 13CO emission
values that are less than 4σ in individual channels, and they arecorrespond to the strongest 12CO emission, shown in
not included in any calculations of mass, momentum, or energy
Figure 6(a) (see also Figure 5), indicating entrained gas even(see Section 3.3).
at lower velocities than what is probed by the optically thick
Several “infrared H-H objects” (e.g., described by Nakamura
12
CO. We note also that since protostars in Serpens South are
et al. 2011a) and H2 features (Teixeira et al. 2012) may bein such close proximity, and given that we expect eachassociated with the CO outﬂows discussed here. Nakamura
protostellar source to drive a bipolar outﬂow, it is likely thatet al. (2011a) report the infrared H-H objects, likely resulting
what appear as broad or wide-angle outﬂows are in factfrom shocks by protostellar outﬂows, which they call K1 andcomprised ofFigure
severalcollimatedthe7. 12COmore
outﬂow emission,
showingoutﬂows
low-velocity from(left), mid-velocity(middle),and high-velocity(right) outﬂowing
features.
Blue H-H
and redobjects
contours have
K2 toward
the southwestof the gas
map.
These
represent sources
blue- and redshiftedgas, and here
velocityor
intervals
aretheindicatedfor each panel, where outﬂow velocity, vout, is relative to the central cloud velocity of 8 km s−1.
distinct continuum
mentioned
from
other
bow shapes indicating that they are moving toward theContours begin at 5σ for each zeroth-moment map and increment by 15σ. The beamsize and scale bar are given in the lower right corner. Background is Spitzer 8 μm
emission (Gutermuth et al. 2008). Boxes 1, 2, and 3 refer to the regions shown in Figures 8(a)–(f) and 9, respectively. Ellipses and the arrow with solid cyan and pinklines indicate outﬂow features mentioned in the text, and plus signs and diamonds
mark the positions of continuum sources detected by our CARMA observations,8labeled in panel (c). K1 and K2 are IR H-H objects identiﬁed by Nakamura et al. (2011a) and correspond to the southwestern tips of the molecular hydrogen objects
MHO3249b and MHO3250 identiﬁed by Teixeira et al. (2012); P4 is an IR source. MHO features from Teixeira et al. (2012, see their Figure 1) are indicated withmagenta dashed lines and ellipses.

channels in the vicinity of CARMA-1 and we found

3.3. Mass, Momentum, and Energy of Outﬂowing Gas

13
The Astrophysical Journal,
803:22in(24pp),
2015
10 of CARMA-1 (seeCO emission,
particular,
to Aprilthe north

Plunkett et al.

The mass, momentum, and energy of outﬂowing gas were
Figure 6), although the emission is somewhat diffuse and does
calculated in order to assess the overall contribution of outﬂow
Table 3 activity within the mapped region, and we compared these withnot conclusively point to this source.
In summary, here we provide additional evidence
oneMomentum,orOutﬂow that
Mass,
Energyturbulent
momentum and energy and gravitational potential
more molecular outﬂows are being driven from the location ofenergy. First, we calculated H2 column density (NH2 ) for each
Mass
CARMA-1. However, a clear picture of a bipolar outﬂow is not
MethodMomentumpixel ( x, y ) and velocity channel (with channel
width %v ) as a
immediately apparent. Further, the collimated low-velocitya−1 to the(Tex)(Maccording(M : ) of excitation
Redb temperature(M :)(M :)lobe, does Bluefunction: km s )blueshifted feature F, which is a candidate All
outﬂow
following relations (following Dunham et al. 2014):
not have+anchannel-by-channel
equally collimated ccounterpart in the
2.8red. It seems
1.2
1.6
13.6
(1) Tex(x, y, v), pixel-by-pixel*that either
NH2 ( x , y , v , Tex )  f ( J , Tex, X CO ) Tmb,12 %v(2)c the outﬂow structure of this source is irregular or the(2) Tex(v), channel-by-channel
2.9
1.2
1.7
13.9
emission is being overwhelmed by stronger emission from the(3) Tex = 10 K
2.2appear more
1.1
1.2
Tex ) EJ 1
surrounding region, which makes the outﬂow
3k 2J 1 Q (10.8
1clumpy rather than collimated. Alternatively, if4.6
gas has already(4) Tex = 30 K
2.3 f ( J , Tex, XCO )  XCO,12
2.4 8π 3Oμ2 J 1 g22.4e kTex . (3)Lbeen swept away from the immediate region by(5) Tex = 50 K
7.2other stronger
3.5
3.7
34.7
outﬂows such as those to the northeast, there remains little gas
The opacity correction is explained in Appendix A, with theto be entrained by this driving source CARMA-1.
main-beam temperature of 12CO corrected for optical depth,abcBlue channels correspond to ∣ vout ∣ = 3.3–30.9 km s−1.
Red channels correspond to ∣ vout ∣ = 3.0–30.6 km s−1.
Means: Tex (x, y, v )  17 o 8 K for (1), and Tex (v )  15 o 4 K for (2).
The Astrophysical Journal, 803:22 (24pp), 2015 April 10

given by Equation (A4). Appendix B describes themethod(s) to solve for excitation temperature, Tex. In the above
equations, J = 0 is the lower energy level of the rotationaltransition, and gJ  2J 1  1 is the degeneracy of the lower
state. For XCO we assume a standard abundance ratio[12CO H 2 ]  104 (Frerking et al. 1982). For 12CO J = 1 − 0,
O  115.27 GHz, and μ is the 12CO dipole moment with avalue 0.11 Debye. Q (Tex ) is the partition function, which hereis approximated as Q (Tex )  kTex hB, and B = 57.635 GHz is
the rotation constant for CO. E J 1  7.64 q 1016 erg is theenergy of the upper level such that the energy of the transition
is E  E J 1  E J , and E J 1 k  5.53 K.
Having measured column density for each pixel, mass is thendetermined to be*,
Tmb,12

Energy(×1045 erg)
1.0
1.0
0.8
1.7
2.6

12
Plunkett et al.

In Figure 10 we show the outﬂow mass distribution per
velocity bin for Serpens South (panels a and b). For the blueand redshifted mass spectra, we ﬁt power laws of the formm (v ) r vH for the velocity range 3  ∣ vout ∣  15 km s−1, andwe found slopes of H  3.5 o 0.1 and H  3.3 o 0.1,
respectively. The slope may be an evolutionary indicator ofmass ejection in a cloud, such that the slope steepens with time
(Smith et al. 1997). We performed a similar analysis of theoutﬂow mass within the slightly more evolved region NGC
1333 (Plunkett et al. 2013), assuming Tex  17 K, which is themean Tex in Serpens South, and the blue- and redshifted mass
spectra for that region can be ﬁt with power laws with slopes ofH  5.1 o 0.2 and H  4.9 o 0.1 (see Figure 10, panels a andc). These two regions are therefore presented as examples forwhich the shallower (steeper) slope pertains to the less (more)
M (x , y , v)  2.72m H NH2 (x , y , v) A ,(4)evolved region.
It has been previously suggested that a break in γ at _10
where mH is the mass of a hydrogen atom, such that 2.72mH is
km s−1can be interpreted as two distinct outﬂow velocity
the mean molecular weight of the gas (Bourke et al. 1997), andcomponents (e.g., Richer et al. 2000; Su et al. 2004). In that
A is the area of a pixel. We sum over the cube in spatial andcase, the high-velocity gas corresponds to a recently accelerspectral dimensions
includingthat
havesignalgreater
ページ 4/12
Figure 10. Outﬂow
mass spectrapixels
for Serpens
South(circles)and NGC
1333 (triangles),atedwith blueand redshiftedandvelocity
channels slopeplotted independently
with theircomponent
steeper(γ ∼ 3–4),
and the low(a) shows
both regions
for comparison,aqua and magenta lines (solid for Serpens South, dashed for NGC 1333) are the best-ﬁt power
than 3σ torespective
get thecolors.
totalPaneloutﬂowmass,
momentum,andandenergy:
velocity gas is a slower, coasting component with a less steep
laws to the blue- and redshifted mass spectra, respectively. Panels (b) and (c) show Serpens South and NGC 1333 separately, where mass spectra with no velocitynor NGC 1333 showed aslope(H for_ Serpens2). Neither
symbols
South andSerpens
NGC 1333,South
respectively.
dependent correction for opacity (i.e., assuming U  1) are also given by plus and cross
Mout  4area 4vel M (x , y , v)
(5)
break in their power spectra according to our primary analysis.
−1
note that
our analysis included a velocity-dependent
Appendix
B,areaand4thehighest T v )values
found in redshifted∣ vWe
out ∣  3 km s , so as not to confuse emission from the
Pout  4v  vare(6)
vel M (x , y , excloudvelocity channels.
line, suggest,which especiallycorrection
opacityof the et12CO
ambient
cloud.for
However,as Dunhamal. (2014)
In addition to the mass calculation, we calculate outﬂow
typicalescape forvelocities
individual
sourcesare lessthanaccounted
more from“hidden”
mass
at low
velocities
than in1−1
2
and ,energy to (7)
be
to be
Pout
this,
and
therefore
ourmass calculation
potentially Whenomits some
Emomentum
4area
4vel
M=
(x ,13.6
y , vM)e vkm
 sv cloud
out 
the
case
of
the
optically
thin
assumption.
assuming that2 q 10 45 erg. Table 3 lists these diagnostics using
outﬂowing
an admittedly
Eout _ 1.0
lineMaking
is opticallythin (simple), breaks into the power
the 12COgas.
U  1extrapolationeach Tex method, and we consider these values to be lower
velocities as low as vout  0.3 km s−1, based on the spectral ﬁtspectra are apparent at outﬂow velocities of ∼5–10 km s−1 (seewhere v islimits
the velocity
corresponding
given channel,
because they
are uncorrectedtofora inclination
angle ofandtheshown in Figure 10, we ﬁnd that the total mass may increase by

LETTER

doi:10.1038/nature14166

The formation of a quadruple star system with wide
separation
42. “The formation of a quadruple star system with wide separation”!
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Jaime E. Pineda et Jaime
al.! E. Pineda1, Stella S. R. Offner2,3, Richard J. Parker4, Héctor G. Arce5, Alyssa A. Goodman6, Paola Caselli7, Gary A. Fuller8,
Tyler L. Bourke9,10 & Stuartt A. Corder11,12
http://ads.nao.ac.jp/abs/2015Natur.518..213P!
Nature Letter!

