星形成ゼミ	
  2017/4/14	
  SFN	
  #291	
  #1-­‐#5		
  長谷川哲夫（国立天文台）	

R. Avramenko1 , S. Wolf1 , and T.F. Illenseer1
University of Kiel, Institute of Theoretical Physics and Astrophysics, Leibnizstrasse 15, 24118 Kiel, Germany
e-mail: [ravramenko;wolf;tillense]@astrophysik.uni-kiel.dearXiv:1702.02862v2 [astro-ph.SR] 21 Feb 2017

ABSTRACT
Context. A substantial fraction of protoplanetary disks form around stellar binaries. The binary system generates a time-dependent

non-axisymmetric gravitational potential, inducing strong tidal forces on the circumbinary disk. This leads to a change in basic
physical properties of the circumbinary disk, which should in turn result in unique structures that are potentially observable with thecurrent generation of instruments.
Aims. The goal of this study is to identify these characteristic structures, constrain the physical conditions that cause them, andevaluate the feasibility of observing them in circumbinary disks.
Methods. To achieve this, first we perform 2D hydrodynamic simulations. The resulting density distributions are post-processed with
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a 3D radiative transfer code to generate re-emission and scattered light maps. Based on these distributions, we study the influenceof various parameters, such as the mass of the stellar components, mass of the disk, and binary separation on observable features inTable 1. Hydrodynamic simulation parameters and ranges.
circumbinary disks.
Results. We find that the Atacama Large (sub-)Millimetre Array (ALMA) as well as the European Extremely Large Telescope(E-ELT) are capable of tracing asymmetries in the inner region of circumbinary disks, which are a↵ected most by the binary-disk
interaction. Observations at submillimetre/millimetre wavelengths allow the detection of the density waves at the inner rim of the disk
ParameterP
and inner cavity. With the E-ELT one can partially resolve the innermost parts of the disk in the infrared wavelength range, including
Massofindividual
binarycomponents
M[M]
B
the disk’s rim, accretion arms, and potentially the expected circumstellar disks around each of the binary components.
Key words. Accretion, accretion disks – Binaries: general – Hydrodynamics – Radiative transfer• Circumbinary	
  diskの観測的特徴を明らかにしたい	
  turn result in unique structures such as a cavity in the disk centre,• 2D流体計算＋3D輻射輸送	
   accretion arms, density waves, and spiral arms (Günther & Kley
The development of circumstellar disks around single stars and 2002; Günther et al. 2004). Previously it was shown that sometheir formation
during the collapse of molecular clouds
were manuscript• ALMAとE-­‐ELTならば、disk内域の非対称構造を観測できる	
of these structures,such as the cavity and parts of the accretion
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no. 30084_am
1. Introduction

Binary mass ratio
Initial disk mass
Semi-major axis
Inner boundary radius
Outer boundary radius
Eccentricity
Simulation timestudied extensively in the past three decades (McKee & Ostriker arms, are observable with the current generation of instruments
2007;Table
Andrews
et al. 2010).simulation
However,at the sametime it was (Ruge et al. 2015). There are also examples ofTable
Radiative transfer simulation parameters and ranges.
more2.complex
1. Hydrodynamic
parameters
and ranges.
found that a substantial fraction of protoplanetary disks form systems, such as GG Tau, which is a triple star system (Piétu
around binaries (Duquennoy & Mayor 1991; Kraus & Hillen- et al. 2011; Di Folco et al. 2014). Some of the disk structures,
brand 2009).
for example spiral arms
potentiallyassociated withembeddedParameter
Parameter value/range
Parameter
Parameter
value/range
planets,M
have
also been observed
Garufiet al.
The assumed mechanism bywhich
a interstellar
Binarytemperature
T B [K]
4500, 5600, 6600
Massof individual
binary cloudcomponents
0.5, 1, (Muto
1, 5 et al. 2012;
B [M ]
evolves into a binary is either byBinary
multiplefragmentation
or stel- 2013; Stolker
Binary radiusRB [R ]
0.42, 0.98, 1.4
mass
ratio
q et al. 2016).
1
lar capture (Bate 2000; Bate et al.
2002;
Wolf
et
al.
2001).
DurBinaryluminosityL[M]
0.45, 0.87, 8.2
Initial disk mass
[M ] of a planet
0.02further complicates the case
B
TheMdiskpresence
ing the collapse both componentsSemi-major
can develop
accretion disks in
Innerradiusr[AU]
0.6 · a
axis
a
[AU]
10,
20,
30
in
Nelson 2013). A binary can also feature two types
the same way as a single star system
et al.
2011; Müller (Pierensr&
Outerradiusr[AU]
200
Inner(Regály
boundaryradius
[AU]
0.6
·
a
out
in orbits: planets that revolve around a single binary& Kley 2012). In addition to that,the dust
and gasfrom outside of planetary
Simulation
wavelengths
[µm]
0.05
. . . 2000
Outer
boundaryradius
r
[AU]
1000
sim
out
around both stars (Welsh et al.
of the binary orbit form a third Eccentricity
disk, the so-called cirumbinary component "and planets that orbit
0
disk (Rodriguez et al. 2010; Romero et al. 2012; Lines et al. 2012; Orosz et al. 2012; Schwamb5et al. 2013). Furthermore, theSimulation timet [yr]
5 · 102015), which revolves around the centre of mass of the binary. growth ofsimplanetesimals from smaller objects might be hindered(Meschiari
2012;
Marzarietal. 2013). The migration
behaviour
3. Results
Furthermore, it is possible for a newly developed single star sysof
planets
inabinarysystemcanalso
beproblematicsince thetem toTable
capturea transitingstarsimulation
to form aparameters
binary. Inandthisranges.
case the2. RadiativetransferIn
thisthechapter,
planetcanbeejectedifitentersa
4:1
resonancewithbinary we present the results of the simulations deprimary could retain a part of its original disk and the captured
scribed in Sect. 2. We discuss the most notable features of bi(Kley&
Haghighipour
2014;
Pierens&
Nelson2008).
star would accrete a disk of its own. The remaining matter would
nary systems and their influence on the radiation emitted by theform thecircumbinary disk.
force on disk,
the i.e. observable quantities.
Parameter
Parameter value/range As mentioned above, the binary exerts a tidal
surrounding
Regardlessof the exact formation
the major
Binary temperature
T Bmechanism
[K]
4500,
5600,dif6600 surrounding disk, which should result in characteristic strucferenceBinary
from single
binarygenradiusstar systems is Rthat] orbiting0.42,
0.98,
1.4 tures. However, specific questions regarding the exact nature of
B [Rthe3.1. Characteristic
structures in the density distribution oferates aBinary
time-dependent
non-axisymmetricgravitational
potenspatial
luminosity
LB [M ]
0.45, 0.87, 8.2 these structures are still open: How do the characteristic
circumbinarydiskstial, inducing
strong tidal forces onrthe
surrounding
gas
and
dust
scalesandtimescales
dependonparameters
such
as
the
mass
and
Inner radius
[AU]
0.6
·
a
in
(Artymowicz& Lubow 1994). Thisin basic binary separation? How does the overall density and temperature
Outer radiusr leads[AU]to a change 200q
Mdisk [M ]a [AU]
rin [AU]
rout [AU]"tsim [yr]Fig. 1. Normalis
crosses indicate t

Another str

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Fig. 4. Normalised dust density distribution for a = 20 AU, MB = 1 M . The crosses indicate the positions of the binary components.
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Fig. 18. Synthetic ALMA observation at# 7; MB = 1 M and a = 30 AU.

