星形成ゼミ 2015/ 4/17 Newsletter #267の41–45本目!

!

!

43. “ASSESSING MOLECULAR OUTFLOWS AND TURBULENCE IN THE
PROTOSTELLAR CLUSTER SERPENS SOUTH”!
A. L. Plunkett et al.!
!
http://ads.nao.ac.jp/abs/2015ApJ...803...22P!
!
ApJ!

!

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

!

41. “Another deep dimming of the classical T Tauri star RW Aur A”!
Peter P. Petrov et al.!
!
http://ads.nao.ac.jp/abs/2015arXiv150304158P!
!
A&A accepted!

!

44. “SMA OBSERVATIONS OF W3(OH) COMPLEX: PHYSICAL AND CHEMICALDIFFERENTIATION BETWEEN W3(H2O) AND W3(OH)”!
Sheng-Li Qin et al.!
!
http://ads.nao.ac.jp/abs/2015arXiv150204467Q!
!
ApJ accepted!

!

45. ”Protoplanetary disk lifetimes vs. stellar mass and possible implications for giant planet
populations”!
Álvaro Ribas et al.!
!
http://ads.nao.ac.jp/abs/2015A%26A...576A..52R!
!
A&A!

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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.

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),h30m(Fuller
00s, decl. &= Ladd 2002; Arce & Sargent 2006), inject
Our aremosaicwas centered cluster
at R.A. = 18
gasmolecular clouds and its low-level rotational stateseasily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 spacing of 30´ and etvertical25.´8. In total,al. spacing
2007;of 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 weremotionsthataffectas standard high-velocity (SHV) with velocities 30ofs pera mosaic
few position,
to and 3 minutes on the phase calibrator ongoing star formation (see also Frank
1743-038, with 33 pointings observed in between each phaseal.primary2014).
Models of clustered star formation must
about 20 km s−1, and extremely high velocity (EHV)
withobservation.
even Uranusetwas thecalibrator
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
MHz (_80 km s−1) and 8 MHz(_20 km s−1), respectively.
includingstar
formation rates (SFRs) and efﬁciencies (e.g., Li
EHV, and in this paper we conduct a census of 31
mostlySHVWith 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. (Images of the data were obtained using the task invert,
specifying a cell size of 2″, and then deconvolved with the
maximum entropy deconvolution mosmem. The clean images
were then restored with a Gaussian beam. We subtracted thecontinuum image from the molecular line maps using
avmaths in the image domain. A summary of the CARMA
observations is given in Table 1.
2.2. IRAM Observations
We mapped the same region with the IRAM 30 m telescopein 2012 April. The center of the map is R.A. = 18h 30 m03s,
decl. = 02n02a to match the CARMA map coverage. We
mapped the region using individual 2a q 2a on-the-ﬂy (OTF)

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 associatedsubstructurelack 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)!
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 bound system19,20. The observed kinematics and separation whelm the combined turbulent and thermal pressure in all the identified(.1,000 astronomical units, AU) of the B5 system is significant because condensations, and a star will probably form in each case. The timescale
it demonstrates a clear mechanism for wide binary formation and pro- for these condensations to undergo gravitational collapse is approximately
vides convincing evidence that the observed condensations will become the gas free-fall time, which we estimate to be 40,000 years (Methods).
This timescale is sufficiently short to ensure that the system’s spatial
a bound multiple star system.

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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
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,
Vitacura, Santiago, Chile. 12National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia 22903, USA.