Detailed knowledge of the underlying distribution of dense gas is the
The initial multiplicity of stellar systems is highly uncertain. A number of mechanisms have been proposed to explain the origin of binary key to determining which structures will go on to form stars. Here weand multiple star systems, including core fragmentation, disk frag- identify the dense gas structures that are most likely to form stars using
mentation and stellar capture1–3. Observations show that protostellar the dendrogram technique21. Dendrogram analysis is a hierarchical strucand pre-main-sequence multiplicity is higher than the multiplicity ture decomposition that uses isocontours to identify individual features,
found in field stars4–7, which suggests that dynamical interactions while also determining where these contours merge with adjacent strucoccur early, splitting up multiple systems and modifying the initial tures to create a new parental structure. We refer to the smallest scale (andstellar separations8,9. Without direct, high-resolution observations brightest) structures in the dendrogram as condensations. These areof forming systems, however, it is difficult to determine the true the most likely places for an individual star to form. Figure 1a shows theinitial multiplicity and the dominant binary formation mechanism. B5 region as seen in dense gas (number density of H2, nH2 >104 cm-3),
Here we report observations of a wide-separation (greater than with the protostar and the identified gas condensations shown by a star
1,000 astronomical units) quadruple system composed of a young and circles, respectively. The mass of the well known protostar B5-IRS1
protostar and three gravitationally bound dense gas condensations. is 0.1 solar masses (MSun; ref. 22), while the masses of condensations
These condensations are the result of fragmentation of dense gas B5-Cond1, B5-Cond2 and B5-Cond3 are 0.36 6 0.09 MSun, 0.26 6 0.12
filaments, and each condensation is expected to form a star on a time- MSun and 0.30 6 0.13 MSun, respectively. Uncertainty in these masses isscale of 40,000 years. We determine that the closest pair will form a dominated by the uncertainty in the temperature used to convert meabound binary, while the quadruple stellar system itself is bound but sured fluxes to masses. The radii of the three condensations are respecunstable on timescales of 500,000 years (comparable to the lifetime of tively 2,800 AU, 2,300 AU and 2,500 AU, while the projected separations
the embedded protostellar phase10). These observations suggest that between the same three condensations and the protostar are 3,300 AU,
filament fragmentation on length scales of about 5,000 astronomical 5,100 AU and 11,400 AU (see Methods). The half-mass radii of the conunits offers a viable pathway to the formation of multiple systems.
densations are about half the condensation radii. This, combined with
Barnard 5 (B5) is a dense core in the Perseus star-forming region (at the mass radius relations (Extended Data Fig. 2), suggests that the central
a distance of 250 pc) that hosts at least one young, forming star11. Imaging regions will collapse faster than the whole condensations and before連星形成。!
of B5 in the emission of the dense-gas-tracing NH3(1,1) line shows it to be interactions between condensations can play a major role in the stars’連星形成として、Core
fragmentation,
stellarhigh-density
captureを含む数多くの提案がinitial separations. Although these separations are large, they are conan example
of a ‘coherent’disk
densefragmentation,
core12, which is a contiguous
region with subsonic levels of turbulence13. Higher-resolution
imaging sistent with initial protostellar pair separations predicted for core fragされて来た。観測的には、原始星・前主系列の連星はfieldstarに比べて高いことが示されている。
14
reveals narrow filamentary structure within the coherent core . We mentation by numerical simulations2. In the simulations, protostellar
observed the NH3(1,1) and (2,2) lines using the Karl G. Jansky Very separations evolve rapidly on timescales of 0.1 Myr, and some systems
これは進化の早い段階で力学的相互作用がおきて連星系を分離し初期の光線間距離が変化するこ
Large Array (VLA)15, which reveals that the filaments in B5 are frag- become unbound while others migrate to closer proximity.
とを示唆する。!
menting and that they are in the process of forming a wide-separation
Projected proximity on the sky does not necessarily imply that objects
multiple stellar system.
are physically related. However, the line-of-sight velocities of the observed
wide-separation (>1000au)の四重連星系(重力的に束縛された高密度ガス３つと原始星から成4,16Nearly half of all stars reside in multiple star systems . Consequently, condensations are similar and the grouping is likely to be physically assoa host of phenomena, ranging from supernova rates to はフィラメントの分裂の結果であ
planet formation, ciated. The velocity dispersion, sv, of the dense gas provides another imる)の観測についての報告。３つの高密度ガスのcondensation
depend on understanding stellar multiplicity17. Because of the observa- portant piece of information, the gas kinetic energy, which is needed to
り、それぞれ4e4 yrで星を形成することが期待される。4つのうち内側の近い二つのペアは連星にtional challenges associated with observing early systems, the dominant determine whether the condensations are transient structures or gravideas for binary formation are based on simulations and analytic argu- itationally bound and likely to form a star. The velocities and velocity
成る一方、他の二つは重力束縛下にあるものの5e5yrで不安定。!
ments, which naturally require a varietyof assumptions4,18. To date, dispersions of the condensations are determined by fitting NH3(1,1)
observations have not captured the formation of a binary system at a and (2,2) line profiles14. The condensations and protostar display the
5000auぐらいでのフィラメント分裂は連星形成の現実的な形成経路である。!
stage where its origin is unambiguous, and prior observations of core same centroid velocity to within 0.2 km s21 and are therefore associatedlack the spatial and kinematic resolution to be used in with the same dense core. The level of turbulence in this region is so low,
観測；VLA, NH3(1,1)
NH3(2,2)!
RESEARCHsubstructure
LETTER
predicting whether observed structures would form protostars and/or sturb < 0.53–0.66 times the sound speed in the gas14, that gravity will overproduce a boundsystem19,20. The observed kinematics
and separation whelm the combined turbulent and thermal pressure in all the identifiedab56′(.1,000 astronomical
units, AU) of the B5 system is significant because condensations, and a star will probably form in each case. The timescale
Barnard 5
for these condensations to undergo gravitational collapse is approximately
it demonstrates a clear mechanism for widebinaryJVLA
NH3(1,1)formation and provides convincing evidence that the observed condensations will become the gas free-fall time, which we estimate to be 40,000 years (Methods).
55′
This timescale is sufficiently short to ensure that the system’s spatial
a bound multiple
star system.
1

Institute for Astronomy, ETH Zurich, Wolfgang-Pauli-Strasse 27, CH-8093 Zurich, Switzerland. 2Department of Astronomy, Yale University, PO Box 208101, New Haven, Connecticut 06520-8101, USA.
Department of Astronomy, University of Massachusetts, 710 North Pleasant Street, Amherst, Massachusetts 01003, USA. 4Astrophysics Research Institute, Liverpool John Moores University, 146
54′
Brownlow Hill, Liverpool
L3 5RF, UK. 5Department of Astronomy, Yale University, PO Box 208101, New Haven, Connecticut 06520-8101, USA. 6Harvard-Smithsonian Center for Astrophysics, 60 Garden
Street, Cambridge, Massachusetts 02138, USA. 7Max-Planck-Institut für extraterrestrische Physik (MPE), Gießenbachstrasse 1, D-85741 Garching, Germany. 8UK ARC Node, Jodrell Bank Centre for
Astrophysics, School of Physics and Astronomy, Alan Turing Building, Oxford Road, University of Manchester, Manchester M13 9PL, UK. 9SKA Organisation, Jodrell Bank Observatory, Lower Withington,
Macclesfield, Cheshire SK11 9DL, UK. 10Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, Massachusetts 02138, USA. 11Joint ALMA Observatory, Alonso de Cordova 3107,
53′
Vitacura, Santiago, Chile. 12National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia 22903, USA.

3

Declination (J2000)

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B5-Cond1

B5-Cond2

1 2 F E B R U A RY 2 0 1 5 | VO L 5 1 8 | N AT U R E | 2 1 3©2015 Macmillan Publishers Limited. All rights reserved

52′
B5-IRS151′

B5-Cond3+32º 50′

10,000 AU
44 s

ページ 5/12

40 s

36 s

10,000 AU

3 h 47 min 32 s

Right ascension (J2000)

Figure 1 | High-angular-resolution image of dense gas and stellar
progenitors. a, Background image from the JVLA of the Barnard 5 region shows
the dense gas traced by NH3(1,1). It reveals two filaments, which together hostthree gravitationally bound condensations (B5-Cond1, 2, 3). Red contours andorange-filled circles show the condensation boundaries and centres, while the star

indicates the protostar (B5-IRS1) location. b, Contour map showing the filaments
in dense gas. Greyscale circles show the distance that could be covered during
40,000 years while moving at the speed of sound, 0.2 km s21. In a and b, filled red
circle at bottom left shows the angular resolution of the observations, and scale bar
is shown at bottom right.

configuration will remain nearly unchanged during collapse, even if the
protostars move as far as possible (in straight lines at the gas sound speed)for the duration of the collapse (,1,700 AU or 6.80 at the distance of Perseus). Figure 1b shows, as dotted circles, the possible ranges over whichhave made the most pessimistic assumptions possible (giving the lowest
stellar masses) in evaluating boundedness.
Depending on the final stellar masses and kinematics, the resulting
multiple stellar system could either be strongly bound or quickly dis-

ページ 6/12

Triangles – minimum during 2010. Filled circles – bright state
between the two minima. Grey dots – observations from 1986–
2005. The colour index is in the Johnson? system.