= 320 µm with configuration

Fig. 5. Left: Optical depth ⌧ at = 526 nm calculated in the midplane. Right: resulting midplane temperature of the same system. Black line
denotes where the optical depth reaches values of ⌧ = 1 and ⌧ = 104 , respectively.

3.2.1. Spectral energy distribution

binary), the decreased flux due to the cavity and the lower disk
temperature (due to the lower net stellar luminosity) result in alower infrared emission in the case of the circumbinary disk.
At longer far-infrared to millimetre wavelengths the di↵erences of the fluxes arising from circumbinary disks with di↵erent
cavity radii (i.e. di↵erent semi-major axes of the central binary)
become negligible (Fig. 7). This is expected since, for large distances from the central binary, the disks structures are similar.
The dependence on the disk mass can be seen in Fig. 8. Boththe short wavelength and long wavelength radiation increases
with disk mass, although the increase is caused by di↵erent
mechanisms.
TheSurface
presenceof morehot dustat in=the
Fig. 17.
brightness
distribution
4.5vicinityµm (left)of
and = 20the binary
tofor
highershort
radiation.
µmleads(right)
MB =levels
1.5 M ofanda = wavelength
20 AU.
For long wavelength radiation, for which the disk is opticallythin in theverticaldirection,
dust mass
to a higher
Table
5. Integrated
fluxhigherof the scatteredandleads
re-emitted
radiation in units
column densityof radiating
particles.
of mJy originatingfrominside a circle with the radius of 2⇥a. The direct

The cavity in the centre of the disk is clearly visible in all calculated density distributions and emission maps and its observability is the topic of Sect. 3.3. Because of the cavity, circumbinary disks tend to have less dust, which can be heated to a hightemperature. This results in much lower flux at near- and midinfrared spectrum in comparison to circumstellar disks (Fig. 6).
Because the cavity radius scales with the semi-major axis a, theflux at short wavelengths is in indirect proportion to a (Fig. 7).
The flux of the corresponding circumstellar disk (i.e. same disk
mass and net stellar mass) is higher for all wavelengths (Fig. 6).
It was simulated using the standard set-up presented in Sect. 2.1
and a star with the mass of Mstar = 2 M . The luminosity of thisstar is higher than the combined luminosity of two stars with individual
masses brightness
Mstar = 1distribution
M . This atleads
to anhigherFig. 13. Surface= 4.5µmoverall(left) and= 20 Fig. 14. Surface brightness distribution at = 4.5 µm (left) and = 20
stellar radiation is not considered.
temperature
to higherµm (right) forandMB thus=1M
and a = fluxes.
10 AU. This opens a potentialµm (right) for MB = 1 M and a = 20 AU.
way for distinguishing between disks around binaries and single 3.2.2. Surface brightness distribution
stars: The outer edges of the circumbinary disk behave nearlyParameter= 4.5µm= 20µmand sensitivity.
Thethesynthetic
observations
are performed
with So far, we only considered
Keplerian.
Throughline broadening
one candetermine thethe synthetic SEDs of the simulated
M =the1.0M, a = 100.42
2702.1 valWe begin
discussion
of AU
the resultswith the highest

Fig. 19. Flux di↵erence between density wave maximum and minimumin units of sensitivity for a = 10 AU, MB = 1 M , and Mdisk = 10 1 M .

Monte Carlo code Mol3D and a temperature profile; in addition,

Either KP
Overall

1.6-2.5⇥10
0.4-2.5⇥1022

0.63
0.6-0.7

0.37
0.3-0.4

UCHII

1.0-1.9

0.6-1.2

0.4-0.7

Whilestarless vs.
star-forming DMRs
both the fraction
0.3-2.4 of 0.2-1.7
0.1-0.7
is insensitive to the column density cuts, the fraction
Draft version February 9, 2017of DMRs associated with methanol masers (fmaser ) inPreprint typeset using L TEX style emulateapj v. 01/23/15
TABLE
1
creases
as a function of column density (see Figure 3).
Various lifetime estimates in this work, described in more detail in §2, 3.2, 3.3, and 3.4. The column density listed refers
In the ‘generous’ and ‘conservative’ cuts, the methanol
to the lower bound on the column densities used to derive the relative lifetimes. The Anchor is the method used tomaserfraction
is 2% and
4% (fthe), respectively,
correanchor relative lifetimes to absolute timescales, CH3 OH masers
or UCHIIRegions.
Finally,
timescales
in the threemaser
right-mostcolumns are lifetimesofvariousphasesin
Myr.
THE LIFETIMES OF PHASES IN HIGH-MASS
STAR-FORMINGREGIONS
sponding to total DMR lifetimes (⌧total ) of 1.9 Myr and
0.9 Myr. Using the alternative column density threshCara Battersby1,2 , John Bally3 , & Brian Svoboda4
Draft version February 9, 2017
old fromKau↵mann& Pillai(2010), called KP10, as
Criteria
Median N(H2 ) Anchor ⌧f f [Myr]Nf f,total
Nf f,starless
Nf f,starry
is about 7-12%, corGenerous
0.6⇥1022
maser
0.7 discussed
2.8 in §2.1,
2.0the maser fraction
0.8
ABSTRACT
22
responding
to
total
DMR
lifetimes(⌧total ) of 0.3 andConservative
1.2⇥10
maser
1.8
1.3
0.5
High-mass stars form within star clusters
from dense,
molecular regions,but is the0.5
process of cluster
22
0.4
Myr.
Given
our
starlessandstar-forming
fractions
Either
0.6-1.2⇥10
UCHII
0.5-0.7of high-mass
1.7-4.8
1.2-3.4
0.5-1.4
formation slow and hydrostatic or quickand dynamic?
We link the physicalproperties
22
KP10 stage
near in a 3.5⇥10
maser
0.3
1.0
0.6and fstarry0.4
star-forming regions with their evolutionary
systematic way, using
Herscheland(fSpitzerdata.
=0.6-0.7=0.3-0.4) and the methanolstarless
22 these regions,
In order to produce a robust estimateKP10of the far
relative lifetimes
the fraction
2.2⇥10ofmaserwe compare
0.4 maser
1.1
0.4
lifetime (⌧0.7=35,000
years; van der Walt 2005)
maser
of dense, molecular regions above a column
with22high-massstar formation,
N(H2) >Eitherdensity
KP associated
2.2-3.5⇥10
UCHII
0.3-0.4
2.4-6.5
1.5-4.1
0.9-2.4
we cancalculate the total DMC lifetime and relative
0.4-2.5 ⇥ 1022 cm 2 , in the ‘starless’ (nosignature of stars& 10 M22 forming)
phasesOverall
0.6-3.5⇥10both and star-forming
0.3-0.7
1.0-6.5
0.6-4.1 the following
0.4-2.4 equations:
phase
lifetimes
using
in a 2 ⇥2 region of the Galactic Plane centered at `=30 . Of regions capable of forming high-mass