3

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

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 normalof RW
Aur A obrightne405(the)-405(whole)404(on)-4059.962607Tm-289(oc-)]TJET
Qq0.654988
brightness levels
flows in the imme
and components
ner part of a stell
Fig.
4.
Left
panel
–
ALFOSC
spectrafrom2014
coveringtheOn the other hFig. 3. Three spectra obtained during the bright state in
1 1996
Fig. 4. Left panel – ALFOSC spectra from 2014 covering the ponents related to
[O
i]
6300
line.
The
strong
lines
at
-90
and
+90
km
s
(solid)showing typical variations in the He i and Na i D lines, and one[O i] 6300 line. The strong lines at -90 and +90 km s 1 (solid) as enhanced blue
are
from
RW
Aur
A,
while
the
singleemissioncomponent
atcorresponding spectrum in 2014 (right).
zero velocity (dashed) is from RW Aur B. Right panel – high res- are from RW Aur A, while the single emission component at Ca ii. Moreover,olution spectra of the same region. Dashed curve shows the [O i] zero velocity (dashed) is from RW Aur B. Right panel – high res- considerably stro
olution spectra of the same region. Dashed curve shows the [O i] locities of the tw
profilefrom bothflux
RWinAurandremained
B in 2014,
where
centralthat
thethe Alines
about
thethesame
af! observed
peak
to RW Aur B, and the red-shifted sharp peak at +90 profile from both RW Aur A and B in 2014, where the central from the opposite
terthebelongs
1star dropped in brightness. This implies that the occultpeak belongs to RW Aur B, and the red-shifted sharp peak at +90 The change in eq
kmbodys belongs
RW
Aur
blue-shiftedis shaling
coveredtothe
starandA.itsThe
inneraccretion emission
flow, wherethekm s 1 belongs to RW Aur A. The blue-shifted emission is shal- as caused by an olow. Numerous
narrow
telluric form,
absorption
components
broad
emission lines
of metals
and also
the innerand
partphoof
tosphericlines fromcomponent
are present.
The solid
curve low. Numerous narrow telluric absorption components and pho- 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 isparts
the panelwind shows(the jets)
not occulted
much.
Thiscentral
setspeak atrestrictions
zero velocity
and
two
peaks
at -140 and+ 140 km s 1 . in the right panel shows the [O i] profile in 1996, with the central cloud containing
certain
to
the
size
of
the
obscuring
cloud.
!forbidden lines are much stronger in the faint than in the peak at zero velocity and two peaks at -140 and + 140 km s 1 . these grains once
TheThephotometry shows that the extinction becomes very grey The forbidden lines are much stronger in the faint than in thebrightthestate.
RW Aur A is inc
whenstar fades implying that large dust grains must be bright state.
ways to explain hpresent in the obscuring body. The large extinction of Av ⇠ 3m
disk, into the lin
in combination with the absence of additional absorption lines
occulting body isin resonance lines of Na i and K i indicate that the dust-to-gas
not rule out thisratio in the obscuring cloud could be very di↵erent from normalthat the stellar w
conditions in the interstellar medium.
that an interactio
The dimmings of RW Aur A are of a di↵erent nature thandisk caused an e
those observed in AA Tau and KH 15D (see Sect. 1). In the
did not find any
hypothesis proposed in R2013 the occultation of RW Aur A in
obscuring cloud.
2010 was caused by a dusty structure connected to a tidally disrupted disk. Since the phenomenon has returned a second time
Acknowledgements.
one must assume that in this case the foreground structure conTelescope, operated
tains more than one condensation passing in front of the star.
(NOTSA) at the ObseFig.
The regionexplanation
covering theresonance
K i 7699theline
and theAn5.alternative
to explore
is whether
dimming
the Instituto de Astroi 7773 bytriplet.
– bright statein 1996,
right panel
isOcaused
someLeftkindpanelof interaction
betweena changing
out-– Fig. 5. The region covering the resonance K i 7699 line and the with ALFOSC, whicOi
7773
triplet.
Left
panel
–
bright
state
in
1996,
right
panel
–
faint
state
in
2014.
Note
the
absenceofanyenhancedKiab(IAA) under a joint aflow and the disk leading to a transport of larger dust grains fromsorptioncomponents
in the spectrum
the dimmed
star;disk.
there faint state in 2014. Note the absence of any enhanced K i ab- KG acknowledges fun
thediskplane
into the line-of-sight
far ofabove
the1inclined
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 .

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!

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into the wind can be caused either by disk erosion (e.g. Schnepf
et al.2014), or by the MHD processes at the disk-magnetosphereboundary (e.g., review by Romanova et al. 2013). In this case the et al.2014), or by the MHD processes at the disk-magnetospherestar and the H↵-emitting region are obscured by a screen related
3 boundary (e.g., review by Romanova et al. 2013). In this case thestar and the H↵-emitting region are obscured by a screen related
to a dusty wind containing large dust grains.
The remarkable phenomenon of the long-lasting deep min- to a dusty wind containing large dust grains.
The remarkable phenomenon of the long-lasting deep minima in RW Aur A rises a number of questions. How recurrent arethese occultations? Are there any similarities to UX Ori stars? ima in RW Aur A rises a number of questions. How recurrent are
Are the dimmings in fact related to a changing stellar wind? It these occultations? Are there any similarities to UX Ori stars?
is premature to draw any deeper conclusions about the cause of Are the dimmings in fact related to a changing stellar wind? It
the phenomenon. Continued spectroscopic, photometric and po- is premature to draw any deeper conclusions about the cause oflarimetric observations of the star will provide more insight, and the phenomenon. Continued spectroscopic, photometric and poit will be interesting to see whether the normal wind features are larimetric observations of the star will provide more insight, andit will be interesting to see whether the normal wind features arerestored when RW Aur A returns to its normal state.
restored when RW Aur A returns to its normal state.

!