Another deep dimming of the classical T Tauri star RW Aur A(Research Note)
The [O i] 6300 Å line is present in the high-resolution spectra

obtained at CrAO and in the low resolution ALFOSC spectrum

1
1
P. P. Petrov1 , G. F. Gahm2 , A. A. Djupvik3 , E. V. Babina1 , S. A. Artemenko
N. Grankin(Fig. 4). ,Inand
ourK.
observations
from 1996, and in earlier publishedobservations (Hamann 1994; Hartigan et al. 1995) the [O i] line

Crimean Astrophysical Observatory, p/o Nauchny, 298409 Republic of Crimea
hasA”!
three emission components: a central one and two at radial
41.
“Another deep dimming of the classical T Tauri star RW Aur
email: petrov@crao.crimea.ua; petrogen@rambler.ru
velocities of about -140 and +140 km s 1 , which form in thePeter P. Petrov et al.!
two flows of a collimated wind (jet) at large distances from the2
Observatory, AlbaNova University Centre, Stockholm University, SE-106 91 star.
Stockholm,
Sweden! Stockholm
http://ads.nao.ac.jp/abs/2015arXiv150304158P!
In a seriesof spectra from 1999 presented in Alencar etal. (2005) the two jet components are at -170 and +120 km s 1 .
!3 Nordic A&A!
Optical Telescope, Rambla José Ana Fernández Pérez 7, ES-38711 Breña Baja, SpainarXiv:1503.04158v1 [astro-ph.SR] 13 Mar 2015

1

In the 2014 spectra the two components appear at much lower
velocities, namely about -90 and +90 km s 1 .
In addition, the equivalent widths of the forbidden lines of
ABSTRACT[O i] and [S ii] have changed drastically and are several times
Context. RW Aur A is a classical T Tauri star (CTTS) with an unusually rich emission line
spectrum.
2014 the star
faded
by ⇠3
stronger as In
comparedto thebrightstate. Considering that alsomagnitudes in the V band and went into a long-lasting minimum. In 2010 the star su↵ered thefromintensity
a similarof
fading,althoughless deep. is lower, we conclude thatthe stellar
continuum
These events in RW Aur A are very unusual among the CTTS, and have been attributed to occultations
bypassing
dust
clouds.
the flux in the lines did not change equally much. This implies
Aims. We want to find out if any spectral changes took place after the last fading of RW Aur A with the intention to gather morethat the occulting body covers the star and its immediate surinformation on the occulting body and the cause of the phenomenon.
rounding
but does
not extend further out to cover also the exMethods. We collected spectra of the two components of RW Aur. Photometry was made before
and during
the minimum.
tendedTheasymmetryintensity of the blue- and redResults. The overall spectral signatures reflecting emission from accretion flows from disk
to starwind.
did notchangeafter theinfading.
shiftedemissioncomponents,
seen
However, blue-shifted absorption components related to the stellar wind had increased in strength in certain resonance lines, andthein the CrAO spectrum of
profiles and strengths, but not fluxes, of forbidden lines had become drastically di↵erent. 2014, was present also in 1999 (Alencar et al. (2005).
Conclusions. The extinction through the obscuring cloud is grey indicating the presence of large
dust grains.
At theFinally,
we didnotsamefindtime,
any there
trace of additional absorption
are no traces of related absorbing gas. The cloud occults the star and the interior part of thecomponents
stellar wind, inbutthenot resonancethe wind/jet Na
furtheri and K i lines, which would
out. The dimming in 2014 was not accompanied by changes in the accretion flows at the indicate
stellar surface.
There is increase
evidence ofthatcolumn
thea substantial
density of cold gas tostructure and velocity pattern of the stellar wind did change significantly. The dimmings could
relatedto passinginwardbethedimmedstar.condensations
With the resolutionof ALFOSC an aba tidally disrupted disk, as proposed earlier, but we also speculate that large dust grains have been stirred up from the inclined disk
sorption line with EW > 0.2 Å can be detected. Such compointo the line-of-sight through the interaction with an enhanced wind.
nents can hide in the complicated Na i profiles, but the K i line
Key words. stars: pre-main sequence – stars: variables: T Tau – stars: circumstellar matteris– stars:
individual:
RW Aur in the ALFOSC spectrum the K i abnot blended.
However,sorption with EW = 0.4 Å is at -75 km s 1 and must be related tothe outflow. In the bright state the absorption of K i is at stellar
背景；RW Aur A; CTTS天体。 !
showing
moderate
1. Introduction
ing the beginning ofvelocity,
the minimum
indicating
thatnight-to-night
no significantvariability (Fig. 5).
2014に長い減光、V bandで３等。!
changes in the emission line spectrum had occurred.
The classical
T
Tauri
star
(CTTS)RW
Aur
A
standsout
among
2010にも同様の減光が起きているが、その度合いは小さい。! Normally, the star fluctuates in brightness by sometimes
the low-mass pre-main-sequence (PMS) stars with its unusually more than one magnitude
4. Discussion
in the V band and on time scales of
これらの現象はCTTS天体では非常に珍しく、dustrich emission
line spectrum as noted early by Joy (1945).cloudsに寄って掩蔽が起きている。!
The a few days. No distinct period has been found as summarised
From our study of the spectral state of RW Aur A before andstar hasbeen subject to a large number of investigations. Several in e.g. Gahm et al. (1993)
目的；二回の掩蔽現象の間でスペクトルに変かがあれば、現象の原因と掩蔽しているものの情報を調べたい。!
andR2013.we
However,colour thatchangesthe fading
can conclude
there are no changes inspectralfeatures respond to changes in the accretion rate, like with periods of 2.7 after
to 2.8days have emission
been reported(Petrov fromet He I, H and once
結果；円盤から星への高長句流からの反射光に依る全体のスペクトルの特徴には変化がない。一方で、星風と関係する
thecharacteristicline spectrum
narrow components in emission lines of e.g. He i, excess contin- al. 2001, R2013), which
is
close
to
half
the
expectedrotational
and
twice
ionizedmetals.
Therefore,
there
uous 青方偏位した吸収線は共鳴線(resonanceemission (called veiling), and red-shiftedlines)で吸収が強くなった、また禁制線のプロファイルと強度(fluxではない)はabsorption com- period of about 5.6 days as determined from variations of the are no signs that theflows close to the star has changed. On the other hand,ponentsflanking certain strong emission lines. These signatures longitudal magnetic accretion
劇的に変化した。!
fieldof the(Dodinal. 2012).
we
foundthatstarsomedrasticetchanges
hadTheoccurred in the strength
are prominent in RW Aur A and indicate that the accretion from star becomes redder withdecreasing
brightness,
and[S
the
結論；雲で隠されたときの減光はgrey (grey body)で大きなダストが存在する。同時に、関連したガスの吸収線が見られ
of forbidden linesof [O i] andii],frequentand in radial velocity of thea circumstellar disk is heavy and variable (e.g., Petrov et al. drops in brightness fromanaverage
level
of
V
⇡
10.4
appears
two components of [O i] 6300 formed in the bipolar jets of RW
2001;ない。雲は星と星風の内側を隠すがwind/jetの外側は隠されない。2014年の減光は星表面での降着流の変化によって起
hereafter called P2001). RW Aur is a visual binary with to be, at least in part,
relatedvariable
Aur
A. In toaddition,anforegroundincrease in extinction
the strength of wind signatures
a separationof 1.400 , where the faint component B is a weak-line (e.g., Herbst et al. 1994).
きた訳ではない。星風の構造と速度パターンは大きくは変化していない。減光は、潮汐破壊された円盤の固まりが通過
occurred in the resonance lines of Ca ii and Na i.
TTS, a PMS star with little or no evidence of accretion.
In
2014,
RW
Aur
A
enteredasecondlong-lasting
することに関連するかもしれない(初期に提唱されている)が、我々は大きなダストが強い星風と相互作用して視線方向
The increased intensity of the minimum
forbidden lines in RW Aur A,
In 2010, RW Aur A went through a deep minimum reaching in brightness, and this timeevenindeeper
thanthein stellar
2010. According
measuredunits ofcontinuum intensity, indicatesに傾斜した円盤から来たと推考する。!
P.
P.
Petrovetal.:Anotherdeep
dimmingoftheclassical
T
Tauri
star
RW
Aur
A (RN)
an amplitude of ⇠2 magnitudes in V and which lasted for 180 to the data collected by theAmerican Association of Variable
days (Rodriguez et al. 2013; hereafter called R2013). They at- Star Observers (AAVSO) the brightness dropped by more than 2
tributed
this drop to an occultation by part of a tidally disrupted magnitudes between May 1 and October 23. Resolved U BVRI
2010年の減光は180日間。とても長い。!
disk as evidenced in Cabrit et al. (2006) and noted that such a photometry of both components of RW Aur was performed on
望遠鏡；Nordic
Telescope(NOT)!
long-lasting event hadOpticalnever been
observed before in at least 50 November 13/14 2014 by Antipin et al. (2015). They found thatP. P. Petrov et al.: Another deep dimming of the classical T Tauri star RW Aur A (RN)
years. Chou et al. (2013)
collected high-resolution spectra dur- RW Aur A was ⇠3 magnitudes fainter in all bands compared tonormal levels and that the star was fainter than component B on
this date. Furthermore, they concluded that the drop in brightSend o↵print requests to: P. P. Petrov?
Based on observations collected at the Nordic Optical Telescope, ness was caused by an increase in foreground, mainly grey exLa Palma, Spain; Fast-Track Service program 50-409.
tinction.

!

Fig. 1. Left panel: light curve of RW Aur with the two minima during 2010 and 2014. The vertical bar indicates the time ofour spectroscopic observations. Right panel: colour-magnitudeFig. 2. Three spectra obtained during the bright state in 1998
diagram
RWpanel:Aur.lightOpen
circlesminimumduringFig. 1.ofLeft
curveof RW–Aurwith the two
min-2014.
covering the range of typical variations in the Ca ii K line (left)
ima during
2010 and 2014.
The2010.
verticalbar indicatesof state
Triangles– minimum
duringFilledcirclesthe– timebrightand one corresponding spectrum in 2014 (right). Note that in
our spectroscopic
observations.
Right –panel:
colour-magnitude
betweenthe two minima.
Grey dots
observations
from 1986–
Three thespectraobtainedbrightstate in 1998
profileis ofduring
P Cygthetypeindicating
an increased winddiagramof RWindex
Aur. Opencircles– minimum
during 2014. Fig. 2. 20142005.
The colour
is in theJohnsonsystem.
covering
the rangeof typicalvariations
the Caandii K inlinethe(left)
Triangles – minimum during 2010. Filled circles – bright statefeaturecentredat -130km s 1in. Here,following figurescorresponding spectrum in 2014 (right). Note that in
1
between the two minima. Grey dots – observations from 1986– and one
theprofile
wavelengthscaletypeis heliocentricand velocities
2014 theis of P Cyg
indicating an increased
wind are in km s
2005. The colour index is in the Johnson system.
1star.
withrespect
tothefeature centred at -130 km s . Here, and in the following figures
The [O i] 6300 Å line is present in the high-resolution spectra
the wavelength scale is heliocentric and velocities are in km s 1
obtained at CrAO and in the low resolution ALFOSC spectrumwith respect to the star.
i] 6300Å line is present
the high-resolution(Fig. 4).TheIn[Oourobservations
from in1996,
and in earlierspectra
publishedobtained at(Hamann
CrAO and 1994;
in the low
resolution
ALFOSCobservations
Hartigan
et al.
1995) spectrum
the [O i] line
(Fig. 4). In our observations from 1996, and in earlier publishedhas observations
three emission
components: a central one andtwo at radial(Hamann 1994; Hartigan et al. 1995)
the [O i] line
velocities
ofemission
about -140 and +140
km ones 1 ,andwhich
has threea central
two atform
radialin the
ページ 7components:/121
two velocities
flows of ofa collimatedwind(jet)kmats large
distances
about -140 and+140, whichform in fromthe thestar.two
In flows
a series
spectrawind
from
1999
presented
Alencaretof a ofcollimated(jet)at large
distancesinfrom
thestar. In thea seriesfrom 1999
in Alencar
et s 1 .
al. (2005)
two ofjetspectra
components
are presented
at -170 and+120 km1
al.
(2005)
thetwo
jetcomponents
are
at
-170
and
+120
km
s
.
In the 2014 spectra the two components appear at much lower
In the 2014 spectra the two components appear
velocities,
namely about -90 and +90 km s 1 . at much lower
velocities, namely about -90 and +90 km s 1 .
In addition,
widthsofof
forbidden
In addition,thetheequivalentequivalent widths
thetheforbiddenlines lines
of of[O i][Oand[S [Sii]ii]have
drasticallyandandseverali] andhavechanged
changed drasticallyare areseveraltimestimes

!