arXiv:1702.02199v1 [astro-ph.GA] 7 Feb 2017

A

stars on ⇠1 pc scales, the starless (or embedded beyond detection) phase occupies about 60-70% of⌧maserthe dense molecular region lifetime and the star-forming phase occupies about 30-40%. These relative⌧total =(4)TABLE
2 lifetimes
lifetimes are robust over a wide range of thresholds. We outline a method by which
relativefmaserVarious
lifetime
estimates
in
this
work,
as
in
Table
1,
in
units
of
the
numberoffree-falltimes.
A
robustestimateofthecan be anchored to absolute lifetimes from large-scale surveys of methanol masers and UCHII regions.
free-falloftime
not performedin absolute
this work,these
arestarless
simply phase
basedtoon
A simplistic application
this ismethodestimates the
lifetimeof thebesimple geometrical assumptions and the mediancolumndensity
pixelsin eachestimate.
These
are discussed
in =
more
detail⌧in §3.5.
⌧starlessfstarless(5)
0.2-1.7 Myr (about 0.6-4.1 fiducial
cloudfree-fallof
times)
andincluded
the star-forming
phase
to be 0.1-0.7
Myr
total
(about 0.4-2.4 free-fall times), but these are highly uncertain. This work uniquely investigates the5star-forminggions,nature orof high-column
pixel-by-pixel
our results
demonstrate
the Myr (forabout 3.6 density⇥ 105 gasyears
, ⌧U CHIIand. The
remain2.4thator 1.250%=orfstarry
100%, ⌧respectively)
when (6)
we⌧starrytotal
majority of high-column
densitygas
isinastarlessorembeddedphase.
ing link between the absolute andrelative
lifetimes isassume a starry fraction of 30% (fstarry ) as shown in the
Subject headings:
ISM: kinematics
and dynamics— dust,withextinction— HIIregions,regions — stars:
formationbelow.
the fractionof DMRs
associated
UCHII
equation

particularly O stars. =Wood& Churchwell (1989a) look
• 解析領域：銀河面L	
  	
  30°の２°四方の領域	
  ⌧U CHII
1. forINTRODUCTION(DMRs).
The regions
gas in regions(7)
only the brightest UCHII regions,dense
con-that form high-mass stars, ⌧total =fstarry fU CHII
• star解析データ：Herschel、Spitzer	  high anddensities(10cm ; Lada & Lada
taining
massivestars
stars,
more hasrecentmore
Whetherclusters
and high-mass
formwhileas the the DMRs,
2003) and cold temperatures (10-20 K; Rathborne et al.
result of slow, equilibriumcollapseof clumps(e.g., Tan et al.
sensitive
studiesof Anderson(e.g. 2011) show HII
absolute
lifetimes of the ‘Generous’ and ‘Conserva• over
解析の考え方：	  2010; Battersby et al. 2011) and isThetypically
detected
et al. 2006)
severalfree-fall
times
or
if
they
collapseregionsoverwider
evolutionary
stages,
after
much
of
the
by
submm
observations
where
the
cold,
dust
contintive’
thresholds
starless phase (using Equations 5 and 6)
quickly on the order of a free-fall time (e.g., Elmegreen
dense
gas
cocoonhas
beendispersed.
Wesuggestthat•
1pc（距離3kpcで1’）スケールで解析	
  uumemissionpeaks.
Given
theappropriate
viewing
an2007; Hartmann & Burkert 2007), perhaps mediated byis then 0.8-1.7 Myr and the “starry” phase would be 0.4these regions
can also
be seen in silhouette as Inworks investigate
of newer
catalogs
fromlarge scale accretionfuture
along filaments(Myers 2009),there- use gle,
0.7
If we
instead assume the KP10 column density
Dark
CloudsWood(IRDCs)absorbingdi↵use,mid• theA)	
  CORNISH
星形成可能	  DMR	
  かfrared
否か：N(H2)	
  >	
  0.4-­‐2.5	
  
×	
  the
1Myr.
0^22	
  
cm^-­‐2	  mains an open question.
The
stars thatform (e.g.;inthesePurcell
surveyetal.
2013).
&
thresholdforhigh-mass
star formation, fU CHII would
IR, Galactic background light. Battersby et al. (2011)
regions may disrupt Churchwell
and re-distributethe molecularma- three(1989a)searchedregions
in
our
l
=
30
showed
that
by
using
a
combinationofdata(including
not change, only the star-forming and starless fraction,
大without
質量星が誕生しているか否か：>	
  10	
  M_sunの星誕生の有無	  terial from which •they B)	
  formeddissociating it, alfield of(allclassifiedstudy) forofUCHII
remeasurements
their temperatures
columninand
thiscase,densi37% and 63%, respectively. The absolute
lowing future generations
star
formationas
in “starry”
the cloud in ourfrom Hi-GAL;
Molinari&et al. 2010), we can sample• long
解析結果	
  
gions
and (20-40
foundMyr,them
towardTherefore
the Wood
with overall
GMC
lifetimese.g.;Kawa- two.ties
lifetimes
inferred
for these phases based on association
5
DMRsofandas starless or star-forming inmura et al. 2009). The
scenario of (1989a)
quick, dynamic
star lifetime
Churchwellabsolute
3.6classify⇥ 10them
years
with
UCHII
regions
is のa ⌧期間は	  Myr, a ⌧starless
total of 1.0-1.9
a moresystematicallyrobustway thanjust using one yes	  • over
yes	
  の
うち	
  
B
)	
  
n
o	
  
の
期間は	
  
6
0-­‐70%、B)	
  3
0-­‐40%	
  formation sustained
a long time by continued inflow(⌧期間の相対比：A)	
  
to
2/3
of
‘starry’
U CHII ) corresponds, very roughly,of 0.6-1.2
and a ⌧starry of 0.4-0.7 Myr. See Table 1
wavelength (e.g. an IRDC would not
be ‘dark’Myr
on theof material is motivated by a variety of observations (disDMRs
theyのsurveyed
our analyzed
far sidefield.
of the Galaxy, so a mid-IR-only
wouldof the relative and absolute lifetimes forfor aselection
summary←A)	  しきい値を変えても変わらない	
  robust	
  な結果	
  cussed in detail in Elmegreen
2007;
Hartmann&inBurkertthose DMRs).
We by
make
the ofassumption
approximately
50-100%2007), and more recently
the lackstarless mas-that exclude
variousmethods.
Most
previous
studiesofhigh-mass
star-forming
region ×	
  0.6-­‐4.1）	
  sive protoclusters•(Ginsburg
et al. 2012;
Urquhart
nal.o	
  0.2-­‐1.7	
  Myr、B)	
  
y
es	
  
0
.1-­‐0.7	
  M
yr（t_ﬀ	
  of期間の絶対値：B)	
  our ‘starry’pixelsareetassociated
with
UCHII
regions
lifetimes
have
focusedondiscrete‘clumps’ofgas,
usually
2014a; Csengeri et al.(f2014) observed
through blind surbased on
thedustfollowing
U CHII ). This assumption is identified
in thecontinuum (e.g. Heyer et al. 2016; 3.5. Free-fall times
veys of cold dust continuum
emission in the Galaxy. Ad←メタノールメーザーとUCHIIから評価、不定性大きい	
  
lines
of
evidence:
1)8µm
emissionisoften
indicativeof et al. 2011). However, ofCsengeri
et
al.
2014;
Dunhamditionally, observations of infall of molecular material on
4 7