5. Conclusions

5. Conclusions
We have observed the binary T Tauri star RW Aur A and B spectroscopically and photometricly during the deep, long-lasting We have observed the binary T Tauri star RW Aur A and B specminimum that occurred on the primary in 2014. This is the sec- troscopically and photometricly during the deep, long-lasting
ond time over at least 50 years that RW Aur A su↵ers from such 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 such4

ページ 8/12

References

Alencar, S. H. P., Bas
Antipin, S., Belinski,
Bouvier, J., Grankin,
557, A77
Cabrit, S., Pety, J., Pr
Chou. M.-Y., Takami
Dodin A.V., Lamzin S
Gahm, G. F., Gullbrin
A&AS, 100, 371
Grankin, K. N., Meln
2007, A&A, 461
Grinin, V., Rostopchi
Grinin, V., & Tambov
Grinin, V., Stempels,
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

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

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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
outflow orientation traced by water masers and radio jets. A north-south velocity gradient across W3(OH) is also observed. However, with current observations

1
Department of Astronomy, Yunnan University, and Key Laboratory of Astroparticle Physics of Yunnan
Province, Kunming, 650091, China; slqin@bao.ac.cn
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
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
Korea Astronomy and Space Science Institute 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
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。!
azhuang 050024, China
W3(H2O)と、Sgr B2(B) hot core, Orion KL hot core, Orion Compact Ridgeを比較し、分子種の化学的起源について議論する。!

!

CH3OH, 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!

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。!
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。!

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Astronomy&
Astrophysics

A&A 576, A52 (2015)
DOI: 10.1051/0004-6361/201424846
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
ABSTRACT
Aims. We study the dependence of protoplanetary disk evolution on stellar mass using a large sample of young stellar objects innearby young star-forming regions.
Methods. We update the protoplanetary disk fractions presented in our recent work (Paper I of this series) derived for 22 nearby(<500 pc) associations between 1 and 100 Myr. We use a subsample of 1428 spectroscopically confirmed members to study the
impact of stellar mass on protoplanetary disk evolution. We divide this sample into two stellar mass bins (2 M boundary) and two age
bins (3 Myr boundary), and use infrared excesses over the photospheric emission to classify objects in three groups: protoplanetarydisks, evolved disks, and diskless. The homogeneous analysis and bias corrections allow for a statistically significant inter-comparisonof the obtained results.
Results. We find robust statistical evidence of disk evolution dependence with stellar mass. Our results, combined with previous
studies on disk evolution, confirm that protoplanetary disks evolve faster and/or earlier around high-mass (>2 M ) stars. We also finda roughly constant level of evolved disks throughout the whole age and stellar mass spectra.
Conclusions. We conclude that protoplanetary disk evolution depends on stellar mass. Such a dependence could have important
implications for gas giant planet formation and migration, and could contribute to explaining the apparent paucity of hot Jupiters
around high-mass stars.
Key words. protoplanetary disks – planet-disk interactions – stars: formation – planetary systems – stars: pre-main sequence

1. Introduction
Studying the evolution of protoplanetary disks around youngstellar objects (YSOs) is crucial to better comprehend how planets form. Understanding the e↵ects of fundamental parameterssuch as age, stellar mass, environment, or binarity provides insight into the planet formation processes, and could shed lightonto important relations between planetary systems and theirstellar hosts. Several studies on this topic (see e.g. Williams &Cieza 2011; Alexander et al. 2014; Espaillat et al. 2014) havealready drawn a comprehensive picture of protoplanetary disk
evolution.
In the first paper of this series (Ribas et al. 2014, hereafterR14) we presented a large sample of ⇠2300 spectroscopicallyconfirmed YSOs in 22 young (<100 Myr) and nearby (<500 pc)
star-forming regions and associations; R14 used this large sample with very low contamination and high completeness levels
to derive accurate disk fractions and protoplanetary disk dissipation timescales. Their results, in good agreement with previous estimates (Haisch et al. 2001; Hernández et al. 2007b, 2008;
Mamajek 2009; Fedele et al. 2010; Murphy et al. 2013) served asa benchmark to test the reliability and robustness of the sample
and methods used in R14.

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Appendix A is available in electronic form at
http://www.aanda.org

ページ 11/12

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The evolution of protoplanetary disks and planets are closely
connected as the disk properties define the initial conditions ofplanet formation, and the planet evolution subsequently a↵ects
the disk properties (e.g. Kley & Nelson 2012; Baruteau et al.
2014). The disk lifetime is a fundamental parameter to understand the interplay between a disk and its planets, since it determines the time available for the planets to form and migrate(e.g. Burkert & Ida 2007; Alexander & Armitage 2009). A dependence of the disk lifetime on the stellar mass could havestrong implications for planet formation and migration theories(see e.g. Raymond et al. 2014; Helled et al. 2014; Baruteau et al.
2014, for recent reviews on this topic), and ultimately result indi↵erent planetary populations around low-mass and high-massstars. Previous studies have already found hints of protoplanetary disks evolving faster around massive stars (e.g. Carpenter
et al.! 2006; Dahm & Hillenbrand 2007; Kennedy & Kenyon
2009; Roccatagliata et al. 2011; Fang et al. 2012; Yasui et al.
2014), but these studies focused on individual regions and/or
were potentially a↵ected by di↵erent systematics (e.g., lownumber statistics, sensitivity biases, contamination from background giant stars, di↵erent completeness levels). As a result,a direct comparison of their results is difficult and a statisticallyrobust confirmation of a possible dependence of the disk lifetime
with stellar mass has been elusive so far.
In this second paper of the series, we study the influence ofstellar mass on disk evolution using the large and homogeneous

Article published by EDP Sciences

A52, page 1 of 13