!

1

and one corresponding spectrum in 2014 (right). Note that in2014 the profile is of P Cyg type indicating an increased windfeature centred at -130 km s 1 . Here, and in the following figures
the wavelength scale is heliocentric and velocities are in km s 1with respect to the star.
P. P. Petrov et al.: Another deep dimming of the classical T Tauri star RW Aur A (RN)

P. P. Petrov et al.: Another deep dimming of the classical T Taura dimming. Photometric data indicate that the star has been occulted by a foreground cloud containing large dust grains. Our a dimming. Photo
primary concern has been to search for spectral changes that culted by a foreg
primary concern
could be related to the phenomenon.
We conclude that the unusually rich emission line spectrum could be related t
We conclude
of RW Aur A on the whole remained the same as at normalbrightness levels. For instance, spectral signatures of accretion of RW Aur A obrightnesslevelsflows in the immediate surrounding of the star have not changed,
and components in emission line profiles of H↵ that trace the in- flows in the imme
and components
ner part of a stellar wind/jet remained the same.
On the other hand, significant changes occurred in line com- ner part of a stell
Fig.
4.
Left
panel
–
ALFOSC
spectrafrom2014
coveringtheOn the other hponents related to more extended parts of the stellar wind, such
Fig. 3. Three spectra obtained during the bright state in
1 1996Fig.
4. Left panel– ALFOSC
spectrainfrom
2014 lines
covering
the ponents related to[O i] 6300
line.variations
The strongin lines
at i-90km s and(solid)as enhanced
blue-shifted
absorption
resonance
showing
typicalthe HeandandNa i+90
D lines,one
1from e.g.
[OCai] ii.6300
line.
The
strong
lines
at
-90
and
+90
km
s
(solid)
are
from
RW
Aur
A,
while
the
singleemissioncomponentatasenhanced
blue
Moreover, forbidden lines of e.g. [O i] and [S ii] became
corresponding spectrum in 2014 (right).
from RW Aur
A, while
thethesingle
component
zero velocity (dashed) is from RW Aur B. Right panel – high res- areconsiderablystrongerthan inbrightemission
state, andthe radial atve- Ca ii. Moreover,zero
velocity(dashed)
is
from
RW
Aur
B.
Right
panel
–
high
resolution spectra of the same region. Dashed curve shows the [O i] locities of the two components tracing more extended emission considerably stro
olutionspectraof the directedsame region.
curve shows[O i] locities of the tw
profilefrom bothflux
RWinAurandremained
B in 2014,
where
centralfrom theoppositely
flowsDashedin the wind/jet
were thedi↵erent.
that
thethe Alines
about
thethesame
af! observed
from both
RW Aur width
A and of
B thesein 2014,
theexplained
central from the opposite
peak
to RW Aur B, and the red-shifted sharp peak at +90 profile
The change
in equivalentlineswhere
can beterthebelongs
1star dropped in brightness. This implies that the occultpeak
belongs
Aur B, bodyand thered-shifted
sharp+90of The change in eq
kmbodys belongs
RW
Aur
blue-shiftedis shalas caused
bytoanRW
occultingthatcovers only
thepeak
inneratparting
coveredtothe
starandA.itsThe
inneraccretion emission
flow, wherethe1km
belongs
to andRW notAurtheA. remoteThe blue-shifted
emission is shal- as caused by an olow. Numerous
narrow
telluric form,
absorption
components
thes stellar
wind,parts.
broad
emission lines
of metals
and also
the innerand
partphooflow.
Numerous
narrow
telluric
absorption
components
andphotosphericlines fromcomponent
are present.
The solid
curve
The picture is not clear regarding the origin of the obscuring the stellar wind, athestellar windradiatingin H↵. B
However,the more
extended
tospheric lines from component B are present. The solid curvein theofrightthe were[O i] profile
in 1996,
with theThe picture is
cloud containing large dust grains, but one would expect thatparts
the panelwind shows(the jets)
not occulted
much.
Thiscentral
sets1
in the right panel shows the [O i] profile in 1996, with the central cloud containing
peak
at
zero
velocity
and
two
peaks
at
-140
and
+
140
km
s
.
these grains once resided in the disk plane. Since the disk ofcertain! restrictions to the size of the obscuring cloud.
peak at zero velocity and two peaks at -140 and + 140 km s 1 . these grains once
TheTheforbidden
lines
are
much
stronger
in
the
faint
than
in
the
RW Aur A is inclined by more than 30 degrees one must findphotometry shows that the extinction becomes very grey The
forbidden lines are much stronger in the faint than in thebrightthestate.
ways to explain how this material was transported far above the RW Aur A is inc
whenstar fades implying that large dust grains must be brightstate.
disk, into the line-of-sight. One explanation has been that the ways to explain hpresent in the obscuring body. The large extinction of Av ⇠ 3m
occulting body is part of a tidally disrupted disk. Our results do disk, into the lin
in combination with the absence of additional absorption lines
not rule out this hypothesis, but since we have found evidence occulting body isin resonance lines of Na i and K i indicate that the dust-to-gas
that the stellar wind has changed after the fading we speculate not rule out thisratio in the obscuring cloud could be very di↵erent from normalthat an interaction between this enhanced wind and the dusty that the stellar w
conditions in the interstellar medium.
disk caused an ejection of large dust grains into the wind. We that an interactio
The dimmings of RW Aur A are of a di↵erent nature than
did not find any traces of absorption by cool gas related to the disk caused an e
those observed in AA Tau and KH 15D (see Sect. 1). In the
did not find any
obscuring cloud.
hypothesis proposed in R2013 the occultation of RW Aur A in
obscuring cloud.
2010 was caused by a dusty structure connected to a tidally disAcknowledgements.
Basedonobservations
made
with
the
NordicOpticalrupted disk. Since the phenomenon has returned a second time
Telescope, operated by the Nordic Optical Telescope Scientific Association Acknowledgements.
one must assume that in this case the foreground structure con(NOTSA) at the Observatorio del Roque de los Muchachos, La Palma, Spain, of Telescope, operated
tains more than one condensation passing in front of the star.
Instituto de Astrofisica de Canarias. The data presented here were obtained (NOTSA) at the ObseFig.
The regionexplanation
covering theresonance
K i 7699theline
and the theAn5.alternative
to explore
is whether
dimmingwith ALFOSC, which is provided by the Instituto de Astrofisica de Andalucia the Instituto de Astro
5. underThe region
coveringwith
thetheresonance
KCopenhageni 7699 lineandandthe with ALFOSC, whici 7773 bytriplet.
– bright statein 1996,
right panel(IAA)a joint agreement
University of
NOTSA.
isOcaused
someLeftkindpanelof interaction
betweena changing
out-– Fig.
i 7773
triplet. funding
Left panel– bright
state in 1996,
panel– (IAA) under a joint a
faintand
state
2014.
Noteto theabsenceofoflarger
any enhanced
i ab- OKG
acknowledges
from theLabEx OSUG@2020
thatright
allowedpurchasflow
thein
disk
leadinga transport
dust grainsKfrom
ing theimagingsystemNote
installed
the 1.25-moftelescopeat CrAO. K
PPi thanks
state
in 2014.
theatabsence
any enhanced
ab- KG acknowledges fun
sorptioncomponents
in the spectrum
the dimmed
star;disk.
there faintthe
diskplane
into the line-of-sight
far ofabove
the1inclined
MarinaRomanova
forusefull
discussions.
is onlya relatively
absorption
at -75 km
s . was inves- sorption components in the spectrum of the dimmed star; there ing the imaging syst
The
problemof dustweak
survivalin disk windsof CTTS
Marina Romanova fo
1
tigated by Tambovtseva & Grinin (2001). The emergence of dust is only a relatively weak absorption at -75 km s .

!
!
!

!
!
!
!
!
!
!
!
!
!
!
!
!
!

into the wind can be caused either by disk erosion (e.g. Schnepf
et al.2014), or by the MHD processes at the disk-magnetosphere References
al.2014), or by the MHD processes at the disk-magnetosphere References
boundary (e.g., review by Romanova et al. 2013). In this case the etAlencar,
S. H. P., Basri, G., Hartmann, L., & Calvet, N. 2005, A&A, 440, 595
reviewby Romanova
2013).
case
the Alencar, S. H. P., Basstar and the H↵-emitting region are obscured by a screen related
Antipin, S.,(e.g.,
Belinski,A., Cherepashchuk,A.,et&al.
et al.
2015, In
IAUthisIBVS,6126
3 boundarystar
and theH↵-emitting
regionL.area screen
relatedBouvier,J., Grankin,K., Ellerbroek,E., obscured
Bouy, H. & by
Barrado,D. 2013,
A&A, Antipin, S., Belinski,
to a dusty wind containing large dust grains.
557, A77
Bouvier, J., Grankin,wind containing large dust grains.
The remarkable phenomenon of the long-lasting deep min- to a dusty
Pety, J., Presenti,
N, & Dougados,
C. 2006,
A&A, 452,deep897 min557, A77
TheS.,remarkablephenomenonof thelong-lasting
ima in RW Aur A rises a number of questions. How recurrent are Cabrit,
Chou. M.-Y., Takami, M., Manset, N. & et al. 2013, AJ, 145, 108
Cabrit, S., Pety, J., Pr
ima
in
RW
Aur
A
rises
a
numberofquestions.
How
recurrent
arethese occultations? Are there any similarities to UX Ori stars? Dodin A.V., Lamzin S.A., & Chuntonov G.A., 2012, Astronomy Letters, 38, 167 Chou. M.-Y., Takami
there any C.,
similarities
Ori stars?
Are the dimmings in fact related to a changing stellar wind? It these
Gahm,occultations?
G. F., Gullbring,Are
E., Fischerström,
Lindroos, K.toP.,UX& Lodén,
K. 1993, Dodin A.V., Lamzin S
dimmings
A&AS,
100, 371 in fact related to a changing stellar wind? It Gahm, G. F., Gullbrin
is premature to draw any deeper conclusions about the cause of Are the
Grankin,K.
N.,
Melnikov,S.
Yu.,
Bouvier,J.,
Herbst,W.,&
Shevchenko,
A&AS, 100, 371
ofS.
the phenomenon. Continued spectroscopic, photometric and po- is premature to draw any deeper conclusions about the cause V.
2007, A&A, 461,
183
phenomenon.
Continued
spectroscopic, photometric and po- Grankin, K. N., Melnlarimetric observations of the star will provide more insight, and theGrinin, V., Rostopchina, A. N., & Shakovskoi, D. N. 1998, Astron. Lett., 24, 802
2007, A&A, 461
the2002,
starAstron.
will provide
it will be interesting to see whether the normal wind features are larimetric
Grinin, V., observations& Tambovtseva, of
L. V.
Lett. 28, more
601 insight, and Grinin, V., Rostopchi
itGrinin,
will beV.,interestingwhethernormalwind489,features
Stempels, H.to
C.,seeGahm,
G. F., &theet al.
2008, A&A,
1233 are Grinin, V., & Tambov
restored when RW Aur A returns to its normal state.
Hamann,when
F. 1994,RW
ApJ,Aur93, 485
Grinin, V., Stempels,
restored
A returns to its normal state.