3Theofabsolute
tentimes‘clumps’containquiescencelifetimes derived in §3.3 and 3.4 are comUVet excitation(Baniaal. 2010,
showthesethat
nearlyall sub-regions
large scales (Schneider
al. 2010; Peretto
et al.et2013)
pared
with
fiducial cloud free-fall times. To calculate
and active
star formationwith
and cannot simply be classified
suggest that GMCs are
dynamic andevolve quickly,
GLIMPSEbubbles,
8 µmbut
emission,
are associatedas either. To lump these regions together and assign thethe ‘stathe ‘cocolumnbecaussolutelifetimestarlesstimes irelative3.

Since
uncertatimes othe ranWood &
times oon thefronts a
They emain-seHiGAL - Spitzer Lifetimes

3

Fig. 1.— A visual demonstration of the identification of DMRs and their classification as starless (pixels within the white contours) or starforming (pixels within the black contours). Shown in the top left panel is a zoom-in map of a portion of our total field with red/green/blue =70/8/4.5 µm emission. Shown in the top right panel and background throughout is the same zoom-in map with temperature shown in red,
8 µm emission in green, and column density in blue. Predominantly blue regions in this map are starless DMRs (cold, high column density,and 8 µm dark), while predominantly white regions are star-forming DMRs (warm, high column density, and 8µm bright). Red/yellow
regions are warm but have low column densities and thus are not classified as DMRs. In the bottom panels, DMR identifications as starlessor star-forming are shown in white or black contours, respectively, with the ‘generous’ criteria on the left and ‘conservative’ on the right.
The maser locations are shown by green contours. Both identification methods yield relative lifetimes of approximately 70% starless and30% star-forming, though the associated relative maser lifetime is 2% on the left and 4% on the right. (The sharp edges are source masks
from Battersby et al. 2011).

tive’ estimates.
Alternatively, we apply the Kau↵mann & Pillai (2010)
criteria (hereafter KP10) for high-mass star formation.

mine whether each DMR is ‘starless’ or ‘star-forming’.
The mid-IR signature is determined from the contrast at
8 µm, i.e. how ‘bright’ or ‘dark’ the pixel is above the4Battersby et al.
No

Is its temperature within 2σ (or 3σfor `generous’ cutoff) of the mid-IRdark distribution? (Conservative:
11-28 K, Generous: 9-32 K)

Dark

Is it above thecolumn density
threshold toform a massive
star?

Yes

What is it’s
mid-IRsignature?

Extremely rare - excluded

NeutralYes

Temperature
below 25 K

Starless

Temperature
above 25 K

StarryStarless

No
Is its temperature within 2σ (or 3σfor `generous’ cutoff) of the mid-IRbright distribution? (Conservative:
21-42 K, Generous: 12-47 K)

BrightNot included in thislifetime estimate

No

Yes

StarryExtremely rare - excluded

Fig. 2.— Flow chart showing the decision tree to assign pixels as DMRs, and subsequently classify them as ‘starless’ or ‘starry,’ i.e.
Battersby et al.
star-forming.

6

Number of Pixels

Percent

HiGAL - Spitzer Lifetimes
7
tive’ and ‘generous’ thresholds respectively), then it isfilter smoothing. PAH emission at 8 µm toward Photon42]100
2.5•10
Criteria
fstarlessAnchorclassified as ‘starless.’ If a DMR is mid-IR-brightandDominated N(HRegions
is notfstarry
likely to
be an⌧total
issue[Myr] ⌧starless ⌧starry2 ) [cm (PDRs)
Starless
Generous warm dust range
22
its temperature is withinthenormal
since
the
low
columndensity
towards
PDRs
will
Generous
0.4⇥100.71of dust
0.29
maser
1.9
1.4
0.6
Starless Conservative
22
Conservative
0.8⇥10
0.72
0.28
maser
0.9
0.6
0.3
(2 or 3- for the ‘conservative’
and‘generous’
threshpreclude
theirinclusion
as
DMRs
in
the
first
place.
Starry Generous
StarryConservative
Either
0.4-0.8⇥1022
0.7
0.3
UCHII
1.2-2.4
0.8-1.7
0.4-0.7
olds
respectively),
then
it
is
classified
as‘star-forming.’
480
2.0•10
Maser
KP10 near
2.5⇥1022 3. LIFETIMES
0.63
0.37
maser
0.3
0.2
0.1
Slight changes in the temperature distributions (e.g., inKP10 far
1.6⇥1022
0.63
0.37
maser
0.5
0.3
0.2
cluding all DMRs down to 0 K as starless and up to 100
Relative
Lifetimes1.0-1.9
Either3.1.
KP Assumptions
1.6-2.5⇥1022 in Deriving
0.63
0.37
UCHII
0.6-1.2
0.4-0.7
70% Pixels that
K as star-forming) have a negligible e↵ect.
22
Overall
0.4-2.5⇥10assumptions
0.6-0.7
0.3-0.4
both
0.3-2.4
0.2-1.7
0.1-0.7
Several necessary
madein inferring
4
HiGAL - wereSpitzer
Lifetimes
9
are mid-IR-bright and cold or mid-IR-dark and1.5•10
warm
60
the relative lifetimes of the starless and star-forming
are extremely rare and left out of the remaining analystages
forIdentification
high-mass star-forming
regions. f i) The sam- Reference
Clump
SF Identification
TABLE 1 starlesssis. When a DMR is mid-IR-neutral, its temperature isple is complete
anddescribedunbiasedin
inmoretimedetailand space.
ii)3.3,
The
IRDC
24
µm
0.653.2,
Chambersal. (2009)
Various
lifetime
estimates
in
this
work,
in §2,
and 3.4.etThecolumn density listed refers
used to classify it as ‘starless’ or ‘star-forming.’
A flowto the lower
bound
onof
the
columndensities usedderiveprobability
the relative
lifetimes.
The
AnchorIRDC
8toµm
0.82 allChambers
et al.
(2009) is the method used to
4
sum
DMRs
represents
with
equal
the
40
1.0•10chart depicting this decision tree is shown inanchor
Figure2.
relative
lifetimes
absolute
timescales,
orThisUCHIIRegions.
Finally,
the timescales in the three
3 OH masers
24
µm
0.33
Parsons
et al. (2009)
phases to
inIRDCtheformationof aCH
massivestar.ofcanberight-most columns 24
arevarious phases
This study is only sensitive to the star-forming
signaIRDCµmlifetimes
0.32-0.80
Perettoin&Myr.
Fuller (2009)30%
achieved eitherwith a constant starformation
rate (SFR)
tures of high-mass stars. Therefore, the term ‘starless’
LABOCA
8/24 µm
0.44
Miettinen (2012)or a large
enough sample. iii)
DMRs are formingor will
ATLASGAL
22 µm ⌧ [Myr] 0.25
Csengeriet al. (2014)refers only to the absence of high-mass stars forming; the3
Criteria
MedianN(H)
AnchorNN
Nf f,starry
2ff
ff,totalform high-massiv) The lifetimes don’t0.23
dependTackenberg
onf f,starless20
5.0•10ATLASGALstars.
GLIMPSE+MIPSGALet al. (2012)
region may support active low- or intermediate-massstar
22
Generous
0.6⇥10and star-forming
maser
0.7
2.8
2.0
0.8
mass.
v)Hi-GAL
Starlessregions
occupysimi24 µm
0.18
Wilcock et al. (2012)22
formation. Using Robitaille et al. (2006) models of dust
Conservative
1.2⇥10maser
0.5proportional
1.8
1.3
0.5
lar
areasHi-GAL
on thesky (i.e.
their
area
is
to
the
8
µm
0.33
Wilcocketal.
(2012)
22
emission toward YSOs, scaled to 5 and 9 kpc distances
Either
0.6-1.2⇥10
UCHII
0.5-0.7
1.7-4.8
1.2-3.4
0.5-1.4
a
BGPSvi) The signatures
mid-IR catalogs
0.80 brightDunham et al. (2011)
mass
of gas).
at 0.38 µm (mid-IR
KP10 nearBGPS3.5⇥1022
masertracers
1.0
0.6
0.4
and apertures, we find that our typical starless flux
00 limitmany
0.47
Svobodaet al. (2016)
or
dark)
and
their
associatedtemperature
distributions
22
KP10
far
2.2⇥10
maser
0.4
1.1
0.7
0.4
of
130
MJy/Srat
8µm(technicallyweuseacontrast
not
ATLASGALMIPSGAL
YSOs
Heyer et al. (2016)
0.0
0.5
1.0
1.5
2.0
are
good
indicatorsof 22thepresenceor absence0.69
of a highEither
2.2-3.5⇥10
UCHII
0.3-0.4
2.4-6.5
1.5-4.1
0.9-2.4
22−2
a specificflux, Density
but thisx is10theapproximatelevel for most
N(HKP
T
and
8
µm
0.60-0.70
Thiswork2 ) from Hi-GAL
dust
Column[cm]
mass
Pixels22above
the thresholdcolumnOverallstar. vii)
0.6-3.5⇥10both
0.3-0.7
1.0-6.5 den0.6-4.1
0.4-2.4of the region at our contrast cut) is sensitive enough tosity
containbeam-diluted
dense
cores
ratherthan
purelya RobitailleFig. 3.— The relative fraction ofdetect
pixels whicharevaststarlessmajority(70%) vs. star-forming(30%)
is shown in thedashed lines
as aal.
functionof the Red MSX Catalog from Urquhart et al. (2008), and the EGO catalog from Cyganowski et al. (2008).
et(2008),
the
(85%)
of
possiblemodel
YSOs
column density cuto↵. This fraction is mostly insensitive to the choice of column density cuto↵ over a reasonable range of thresholds
beam-filling low-surface density gas. We discuss thesewitha mass
10 M. Some YSO
models(N(H2 ) ⇠ 0.3 - 1.3 ⇥ 1022 cm 2 ) ina 2500 beamand toabovethe ‘generous’vs. ‘conservative’
parameters.
The left will
y-axis always
and dashed lines show
TABLE 2

ar census of the nearby, young 32 Orionis groupn P. M. Bell? ,1,2 Simon J. Murphy3 and Eric E. Mamajek4,5

Astronomy, ETH Zürich, Wolfgang-Pauli-Strasse 27, 8093 Zürich, Switzerlandtute for Astrophysics Potsdam (AIP), An der Sternwarte 16, 14482 Potsdam, Germany
hysical, Environmental and Mathematical Sciences, University of New South Wales, Northcott Drive, Canberra, ACT 2600, Australiaof Physics & Astronomy, University of Rochester, Rochester, NY 14627, USA
on Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA2

eceived ?; in original form ?

C. P. M. Bell et al.
32 OriKnown members
O stars
B stars
A stars
Pre-MS stars
Bellatrix

40

ABSTRACT

Key words: stars: kinematics and dynamics – stars: pre-main-sequence– stars:
however, we believe
furtherfundamental
astrometric, spectroscopic, and highparameters – solar neighbourhood – open clusters and associations:generalcontrast imagingof Bellatrix is warranted to see whether it could20µ [mas yr 1]

The 32 Orionis group was discovered almost a decade ago and despite the fact that it represents the first northern, young (age ⇠ 25 Myr) stellar aggregate within 100 pc of the Sun(d ' 93 pc), a comprehensive survey for members and detailed characterisation of the group
has yet to be performed. We present the first large-scale spectroscopic survey for new (predominantly M-type) members of the group after combining kinematic and photometric data to
select candidates with Galactic space motion and positions in colour-magnitude space consistent with membership. We identify 30 new members, increasing the number of known 32 Origroup members by a factor of three and bringing the total number of identified members to
46, spanning spectral types B5 to L1. We also identify the lithium depletion boundary (LDB)of the group, i.e. the luminosity at which lithium remains unburnt in a coeval population. We
estimate the age of the 32 Ori group independently using both isochronal fitting and LDB
analyses and find it is essentially coeval with the Pictoris moving group, with an age of
24 ± 4 Myr. Finally, we have also searched for circumstellar disc hosts utilising the AllWISE
catalogue. Although we find no evidence for warm, dusty discs, we identify several stars withexcess emission in the WISE W 4-band at 22 µm. Based on the limited number of W 4 detections we estimate a debris disc fraction of 32+12
32P.Ori
4 theC.
M.group.
Bell et al.
8 per cent for

0

20

40

60

40

20

0
20
µ cos( ) [mas yr 1]

4060• 距離93pc、Local	
  Bubbleの縁に位置する運動星団	
  harbour a dark companion which could be responsible for its deviant velocity .
• 運動星団は1980年代以降発見され、現在100pc以
We hereby present the first large-scale spectroscopic survey
fornew
low-mass members of the 32 Ori group to better under内に年齢〜10	  –	
  200	  theories
Myrのもの〜10星団が知られる	
  