!

5. Conclusions
We have observed the binary T Tauri star RW Aur A and B spectroscopically and photometricly during the deep, long-lasting
minimum that occurred on the primary in 2014. This is the second time over at least 50 years that RW Aur A su↵ers from such
4

Hamilton, C. M., Johns-Krull, C. M., Mundt, R., Herbst, W., & Winn, J. N. 2012,
751:147
P., Edwards, S & Ghandour, L. 1995, ApJ, 452, 736
5.Hartigan,
Conclusions
Herbst, W., Herbst, D. K., Grossman, E. J., & Weinstein, D. 1994, AJ, 108, 1906
Joy,
A.H.
1945, ApJ, the102,binary168
We have observed
T Tauri star RW Aur A and B specMamajek, E. E., Quillen, A. C., Pecaut, M. J., & et al. 2012, AJ, 143, 72
troscopicallyandG.photometricly
the deep,long-lasting
Petrov, P. P., Gahm,
F., Gameiro, J. F.,during& et al. 2001,
A&A, 369,
993 (P2001)
minimumon theI. 2001,
primary
in375,
2014.
Petrov, P. P.,that
Pelt,occurredJ., & Tuominen,
A&A,
977 This is the sec-

ond time over at least 50 years that RW Aur A su↵ers from such4

ページ 8/12

Hamann, F. 1994, Ap
Hamilton, C. M., John
751:147
Hartigan, P., Edwards
Herbst, W., Herbst, D
Joy, A.H. 1945, ApJ,
Mamajek, E. E., Quil
Petrov, P. P., Gahm, G
Petrov, P. P., Pelt, J., &arXiv:1502.04467v1 [astro-ph.SR] 16 Feb 2015

Liu,6 Álvaro Sánchez-Monge2
ABSTRACTWe report on the Submillimeter Array (SMA) observations of molecular lines
at 270 GHz toward W3(OH) and W3(H2 O) complex. Although previous observations already resolved the W3(H2 O) into two or three sub-components, thephysical and chemical properties of the two sources are not well constrained. Our
SMA observations clearly resolved W3(OH) and W3(H2 O) continuum cores. Taking the advantage of the line fitting tool XCLASS, we identified and modeled arich molecular spectrum in this complex, including multiple CH3 CN and CH3 OH
transitions in both cores. HDO, C2 H5 CN, O13 CS, and vibrationally excited linesof HCN, CH3 CN, and CH3 OCHO were only detected in W3(H2 O). We calculate gas temperatures and column densities for both cores. The results showthat W3(H2 O) has higher gas temperatures, and larger column densities thanW3(OH) as previously observed, suggesting physical and chemical diﬀerences
between the two cores. We compare the molecular abundances in W3(H2 O) to
those in the Sgr B2(N) hot core, the Orion KL hot core and the Orion Compact
Ridge, and discuss the chemical origin of specific species. An east-west velocity
gradient is seen in W3(H2 O), and the extension is consistent with the bipolar–2–
outflow orientation traced by water masers and radio jets. A north-south velocity gradient across W3(OH) is also observed. However, with current observations
we can not assure if the velocity gradients are caused by rotation, outflow orradial velocity diﬀerences of the sub-components in W3(OH).

– 19 –

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Fig. 2.— A: Lower sideband spectra at W3(H2 O) and W3(OH) continuum peak positions.
The black curve is the observed spectral profile and red curve is from LTE modeling using
XCLASS software. Upper panel is for W3(OH) and lower panel is for W3(H2 O). The upper
level energies (Eu ) of CH3 OH and CH3 OCHO are labelled. CH3 OH (90,9 -81,7 ) and (171,17 162,14 ) at 267.4 GHz are blended together.

44. “SMA OBSERVATIONS OF W3(OH) COMPLEX: PHYSICAL AND CHEMICALSMA OBSERVATIONS OF W3(OH) COMPLEX: PHYSICAL
DIFFERENTIATION BETWEEN W3(H2O) AND W3(OH)”!
AND CHEMICAL DIFFERENTIATION BETWEEN W3(H2 O)
Sheng-Li Qin et al.!
AND W3(OH)!
http://ads.nao.ac.jp/abs/2015arXiv150204467Q!
!
ApJSheng-Li
accepted!
Qin,1 Peter Schilke,2 Jingwen Wu,3 Yuefang Wu,4 Tie Liu,5 Ying

– 25 –– 22 –

1
Department of Astronomy, Yunnan University, and Key Laboratory of Astroparticle Physics of Yunnan
Province, Kunming, 650091, China; slqin@bao.ac.cn
Subject headings: ISM:abundances — ISM: individual (W3(OH)) — ISM: molecules
SAM 270GHzで 2“W3(OH)”と”W3(H2O)”を観測して、物理的・化学的特徴を調べた。!
Physikalisches Institut, Universität zu Köln, Zülpicher Str. 77, D-50937 Köln, Germany— radiolines: ISM — star: formation
これら二つの天体はSMA
連続はで分解できた。!
3
Department of Physics and Astronomy, University of California, Los Angeles, CA 90095, USA
CH3CN,CH3OHの複数輝線を含む輝線に対して分子輝線を同定、モデルした。!
4
Department of Astronomy, Peking University, Beijing, 100871, China
HDO,C2H5CN,O13CSと振動励起のHCN,CH3CN,CH3OCHOはW3(H2O)のみで検出された。!
5
1. Institute
INTRODUCTION
Korea Astronomy and Space Science
776, Daedeokdaero, Yuseong-gu, Daejeon, Republic of
Gas温度、柱密度をそれぞれのコアで計算した結果、W3(H2O)の方がW3(OH)に対してガス温度が高く、柱密度も大きかった（以前かKorea 305-348
ら知られていた）。二つのコアで物理的・化学的に異なることを示唆する。!
6 How massive stars form is poorly understood yet, partly because powerful radiation
Department of Physics and Hebei Advanced Thin Film Laboratory, Hebei Normal University, Shiji!
W3(H2O)はHot
coreで進化段階として若い。!
pressure050024,
from
stars
azhuang
China with masses above ∼8 M⊙ should prevent the gas accretion on the protostars
to hotform
more
massivestars,Orionaccordingto Ridgeを比較し、分子種の化学的起源について議論する。!
the monolithic collapse model. Diﬀerent
W3(H2O)と、Sgr
B2(B)
core,
Orion
KL hot core,
Compact!

Fig. 5.— Intensity-weighted velocity maps of HCN v2 =1 at 267.199 GHz, and CH3 OH at– 26 –
266.838 GHz with FWHM beam size of ∼2′′ .7×2′′ .4, PA=–61◦ . The scales on the right show
the velocity range.

theoreticalscenarios
related to high-mass star formation have been proposed, i.e., monoCH3OH,
CH3OCHO,CH3CNはこれらの領域と似た特徴を持ち、化学的起源は同じ!

lithic collapse (e.g., Jijina & Adams 1996; McKee & Tan 2003), competitive accretion andcoalescence (e.g., Bonnell et al. 1997, 1998), and how each of them takes eﬀects depends on
W3(OH)では南北の速度勾配が観測された。しかし、その原因は判別できない。!
the initial environments of their parent clouds (Zinnecker & Yorke 2007, Tan et al. 2014).
Observations are essential to test scenarios of high-mass star formation. Lower spatial resolution observations can not resolve detailed structure due to their large distances, nor canthey explore the clustered environments, and small scale variations in high-mass star forming regions. Higher spatial resolution observations are necessary to characterize kinematics,
physical and chemical conditions of high-mass star forming regions at small spatial scales.

– 18 –
W3(H2O)では東西の速度勾配が見られた、また広がりはH2Oメーザーち電波ジェットでトレースされている双曲流の方向と一致。!

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The W3(OH) complex, located at 2.04 kpc (Hachisuka et al. 2006), is one of thenearest and well-studied high-mass star forming regions, and harbors two objects, W3(OH)and W3(H2 O). Radio observations suggested that W3(OH) is an UC Hii region, ionized byyoung OB stars, and is rich in OH masers (Reid et al. 1995; Wilner et al. 1999; Fish &Sjouwerman 2007). W3(H2 O), also known as W3(OH)-TW (Turner & Welch 1984), locates6′′ east of W3(OH) and is rich in H2 O maser and organic molecules (Wyrowski et al. 1999;
Chen et al. 2006; Zapata et al. 2011; Hernández-Hernández et al. 2014), and presents
hot core properties (Kurtz et al. 2000). Single dish observations have shown overall inflow
in the W3(OH) complex (Wu & Evans 2003), and outflows as well as a possible disk wereidentified by Zapata et al. (2011). Subarcsec resolution observations of continuum andCH3 CN lines showed a high-mass protobinary system in W3(H2 O), with the two sub-coresFig. 1.— Continuum image of W3(OH) complex in both contours and grey scale. The
ページ
/from
12
having(Chen et al. 2006). However, physical and chemical
contours9
are10 to 90% diﬀerentof the maximum physicalintensity (2.3 Jyproperties
beam ). The synthesized
beam of ∼2 .7×2 .4 (PA=–61 ) is shown in lower right corner.
properties of W3(H2 O) and W3(OH) are still not well characterized.
−1

′′

′′

◦

In this paper, we present results from Submillimeter Array (SMA1 ) observations towards
1

The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the!