DUCTIONrole in constrainingof starand planet
formation.
Given In Section 2 we disstandits stellar
populationand properties.
their youth and proximity tocuss
Earth,
these moving
groupsprovideour candidate
selection
process.
Section 3 details the medium• of ⇠32	  Oisri	  urrounding the Sun out to a distance
100 pcof-groupはMamajek	
  (2007)が最初の報告	
  opticalspectroscopic
observations
the best available samples toresolution
investigate
the early
evolutionof low- in addition to describto as the ‘Local Bubble’ on account of the relatively
inge.g.
howMamajek
the spectroscopic
to intermediate-mass stars (see& Bell properties
2014). Inwere
addi-determined. In Section 4• and 新メンバー探し	  of the interstellar medium
the accompanying lack
we combine
variousindicatorsof stellarand
youth and group membertion to their obvious importanceregardingstellarastrophysics,-forming regions. Hence the discovery just over threeship to identify
new bona
fidethese
members,
and place these in context
in particular the pre-main-sequence(pre-MS)phase,
groupsvan Leeuwen (2007). Also plotted are pre-MS stars from the cata• stars
UCAC4カタログ：半径10°以内に7E5星→固有運動とby comparing them against the findings of previous surveys. Fiof young, seemingly isolated T-Tauriin close
also represent the most readily
accessible
targets
for direct
imaginglogue of Ducourant et al. (2005) and the 15 known stellar members
nally,
Section5
synthesises
ourstellarcensusresults
todiscussthethe Sun was a watershed moment and precipitated a(and other measurements/studies)ofdusty
circumstellar
discs;
esof the 32 Ori group, compiled from a combination of Hipparcos,
global properties of the 32 Ori group, including its age, circumstel距離から5349星→色等級図で100星＋追補50星	
  ervational effort to characterise the young
population
PPMXL (Röser, Demleitner & Schilbach 2010) and the Fourth U.S.
pecially during the epoch oflarterrestrial
planet(see e.g.
disc frequency
andformation
spatial structure.
Bubble (see e.g. Rucinski & Krautter 1983; de la Reza
Naval Observatory CCD Astrograph Catalog (UCAC4; Zacharias
Canup 2004). Finally, constituent members of these groups are• et al.分光観測@Siding	  Springs	
  ANU2.3m、R=7000：
Gregorio-Hetem et al. 1992; Webb
1999). Our
et al. 2013).
ideal candidates for the discovery and characterisation of young,
While the majority of stars in Fig. 1 have small proper mog of this young population has increased dramatically
substellar objects and, of course,
extrasolar SELECTION
planets (see e.g. La2 CANDIDATE
124+11星について、スペクトル型＋視線速度＋Hα
tions (. 10 mas yr ) and are associated with the more distantarly discoveries with the advent of all-sky
astrometricgrange et al. 2010).
(⇠ 400 pc) Orion OB1 and Orionis associations, there is a clear
Members of a moving group typically cover both a large range
V surveys. Not only have hundreds of additional young
concentration of stars in proper motion space in the vicinity ofof
distances
and
a
large
area
on
the
sky;
unlikecompact
clus輝線とLi吸収線→
新メンバー30星	  entified and spectroscopically characterised within this32 Ori. Within this concentration there are three A-type stars withmore importantly, it has been demonstrated that many ofomprise kinematically distinct, yet spatially dispersed
n which the members share a common motion throughKnown 32 Ori group members

Figure 1. Proper motions of stars within 10 New
of 32
Ori. Proper motions formembers
the O-, B- and A-type stars (including 32 Ori and Bellatrix; see Introduc15 taken from the revised Hipparcos reduction, whereas the pre-MStion) arestars are from the catalogue of Ducourant et al. (2005). Those for the known
members are compiled from a combination of the revised Hipparcos reduction, PPMXL and UCAC4 (see Table 1). The majority of stars have small
10proper motions and are associated with the more distant Orion OB1 andOrionis associations at ⇠ 400 pc, however there is an appreciable concentration of stars with proper motions similar to that of 32 Ori.
5

Declination [Degrees]

2

20

0

5

1

10
6.5

ters whose members essentially share the same proper motion, the

1.1 The 32 Orionis Groupproper motions and tangential velocities of individual group members can vary systematically and significantly. To identify potential

6.0

5.5
Right ascension [Hours]

5.0

4.5

Hipparcos entries (HR 1807 [HIP 25453], HD 36823 [HIP 26161]
FigureSpatial[HIP
distributionof whichthe 15 known
members of
the 32 Ori
andHD3.35714
25483])
are all stellar
approximately
co-distantgroup32
listedTablewhich
1, in addition
the 30 new
members wehavegroup
identiwith
Ori,in andtogetherto provide
a weighted
mean

Olling 2011), and somembers. Table 1 p
tions and radial velo
32 Ori group, whileproper motion, radispectral types listed
HR 1807 have beencalibrated spectra a
grid of MK spectra
depicts the U V W
young groups andother young groups
majek (2016), excep
Interestingly, not onvelocity to the Pic
planes, but the two
et al. 2015 and the dTwo of the 15
TYC 112-917-1) ha
the BPMG by Elliot
Nearby Young Asso2014; Malo et al. 20
membership probab
ties that they are ass(' 90 per cent). Furtion 2.1 for details)
velocity from Mam
BPMG members theof the group, at appof the ‘classic’ mem
also Mamajek & Be
TYC 112-1486-1 anfor this study and ingroup properties list
Shvonski et al.
survey of the 32 Or
covering all bandpa
et al. 2010). In thistometry with opticalemission arising fro
port that 4/14 memb
(f24 = 88.45 ± 0.36
HD 36338 (f24 = 1
3.87 ± 0.09 mJy). H
91.0 ± 4.2 mJy). No
and correspond to ex
tospheric levels. Mo