Fig. 3.— Sample images from specific molecular species. The contours are from 10 to 90%of the maximum integrated intensity for each molecular transition. The peak values of HCN
v2 , C2 H5 CN, CH3 CN, CH3 OCHO, higher and lower energy transitions of CH3 OH are 19.6,
11.6, 41.7, 12.2, 11.5, 78.5 Jy beam−1 km s−1 , respectively. In each panel, the synthesized
beam is shown in lower-right corner, and the cross symbol indicates the peak position of thecontinuum source.

ページ 10/12

Fig. 6.— Abundances of molecules in W3(H2 O) relative to H2 (maroon bars), compared toabundances in the Sgr B2(N) hot core (green) from Neill et al. (2014), the Orion KL Hot
core (turquoise) and Compact Ridge (blue) from Crockett et al. (2014).

Á. Ribas et al.: Protoplanetary disk lifetimes vs. stellar mass and possible implications for giant planet populations

Astronomy&
Astrophysics

c ESO 2015

Protoplanetary disk lifetimes vs. stellar massand possible implications for giant planet populations?

1,2,3stellar
1,445. ”Protoplanetary diskÁlvaro
lifetimesvs.
possible
implications
for giant planet
Ribas, Hervémass
Bouy2 ,and
and Bruno
Merín
populations”!
1 Álvaro Ribas et al.!
European Space Astronomy Centre (ESA), PO Box, 78, 28691 Villanueva de la Cañada, Madrid, Spainaribas@cab.inta-csic.es!2 e-mail: http://ads.nao.ac.jp/abs/2015A%26A...576A..52R!
Centro de Astrobiología, INTA-CSIC, PO Box – Apdo. de correos 78, 28691 Villanueva de la Cañada, Madrid, Spain!3 Ingeniería
A&A!
y Servicios Aeroespaciales-ESAC, PO Box, 78, 28691 Villanueva de la Cañada, Madrid, Spain!4 Herschel Science Centre, ESAC-ESA, PO Box, 78, 28691 Villanueva de la Cañada, Madrid, Spain

Received 21 August 2014 / Accepted 19 January 2015

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Protoplanetary disk fraction (%)

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!
!

A&A 576, A52 (2015)
DOI: 10.1051/0004-6361/201424846

PDFs of the “evolved” disk fractions obtained in Sect. 4.3 forhigh-mass stars on the one hand and low-mass stars on the otherhand. The PDFs are completely compatible both for low-massand high-mass stars, pointing to a roughly constant level of
evolved disks within the uncertainties at any age regardless ofstellar mass.

Finally, the implications of the faster disk dispersal aroundhigh-mass stars for terrestrial planets are harder to determine. As

60-70%
5.1. Dependence of protoplanetary
disk lifetimes on stellar
massPrevious observational studies had already found lower protoplanetary disk fractions around high-mass stars and suggested
that their protoplanetary disks disperse earlier and/or evolve35-40%
faster (see e.g. Williams & Cieza
2011, for a review). A number of these studies focused on individual regions, making it
hard to confirm a global dependence of protoplanetary disk evolution on stellar mass independently of the environment andmethod used to identify the disks (e.g. Lada et al. 2006; Dahm
& Hillenbrand 2007; Currie et al. 2009; Carpenter et al. 2006, 5-20%
in
IC 348, NGC 2362 and Upper Sco). Kennedy & Kenyon (2009)
performed an original and coherent study of disk frequencies
in nine young associations by combining optical spectra and0%
IR photometry. In contrast to our study, they found only a weak
dependence of disk dispersal
on stellar
mass. The ages used in
1 Myr
3 Myr
11 Myrtheir study have been improved and sometimes changed significantly. For example, the age of Upper Sco was recently reestimated to 11 Myr Pecaut et al. (2012) from an original 5 Myr(Preibisch & Zinnecker 1999), resulting in a noisier distribution
Fig. 5. Evolution of protoplanetary disks as a function of mass, as deof disk fraction as a function of age in their study.
rived
in this quantitatively
study.
Our results
demonstratethat protoplanetarydisks dissipate significantly faster and/or earlier around highmass stars, as we find between 10 and 30% more protoplanetarydisks around low-mass stars regardless of the age.
Our data additionally show that the disk fraction for highmass stars falls⇡0% Myrs,
anytime fastening
between 3 andthe
11 Myr,andthetofirstprocesses
responsible of disk dispersuggest that protoplanetary disks disperse up to two times fastersal (Alexanderal. 2014).
around high-massstars than aroundet
low-mass
ones. More datafor regions with intermediate ages are needed to better constrain
this ratio.
Interestingly, other studies found a higher (80 85%) pro5.2.fraction
A constantlevel oftoplanetary disk
and no dependence
withevolved
stellar massdisks
in the younger Trapezium cluster (<1 Myr, Lada et al. 2000)
and NGC 2024
⇥ 105 yr, Haischet al. 2000).
Weclusterfind(<0.4
the 1.4
fractionof evolved
sources to remain constant with
Combined with the present study, these results suggest that

Age (Myr)This e↵ect is probably not critical for our analysis of protoplanetary sources (Sect. 4.1). These dense and massive disks
have complex physical structures, and the IR excess at a cer-under networks of the SCIENCE and Human Capital and Mobility programs,

Aims. We study the dependence of protoplanetary disk evolution onthestellar
mass usinga largeof young stellarinLandessternwarte,
Heidelberg
and sample
Institut d’Astrophysiquede Paris.objects
This

Low-mass stars
High-mass stars

5. Impact of stellar mass on disk evolution

Fig. 4. Comparison of the obtained PDFs for evolved sources as a function of age and for the two stellar mass bins.

A&A
576, A52 (2015)
ABSTRACT80%

A&A 576, A52 (2015)

tain wavelength is emitted from relatively large areas of the disk.
makes use of data products from the Wide-field Infrared Survey
nearby youngstar-forming
regions.
previously
mentioned,
their formation follows the initial stage of publication
Moreover, the protoplanetary definition is based on the presence
Explorer,
whichwork
is a joint
project of
the this
Universityof California,
Methods.theWe
protoplanetarydisktofractions
presented
our recent(PaperI ofseries)derivedLosfor
22 nearby
Á.Angeles,
Ribas
et al.: Protoplanetary disk lifetimes vs. stellar mass and possible
implications
for giantplanet populations
CAupdatemodel,theaccretingplanetesimals
form a planetary
em- in andof excessover a relatively
wide wavelength
range, ensuring that
the Jet Propulsion Laboratory/California Institute of Technology, funded(<500 pc)bryo
associations
between
1 and
100which
Myr. will
We notuse becomea subsampleof National
1428 spectroscopicallyconfirmed members
thethe corresponding emission comes from broad areas of the disk
and eventuallya rockyplanet,a by theAeronautics and Space Administration.
This research tousedstudy
theregardless of the stellar mass.
the Canadian
AstronomybytheNational
age both for low- and high-mass stars: 5 15% of low-mass stars
giantmass
if theonplanetcannot accreteenoughgas forWe
anydividereason.thisfacilities
impact ofgas
stellarprotoplanetarydiskevolution.
sampleof into
two stellar
massData
binsCentre(2orbiting
Moperatedboundary)
and
two
age
other stars. The global planet formation process may be
On the other hand, our definition of evolved disks relies priResearchCouncilof
CanadawiththesupportoftheCanadianSpace
Agency.
– the disk fractions before 1 Myr are in the continuity of the
However,
this process
is fasterthanexcesses
gas giant overformation:
earthbins (3 Myrboundary),
and useinfrared
the photosphericemission to classify objects in threegroups: protoplanetary
display
this kind of excess, and the fraction increases to 20 30%
a combination
marily on IR excess at 24 µm (MIPS1). The area of the disk
results presented
here;This publication makes use of data products fromthe Two Micronof
Allthese
Sky two scenarios.
massplanets
canbeformedin
1
Myr(Raymondetal.
2014),
andprobedat
24µm
is
expected
tovary
significantly
betweenmasdisks, evolved disks, and diskless. The homogeneous analysis and biasSurvey,
corrections
a statisticallysignificant– the dependenceof disk fractionsobjects
on stellar mass
which isallowa joint for
projectof the University
ofIfMassachusetts
theCA inter-comparisonis theanddominantmechanism, then planets are likely
for high-mass(seeappears
Fig. af4). Again, the dependence withsive and low-mass stars, and might make the comparison of theso the disk
dissipation dependence with stellar mass may not be Infrared Processing and Analysis Center/California Institute of Technology,
ter the first Myr of the star’s lifetime, and persists until noof the obtained
results.
to
form
below
the
ice-line(Ida&
Lin
2004a):
watercanfreezestellarofaroundthe⇠10disk
excess fractions and their interpretation more difficult.
protoplanetarydisksmass
are found
Myr. radii probed at a certain wavelength
as relevant as for giant planet formation. Moreover, terrestrial funded by the National Aeronautics and Space Administration and the National
Results. We find robust statistical evidence of disk evolution dependencewith stellar
mass.
Ourinresults,
combinedwithprevious
To test how much this dependence might a↵ect our results,
into
ice
aroundthegrainsofdust,andincreasetheefficiency
of
ScienceFoundation.
This
work
is
based
part
on
observations
made
with
the
shouldbeconsiderandit
could,
in fact, account for the higher
planets do not undergo type-II migration, and therefore their orwe computed a grid of radiative transfer models of circumstel- Based on these two suggestions, we propose a scheme of disk
studies onbits
disk
confirm
thatdisk
protoplanetarydisks evolve
and/orearlier around(>2) stars.Laboratory,
Wetoalso
find
SpitzerSpace Telescope,
which is high-mass
operated by the
Jet M
Propulsion
theiraggregationform
the required planetary embryo. On the
areevolution,less dependent
on theproperties. However,their fasterof evolved
disks
aroundlar disks using the code MCFOST (Pinte et al. 2006). We found evolution as afraction
function of stellar
mass (see Fig.
5). Thetimescalehigh-mass stars. Nevertheless,
California
Technology under a contract with NASA. This researcha roughlyorbits
constant
levelof evolved
throughout
thegiant
whole
age andstellarInstitutemass ofspectra.
other
hand,
formation
at
very
large
radii
from
the
host
star
is
that
synthetic
spectralenergydistributions(SEDs)
classified
asof
protoplanetary
disk
formation
andevolutionismostlyindecan bestronglya↵ecteddisks
by themigration ofplanets
has made use of the Exoplanet Orbit Database and the Exoplanet Data Explorer
the obtained values at these ages can be established to be within
evolved according to our criterion always corresponded to higher pendent of stellar mass until the first Myr. At 3 Myr, significant
Conclusions.
diskevolution
depends
on stellar mass. Such a dependence
could haveimportantin the We
sameconclude
system. Ifthat
theprotoplanetary
migration historyof gasgiants has
also unlikely,
since
the solid accretion rate of planetary embryos
at exoplanets.org.
5 already30%.observable between the evolution of prolevels of disk processing regardless of the star’s temperature. di↵erences are
implications
for gas giant
planetmass,
formation
migration,
anddifcould contribute to explaining the apparent
paucityof hotJupiters to the star. As an example, the opa dependencewith stellar
it couldand
indirectly
producedecreaseswiththe distance
These simple simulations suggest that our classification scheme toplanetary disks around low- and high-mass stars. By ⇠10 Myr,
evolveddisk
definition
ferent orbits
for terrestrial planets too. The identification of any
around high-massstars.
is efficient and robust enough to allow a comparison of the re- no protoplanetary The
disks arefound aroundhigh-mass
stars, butproposed here includes transitimal zone to form planets for a solar-type star is estimated tosults obtained for high and low-mass stars.
⇡10 15% oftional
low-massdisks,stars stillhot
harbora disk. Stronger
possible trend of this e↵ect requires further modeling and study Referencesdebrisdisks,radiaedge-on disks and circumbinary
betweensequence
5 10 AU (Helled et al. 2014). Surprisingly, a sigFinally, a comparison of the evolution of evolved disks can tion fields and higher accretion rates of high-mass stars (Calvet
Key words.
disks– planet-disk interactions – stars: formation– planetary systems – stars:bepre-mainof theprotoplanetary
known planetary
population.
disks.
Unfortunately,
theclassification
does not allow us to disAlexander, R. D., & Armitage, P. J. 2009, ApJ, 704,989
nificant number of Jovian-like planets are found very close tobe performed within the same mass bin, removing most of the et al. 2005; Garcia Lopez et al. 2006; Hillenbrand 2008) are
Alexander, R., Pascucci, I., Andrews, S., Armitage, P., & Cieza, L. 2014, in
previous dependency on stellar temperature. Figure 4 shows the likely to a↵ect
the evolution of
protoplanetarydisks evencriminate
betweenthese
fourduringtypes of disks. We neverthelessstars, andare ofreferred to as hot Jupiters. This particularProtostars and Planets VI, eds. H. Beuther, their
et al. (Tucson:
University
6. Conclusion
note that simple geometricconsiderations
suggest that edge-on
A52, page
7 of 13subgroup of exoplanets suggests that a fraction of these gas giArizona Press), 475
Allard, F., Homeier, D., & Freytag, B. 2012, Roy.
Soc.migrate
London Philos.
Trans.after their formation: once a gap is opened
ants
inwards
disks must be relatively rare contribute only moderately to theWe have studied the disk lifetime dependence on stellar mass usSer.
A, 370,2765
1. Introduction
The
evolutionof protoplanetaryand planets
are closely
by disks
the 655,
planet,population of evolved disks presented here. Hot debris disks are
ing the sample of YSOs compiled by Ribas et al. (2014). After Allen, P. R., Luhman, K. L., Myers, P. C., et al. 2007,
ApJ,
1095 it gets locked to the disk and exchanges angularthe disk
definethe initial
Andrews, S. M.,as& Williams,
J. P.properties
2005, ApJ, 631,momentum1134
with it,conditions
causing theof
planet to move closer to its hosttheir results, we have divided the sample into low- or connected
expected to appear mostly at later stages of disk evolution andStudying the updating
evolutionof protoplanetary disks around young
Andrews,formation,
S. M., Wilner, D.andJ., Espaillat,
C., et al. evolution
2011, ApJ, 732,subsequently
the planeta↵ects et al. 2007). Other migration
high-mass stars (boundary set to 2 M ) and young or old re- planetstar(type-II42migration, Papaloizou
must populate mostly the older bin of our analysis. Circumbinary
Armitage, P. J., Livio, M., Lubow, S. H., & Pringle, J. E. 2002, MNRAS, 334,
stellar objectsgions(YSOs)
iscrucialtobettercomprehendhowplan(boundary set to 3 Myr). We then study the fraction of the 248
disk properties (e.g. Kley & processes
Nelson 2012;
et al.
such asBaruteau
planet-planet
scattering, Kozai resonances, ordisks are expected to be common given the multiplicity freets form. Understandingof fundamental
protoplanetarytheand e↵ects
evolved disks
as a function ofparameters
age and mass. 2014).
Andrews,The
S. M., Rosenfeld,
K. A., Kraus,&similarWilner, D.dynamic
J. 2013,
ApJ,mechanisms
771, to underdisk lifetime
is A.
a L.,
fundamental
parameter(e.g. Kozai 1962; Wu & Murray Fig. 6. Orbital semi-major axes as a function of the mass of the host quency among stars (e.g. Bouy 2011), but only the closest sys129
The large
numberof sources inor
thebinarity
sample allows
us toinconsuch as age, stellar
mass,
environment,
providesstar for the known exoplanet population. Red circles show hot Jupiters
2003;
Naoz
et
al.
2011)
are
likelyresponsibleforthefraction
standtheinterplay
between
a
disk
and
its
planets,sinceitdetems would mimic the SED of evolved disks defined here (e.g.
the existenceof a processes,
dependence ofanddiskcould
lifetimeshedon thelightstellar Barrado y Navascués, D., Béjar, V. J. S., Mundt, R., et al. 2003, A&A, 404, 171
sight into the firm
planetformationof planets
hot Jupiters
whose
spin
is found to be misaligned with their (M > 1 M J , period <10 days), black circles represent the rest of theBarrado y Navascués,D., Stau↵er,
J. R., Morales-Calderón,M., etto
al. 2007,ApJ,
the time
available
for theform
andmigratemass: high-mass stars disperse their disks up to two times faster termines
664, 481
onto important
relations
betweenplanetarysystems
andtheir (e.g.
planet-starorbits (Triaud2010; Narita et al. 2010; Morton sample. The size of the symbols scales with the planet mass. We also as in Ireland & Kraus 2008; Nguyen et al. 2012). The evolved
Burkert&A.,
IdaPaardekooper,
2007; Alexander& Armitage2009).etAal.deand/or
earlier thanlow-mass
ones. We alsofind thatthe fraction
Baruteau,C., Crida,
S.-J., et al. 2014,
in Protostars andshow the estimated stellar mass radii of main-sequence stars (Siess et al. disks in the young (<3 Myr) bin are therefore most likely primarJohnson
2011).
stellar hosts. Several
on this topic(see circumbinary,e.g. Williams
of evolvedstudies
disks (includingtransitional,
and&hot pendence
Planets VI, ofeds. theH. Beuther,al. (Tucson: University
ofstellar
Arizona
Press),
disket lifetimeon& themass667could have2000) and the stellar radii of the host stars. A clear gap of hot Jupiters ily made of transitional and circumbinary disks, and our analysis
Bayo, A., Barrado, D., Stau↵er, J., et al. 2011, A&A, 536,
A63 this theoretical framework, the obtained depenWithindebris disks) etremainsroughlyEspaillat
constant (5et 30%)
during havethe first strong
Cieza 2011; Alexanderal. 2014;
al. 2014)
is seen for M⇤ > 1.5 M .
for planet
formationand migration
theories
Bell, C. P.implications
M., Naylor, T., Mayne,
N. J., Je↵ries,
R. D., & Littlefair,
S. P. 2013,
suggests that these two types of disks can be found even during
10
Myr
of
the
stellarlife.
Thefasterdispersalofprotoplanetarydenceofstellarmass and disk lifetime would have two mainalready drawn a comprehensive picture of protoplanetary disk (seeMNRAS,
434, 806e.g.
Raymond
et
al.
2014;
Helledetal.
2014;
Baruteauetal.
disks around high-mass stars could have important implications Birnstiel, T., Andrews, S. M., & Ercolano, B. 2012,
the very first Myrs of stellar life.
implications:
A&A,
544,
A79
evolution.

!
!
!
!
!
!
!
!
!
!
!
!
!

reviews
this topic),andresult infor giant planet formation and migration, and may contribute to 2014,
Bouy, H.for
2011,recent
in Computational
Star onFormation,
eds. J. Alves,
B. G.ultimately
Elmegreen,
J. M. Girart,& V. Trimble,
IAU Symp., 270,around411. Observational
apparent
hot Jupiters
aroundthesestars.hereafter di↵erent
In the firstthepaperof lack
thisofseries(Ribaset al.
2014,
results
tentatively suggest that massive stars
planetary
populations
low-mass andhigh-mass