32 Orionis group10532 Orionis group71
The latter member alsoHehas, butii with large
He i
Li i
PM = 12.5 mas yr He
1
4
(>10
mas
yr
)
errorson
its
UCAC4
propermotion.
2M0537+1329 (M5.5)2
2M0537+1329(M5.5)
Given the incompletenessof APASS BV gri photometry in
UCAC4,
we also searched for additional low-mass members in the
8
MRUCAC , RUCAC
Ks CMD using the stars selected in Fig. 4
2M0524+0640 (M5)4(M5)identified ⇠30 additional candidates with conas a 2M0524+0640
reference, andsistent CMD positions not present in the MV , V
Ks diagram.
Furthermore, we also looked for potential members in the First
2M0526+0712 (M4.5)
6
6(M4.5)
U.S.2M0526+0712
Naval Observatory
Robotic Astrometric Telescope Catalog(URAT1; Zacharias et al. 2015), which extends approximately
2 mag fainter than UCAC4 in an almost identical non-standard
2M0519+1038 (M3)
pseudo-R
bandpass.
2M0519+1038(M3) We used the same method as above, but al8
41
lowed a PM < 15 mas yr limit due to the two-epoch astrometry having larger proper motion uncertainties. From a similar kine2M0533+1400 (M2)
2M0533+1400
(M2)
matic
and CMD analysis
we identified another ⇠20 candidates only10
present in URAT1 down to R ⇡ 16 mag, giving a total of ⇠150
2
potential
members from the combined searches.
2M0535+1354 (M1.5)
2M0535+1354
The fainter (M1.5)
RUCAC and URAT1 samples are crucial for better
12
estimating the age of the group. The limiting magnitude of APASSis V ' 16 mag, which at a distance of 90 pc, Hcorresponds
to the H
Li i
Li i
0
expected
the lithium
depletion(LDB)
in a6550
5400 location
5600 of
5800
6000
6200
6400boundary
6600
6800
6600
6650
6700
6750
6800
0
1
2
3
4
5
6
⇠25 Myr population. Only by confirming fainter Li-richWavelength
members[Å]V –Ks [mag]32 Orionis group 21
can we can identify
the precise
32 Orionis
group position
19 of the 32 Ori group LDB
and calculate a semi-fundamental (Soderblom et al. 2014) age for
Figure 5. WiFeS/R7000
spectra for a selection of new M dwarf 32 Ori
1.0 group members identified in this study. Fluxes are normalised around 6100 Å and key
Figure 74. UCAC4/APASS DR6 MV , V
Ks colour-magnitude
The diagram.
obvious dataset
which
to compare
censusthe against
group
Section5.3). Å our
10 8 (see10youth indicators
are labelled. Li I 6708HD
absorption decreases in strength with decreasing effective temperature through the early-M spectral types, before
The circles denote potential kinematic
stars
are
known
HR 1807 members
2M0528+0113
36338
10 and theof circumstellar discs is thereturning
recently
publishedlevels
Spitzersurvey10
to undepleted
below
the6 ⇠ 25by
Myr lithium depletion boundary (LDB) at around M4.5 (see also Fig. 7).
THOR
2
4B
THOR
32 Origroup members in UCAC4
which
satisfy the same kinematic se-THOR sharply
10 8
Shvonski et al. (2016). Our10 WISE
analysis corroborates that of
9lection criteria (see Section 2.1). The solid line is the empirical fit to the
Shvonski et al. with neither study finding evidence for the presence
11
Pleiadeswhich has been cor10 (2007)
10 9 single-star sequence by Stauffer et al.
3 SPECTROSCOPICOBSERVATIONS
0.8
base of the convective zone in late-K/early-M dwarfs. For stars beof full or warm dusty discs
based
on a lack of excess
detections AND DATA
1010rected for the effects of both distance and interstellar extinction.
Fe i
Fe i Ni i
Fe i Fe i
Li i Ca i
Fe i
Ni i
ANALYSIS
tween these luminosities, rapid Li depletion ensues and the Li I
in
all
four
IRAC
bands
as
well
as
the
WISE
W
1
and
W
2
bands.
10
1.2
10
feature is no longer visible. At an age of ⇠ 25 Myr we expect Li
12
Of the four debris disc candidates
identified differentiate
by Shvonski etal., we genuine 32 Ori group10
To 10
unambiguously
between11
to be fully depleted in stars with spectral types between ⇠ M1.5
2016 Jan 24
10 11
find that both HR 1807members(THOR 2)andandfield
HD interlopers
36338 (THOR
6) arecontaminants, additionalor otherand M4.5, however it should still be present in both earlier and0.6
22 µm excess
sources, however
we do diagnostics
not detect any
emis-These typically include
70 < 12dkin < 110 pc). During the spectroscopic
follow-up
we also
spectroscopic
areexcessrequired.
later spectral types (see e.g. Mentuch et al. 2008; Jeffries et al.
10
10 13
10 12
0
1
0
1
0
1
10 outside
10 35499
10 TYC-112-1486-1
1.1
sion
from either(THOR radial
5) or
(THOR
observed stars10whose distance10and PM values felltheseHDmeasuringvelocities to10ascertain
whether candidates Cha
share2013). As in the case of H↵, we determined EW[Li] by Gaus2016well
Feb 21as identifyal. determining
both exhibitspace
24 µm
emis- aslimitsotherwise resembled strong candidates. 10), despite Shvonski etsimilar
9three-dimensionalmotions,
NGCprofile fitting if the line was present and estimated an upper
10 but
UCL/sian
10
2547
10
LCCLi
sion at > 4 above typicalphotospheric values.
The
cause of withthis stellar youth (e.g.
limit
the local pseudo-continuum if it was not. No attempt
ing
spectroscopic
featuresassociatedI fromTuc-Hor/
2M0527+1446
2M0527+0645
2M0524+0640
0.4
discrepancy
isunclearasneither
objectis
flaggedasvariableinthewasmadeColumba
to deblend Li I from the nearby Fe I feature at 6707.4 Å,
10
THOR 17A
THOR 27 6708 Å absorption and H↵ emission).
THOR 38101010which is unresolved at our resolution and is typically much weaker
AllWISE catalogue (but1.0see We
Melis
et al. 2012)and visualmembers
inspec- (plus2016 Oct
13 TWAobserved
124 candidate
11 of
the known
32 Ori
10 11(EW[Fe]. 30 mÅ; Soderblom et al. 1993) compared to the extion of the co-added images
suggestnothing
unusual(e.g.
nearbymembers
listed
in
Bell
et
al.
2015)
in
four
runs
between2015
2.2 Colour-magnitudeselection
10 11pected EW[Li] of 32 Ori group members. WiFeS EW[H↵] and10 11sources, non-uniform background).
In
the
case
of
TYC
112-1486-1
September and 2016 February using the Wide Field Spectrograph
BPMG values
IC 2391
0.2
EW[Li]
for all observed candidate and known members of
The kinematic distances determined in Section 2.1
2017
Jan 07 at Siding
this must
couldalso
be abesensitivity
issue
as
the signal-to-noiseof the(WiFeS;
Dopitaet al. 2007) on ratio
the ANU
2.3-m
telescope
10 12
the
32
Ori
group
are listed in Table 3.
12
consistentwiththecolour-magnitude
diagram(CMD)
positionsof10
W 4-band measurement 0.9
is only Observatory⇠3. The remaining
debris
10 12
Spring(SSO). four
Poorerthandisc
expected conditions meant
Per
Pleiadesyoung, low-mass stars above the main-sequence. Our
primary
phocandidates identified in we
thiswerestudy
are all
new members
andfaintest
thus (V > 15 mag) candiunable
to observe
all of thetometric
selection was performed using theUCAC4/APASSDR6 in thedates
were not included
Shvonskiet we
al. study.
13
prioritised
objects
with small 6750
PM values, CMD
6600 so
6650
6700
10 13
10 13
10 and00
0.0
101 the 3730 (out of 10
101
100
101
MV , V Ks 10
CMD.
Fig. 4 shows5349)
Of potential
the six objects
identified
asexhibiting
Wknown4 excesses,
the and candidates
0
20
40
60
80
100
120
140
160
or sky positions
similar Wavelengthtomembers
coinci[Å]
Wavelength[µm]
kinematic members with APASS V - and 2MASSK-band
phoAge [Myr]four sM dwarf debris disccandidates
are ofparticular
Studdent
with ROSAT
X-ray
sourcesnote.
(Vogeset al. 1999). We later took
tometry, in addition to known members of the 32 ies
Orihavegroupdemonstrated
which
that
as stellar inmass
so does
observations
2016decreases,
October and2017theJanuary to revisit several
6. excess
Li I-region
spectra at
of 22
the
new
SB3
and possible memFigure
Spectral
distributionsof six membersof the32 Origroup isforwhich
we haveFigure
identifiedW
4 emissionµm.
Notethatsystemthesatisfy15.thesameenergy
kinematic
selection
criteria.
Alsoplottedemfraction
ofthe
associateddebris
discs
(see
e.g.
Lestradeetal.
2009).
possible
spectroscopic
binaries
identifiedin earlier
ber HD
37825(spectral
type
F5),
showing
clear evolutionof Figureitsruns.
neutraluncertainties
ontheindividual
pointsare
typicallysmallerthanthesymbols
used.
In
all
cases
except
HR
1807,
for
which
we
used
the
ATLAS
9
ODFNEW
16. Debris disc fraction (as estimated from excess emission at either
pirical fit to the ⇠ 130 Myr-old Pleiades single-starsequence
by M dwarf WiFeS
To date,only three
discs
confirmed
via
is Theanhaveimage-slicing
integralfield spectrographwith time.
2016been
February
spectrumis single-lined(RV⇡ µm) withmodels, the BT-Settl CIFIST models have been adopted to compare against
the photometricdata. metal linesdebris
22
or 24
as a function of age for young nearby groups/associations
1
Stauffer et al. (2007), corrected for the effects of bothdistanceandscattered light observations,withsimilarto 2017or younger
28 akmnominals all), while
inages2016arcsec
Octoberfield-of-view
andJanuary
it is0.5
resolved
intosampling
25⇥38
and
arcsecandclusters.
The disc fractions for all other regions have been taken fromextinction (d = 136 pc and AV = 0.12 mag fromthan
Melistheet al.
threecomponents(only 7twoof whichare25;
visibletheCCF). The twowere
main made in
32 2014
Ori group
(AU
Mic,
TWA
and
TWA
seeinChoalong
twenty-five38⇥1
arcsecslitlets.
Observations