Bressan, A., Marigo, P., Girardi, L., et al. 2012, MNRAS,
127 likely to host gas giants than their low-mass counare427,
moreR14) we presented
a large This
sampleofbeen⇠2300
studieshave
Burkert,Previous
A., & Ida, S. 2007,
ApJ, 660,
845 already found hints of protoplaneAcknowledgements.
work haspossible spectroscopicallythanks to the ESAC Science stars.
terpartsstars(Johnsonal. 2010). This trend would be the reOperations
Division
researchfunds Myr)with codeandSC 1300016149,
support from
Caballero,
J. A.evolving
2007, A&A, 466,
917 around massive
confirmed YSOs
in 22
young
(<100
nearby (<500pc)the tarydisksfaster(e.g. etCarpenter
ESAC Space Science Faculty and of the Herschel Science Centre. We thank Calvet, N., D’Alessio, P., Hartmann, L., et al. 2002, ApJ,
568,
1008
sult of Kennedy
two competing
factors: high-mass stars have morestar-forming regions
associations;
R14haveused
thislarge thesamal.! N.,
2006;
Dahm
& Hillenbrand2007;& Kenyon
D’Alessio,P., Watson,
D. M., et al. 2005, ApJ,
630, L185
the referee andfor valuablecomments which
helpedto improve
paper. etCalvet,
massivedisks(e.g.
Andrews
alsocontamination
acknowledge Gaspard andDuchêne,Johannes
Sahlmann, Jorgelevels
Lillo-Box, 2009;
Carpenter,J. M., Mamajek, E.etE.,al.
Hillenbrand,A., & Meyer,
R. 2006,Yasui et al. et al. 2013), and so are more
ple with very We
lowhighcompletenessRoccatagliata
2011; L.Fanget al.M.2012;
likely to form gas giants. On the other hand, the present study
ApJ, 651, L49
José
Antonio
Caballero,and vs.
JustinCreppfor helpful
on disks,
Á. Ribas
et al.: Protoplanetary
disk lifetimes
stellar massand possible
implicationsdiscussions
for giant planet populations
to derive accurate
disk populations,fractionsmigrationand protoplanetarydiskof known
dissi-ex- 2014),
these A.,studiesfocusedonindividual
regions
and/orJovian planet
and biases in the sample
Chabrier, but
G., Johansen,
Janson, M.,& Rafikov,
R.shows
2014,
in thatProtostars
andmassivestars
disperse their disks earlier. Under
A&A 576, A52(2015)
oplanets.Their
H. Bouy
is funded inby good
the Spanish
Ramón y Cajal
fellowship
Planets
VI, eds. H. Beuther,
et al. (Tucson:
University of Arizona
Press), 619 (e.g., lowpation timescales.
results,
agreement
withprevi-pro- werepotentiallya↵ectedby di↵erent
systematics
circumstances,
the formation of giant planets could
ApJ, 667,
308
number RYC-2009-04497. This work has made an extensive use of Topcat Cieza, L., Padgett, D. L., Stapelfeldt, K. R., et al. 2007,someous estimates gram
(Haisch
et
al.
2001;
Hernándezetal.
2007b,
2008;
numberstatistics,
sensitivity
biases,contaminationfromback(TOPCAT http://www.star.bristol.ac.uk/~mbt/topcat/ and STILTS,
Close, L. M., Follette, K. B., Males, J. R., et al. 2014, ApJ,
781,
L30for some Myr (Helled et al. 2014). Therefore, a fast
even
last
Taylor
2005, 2006).
work isbased in part
obtained
as part ofComerón, giant
F. 2008, stars,
The Southern
Sky ASP
Monograph Publications,ed. B.
Mamajek 2009;
Fedeleet al.This
2010;
Murphyet on
al.data2013)
servedasthe grounddi↵erent
completenesslevels).
As nota result,
dispersing
disk
may
live long enough to form gas giants.
Reipurth, 5, 295
UKIRT Infrared Deep Sky Survey. This research made use of the SDSS-III cataa benchmark to
theforreliability
andprovided
robustness
ofP.thedirectcomparisonof their results isIfdifficultand atrend
statisticallythe observed(massive stars being more likely to hostlog. test
Funding
SDSS-III has beenby the Alfred
Sloansample
Foundation, aCurrie,
T. 2009,
ApJ, 694, L171
Currie, T.,
Lada,
C.
J.,
Plavchan,
P.,
et
al.
2009,
ApJ,
698,
1
the Participating
and methods usedin R14.Institutions, the National Science Foundation, and the US robust
giant planets)
is disk
real, theneither a strong correlation between
confirmation of a possible dependenceof thelifetime
Department of Energy Office of Science. The SDSS-III web site is http://www.
Dahm, S. E., & Hillenbrand, L. A. 2007, AJ, 133, 2072
stellarmassA.haselusive
far.stellar mass and the ability to form giant planets exists to
sdss3.org/. SDSS-III is managed by the Astrophysical Research Consortium withDuchêne,G., & Kraus,
2013,been
ARA&A,
51, 269 so theshorter disk?
C., second
Teixeira, R.,paper
Galli, P. of
A. B.,
et al.
2014, compensate
A&A,
A121forthe
Participating
Institutions offormthe SDSS-III
Collaboration including Ducourant,
In thistheseries,we563,
study
the theinfluenceof lifetime, or giant planet forAppendix Aforisavailableinelectronicatthe University of Arizona, the Brazilian Participation Group, Brookhaven Dzib, S., Loinard, L., Mioduszewski, A. J., et al. 2010, mation
ApJ, 718, is
610faster than the disk dispersal time regardless of the
http://www.aanda.org
on diskevolution
using
the large
and homogeneous
National Laboratory, University of Cambridge, University of Florida, the French stellar
Espaillat,mass
C., Muzerolle,J., Najita,J., et al. 2014,
in Protostars
and Planets
VI,!

Finally, we also note that the higher fraction of evolved
star for the known exoplanets1 (planets around pulsars are not disks – which include circumbinary disks – around high-mass
plotted since they represent a completely di↵erent type of ob- stars is also consistent with their higher multiplicity frequencyject). Considering as hot Jupiters the planets with masses above (Duchêne & Kraus 2013).

1 M J and periods <10 days, a clear lack of these objects aroundmassive stars (>1.5 M ) is easily identified. This paucity of close
Jovian planets around massive stars is still a matter of intense
debate. Some authors indeed report a real dependence on stellar
mass (e.g. Johnson et al. 2011), and have even shown hints of it
not being the result of planet engulfing caused by the post-mainsequence evolution of the star (Johnson et al. 2007). On the other
hand, other authors suggested that such claims were made on
incorrectly assigned stellar masses (Lloyd 2011), and that tidal
destruction is responsible for the observed trend (Schlaufman &Winn 2013). By also plotting the radii of the planet host stars in
Fig. 6, it becomes clear that most low-mass stars (<1.5 2 M )stellaror most of the planet formation process happens have radii matching those predicted from main-sequence models
eds. H. Beuther, et al. (Tucson: University of Arizona
Press), mass,
497
Participation Group, the German Participation Group, the Instituto de Astrofisica
de Canarias,
Dame/JINA Participation Group, Johns
Evans, II, N. J., Dunham, M. M., Jørgensen, J. K., et al.in2009,
321 years.
ページ 11/12 the Michigan State/NotreArticle
ページ
theApJS,
first181,
millionpublishedby
EDP
Sciences
A52,
page
1
of
13
(Siess et12
al./12
2000). For more massive stars, the stellar radii inHopkins University, Lawrence Berkeley National Laboratory, Max Planck Fang, M., van Boekel, R., King, R. R., et al. 2012, A&A, 539, A119
2. Type-II migration is active while the gaseous disk is present crease owing to the post-main-sequence phase, overlapping with
Institute for Astrophysics, New Mexico State University, New York University, Fedele, D., van den Ancker, M. E., Henning, T., Jayawardhana, R., & Oliveira,and the planet can exchange angular momentum with it. The the loci of the hot Jupiters semi-major axes. It is not possible to
Ohio State University, Pennsylvania State University, University of Portsmouth,J. M. 2010, A&A, 510, A72
A&A, 446, 501of orbital semi-major axes produced by these
Princeton University, the Spanish Participation Group, University of Tokyo, Franciosini, E., Pallavicini, R., & Sanz-Forcada, J. 2006,
distribution
tell if the massive stars in Fig. 6 have engulfed any close gas giUniversity
ofdisks
Utah,
Vanderbilt
University,
UniversityofThe
Virginia,
University
Garcia Fig.
Lopez,
R., Natta,
A., Testi,& Habart,
E. 2006,
459,
837
Fig. 2. Evolutionof protoplanetary(red) andprocessed (yellow)
disks for eachmass and age bin.
number of sourcesin each bin is shown
3. Frequencies
for protoplanetary(red),L.,
evolved(blue) and diskless(gray)A&A,
objects. The
numberof sources in each bin is shown in brackets.
migrationis controlled by the ratio of the disk ants, but it becomes clear that the post-main-sequence phase ofwithin brackets. Errors are less than 6% (see Table 3).
Errors arethan
8% (see Table3). Megeath, S. T., et al. 2008,of Washington, and Yale University. This work makes use of data from the Gutermuth,
R.lessA.,
Myers,
P. C.,
ApJ, 674, mechanisms
336
Jr., 3.K.ResultsE., from
Lada,
E. A.,of&
C.asJ.a 2000,
AJ,
120,
1396
DENIS Survey. DENIS is the result of a joint e↵ort involving human and Haisch,Table
depletion
timescaletotheviscousdi↵usiontimescaleatthethe analysis
diskLada,frequencies
function ofage andmass.
these stars could account for a fraction of the observed paucity
form of pie charts in Fig. 2 and in Table 3. Errors for the disk same procedure as in the previous case and excluded the regions
financial contributions of several Institutes mostly located in Europe. It has Haisch, Jr., K. E., Lada, E. A., & Lada, C. J. 2001, ApJ,distance553, L153 where Jovian planets form (Armitage et al. 2002;
fractions were estimated using the bootstraping method further with significantly worse limiting spectral type ( Orionis). Theof hot Jupiters around them. Assessing the origin of this lack of
Disk type
Young+Low-mass[%] Old+Low-mass
been4.3,supported
financiallymainly by the
FrenchInstitutR.,
Bodenheimer, P., Podolak,M., et[%]al.Young+High-mass
2014, in Protostars
and [%] Old+High-mass [%]described in Sect.
and are typicallyless than 6%.
final
sample waslimitedNationalto spectral destypesSciences
earlier than M4Helled,(theProtoplanetarydisks(Sect.
4.1)
Trilling
et
al.
2002;
Ida
&
Lin
2004b).
A
fasterdispersalofAfter inspection
of
Fig.
2
some
conclusions
canbedrawn:
corresponding
limiting
spectral
type
of
CrA).
As
a
result,thetoPlanets VI, eds. H. Beuther, et al. (Tucson: University of Arizona Press), 643
de l’Univers, CNRS, and French Education Ministry, the European Southernclose gas giants is beyond the scope of this study, but the fastertal number of sources meeting this criteria dropped to 389. While
Protoplanetary (IRAC4 detection limit)
63 ± 238 ± 617 ± 22+4

5.3. Transitional disks

We stress that the current data in our sample cannot be used toisolate transitional disks among evolved disks. With this limitation in mind, it is nevertheless interesting to compare theA52, page 8 of 13

properties of evolties of transitional
Transitional di
cumstellar disks th2011). Various meexistence of these gformation of planeevaporation or the
et al. 2012; Alexa
connection betwee
strated by several
ions within the gap
Ireland 2012; Quaa cornerstone for t
First, we note
our analysis (betwdisk fractions repo
Kenyon & Hartma
et al. 2007; Merín
Second, previo
transition phase toagreement with the⇡10 Myr disk lifet
a transitional phas
Jovian planets are
disk (Lin & Papal
et al. 2000; Calvealso plausible thattional phase durin
estimated ⇡1 Myrlower limit.
5.4. Implications
migration

To date, there are t
planets: the core amodel (DI) (see H
recent reviews on– In the CA sc
by continuum
Terrestrial plan
core becomes mat an increasin
the disk (eithe
the planet has
is a gas giant p– In the DI mod
own gravity, leand potential f

So far, the CA modproperties of the k
are known to exist
planet occurrence
ities (e.g. Johnson
plained by the CA
more efficient planetary systems cannsystem around HR
may be responsibl
large radii from thstars. It could also