22/24 µm excess fraction

Normalised flux + const.

Normalised flux + const.

F [erg s 1 cm 2 ]F [erg s 1 cm 2 ]MV [mag]

Known 32 Ori group members
Kinematic candidatesPleiades sequence16

C. P. M. Bell et al.
2

9Known members
New Li-poor members
New Li-rich members
Binaries

0

10

2

11
MV [mag]

MV [mag]

4
6

12

8
10

13
12
14

1

0

1

2

3
V –Ks [mag]

4

5

14
4.0

6

LDB position
Dartmouth LDB 18 Myr
Dartmouth LDB 22 Myr
Dartmouth LDB 26 Myr

4.5

5.0
5.5
V –Ks [mag]

6.0

6.5

Figure 13. Left: Example best-fit MV , V
Ks CMD of the 32 Ori group with the Dartmouth models overlaid. The grey points represent individual stars
in our two-dimensional probability distribution (see text) from which we derive a best-fit age of 23 Myr. The lithium depletion boundary (LDB) is clearly
identified at V
Ks ' 5.5 mag (spectral type M4.5). Right: The position of the 32 Ori group LDB as defined in this study (see text). Overlaid are several
lines corresponding to loci of constant luminosity at which Li is depleted at the 99 per cent level as predicted by the Dartmouth evolutionary models. The
best-fit LDB age of 22 Myr in this panel is in excellent agreement with the isochronal age based on the same stellar evolutionary models.

Table 9. Isochronal and LDB ages for the 32 Ori group.
MethodIsochrone

LDB

Final adopted age

Model

Age (Myr)

Mean age (Myr)

Dartmouth
BHAC15
Pisa
PARSEC

23+73
28+34
20+41

25 ± 5 (1 )[±4 (statistical), ±3 (systematic)]

29 ± 3Dartmouth
BHAC15
Pisa
Dartmouth (hBf i = 2.5 kG)

22+43
22+43
21 ± 326 ± 423 ± 4 (1 )[±4 (statistical), ±1 (systematic)]

24 ± 4 (1 )[±4 (statistical), ±2 (systematic)]

a manner is entirely justified (see also Jeffries & Oliveira 2005;
Tognelli et al. 2015).

magnetic Dartmouth models imply an older, yet consistent, age forthe group. These findings are similar to those reported by Maloet al. (2014) and Binks & Jeffries (2016b) for the BPMG.

1

v:1703.02745v1 [astro-ph.GA] 8 Mar 2017

University of Kiel, Institute of Theoretical Physics and Astrophysics, Leibnizstrasse 15, 24118 Kiel, Germany
e-mail: [rbrauer;wolf]@astrophysik.uni-kiel.de
2
University of Heidelberg, Institute of Theoretical Astrophysics, Albert-Ueberle-Str. 2 U04, 69120 Heidelberg, Germany
e-mail: reissl@uni-heidelberg.deMarch 9, 2017
ABSTRACT
Context. Observations of Zeeman split spectral lines represent an important approach to derive the structure and strength of magnetic

fields in molecular clouds. In contrast to the uncertainty of the spectral line observation itself, the uncertainty of the analysis method
to derive the magnetic field strength from these observations is not been well characterized so far.
Aims. We investigate the impact of several physical quantities on the uncertainty of the analysis method, which is used to derive theline-of-sight (LOS) magnetic field strength from Zeeman split spectral lines. These quantities are the density, temperature, velocity,and the magnetic field strength.
Methods. We simulate the Zeeman splitting of the 1665 MHz OH line with the 3D radiative transfer (RT) extension ZRAD. Thisextension is based on the line RT code Mol3D (Ober et al. 2015) and has been developed for the POLArized RadIation Simulator
POLARIS (Reissl et al. 2016).
Results. Observations of the OH Zeeman e↵ect in typical molecular clouds are not significantly a↵ected by the uncertainty of theanalysis method. However, some observations obtained a magnetic field strength of more than ⇠300 µG, which may result in an
uncertainty of the analysis method of >10%. We derived an approximation to quantify the range of parameters in which the analysis method works sufficiently accurate and provide factors to convert our results to other spectral lines and species as well. We
applied these conversion factors to CN and found that observations of the CN Zeeman e↵ect in typical molecular clouds are neither
significantly a↵ected by the uncertainty of the analysis method. In addition, we found that the density has almost no impact on theuncertainty of the analysis method, unless it reaches values higher than those typically found in molecular clouds (nH
107 cm 3 ).
Furthermore, the uncertainty of the analysis method increases, if both the gas velocity and the magnetic field show significant variations along the line-of-sight. However, this increase should be small in Zeeman observations of most molecular clouds considering
typical velocities of ⇠1 km/s.
Key words. ISM: clouds – Line: profiles – Magnetic fields – Polarization – Radiative transfer

1. Introduction
The impact of magnetic fields on the formation of stars andplanets is a matter of ongoing discussions (e.g., Matthews &Wilson 2002; Pudritz et al. 2014; Seifried & Walch 2015). For
instance, a comprehensive understanding of the relative importance of magnetic fields and turbulence in driving the formationof molecular cores has not been achieved so far (Crutcher et al.

(Crutcher 2004; Cho & Yoo 2016). On the other hand, direct
measurements are performed by observing spectral lines thatare split up by the Zeeman e↵ect (Crutcher 2012). The uncertainty of the magnetic field strength derived from these observations range from ⇠5% to a few 100% (Crutcher et al. 2010b).
With the sensitivity and spectral resolution of currently operating instruments, only the LOS component of the magnetic fieldstrength can be obtained from Zeeman observations (Crutcher

Multi-epoch, high spatial resolution observations of multiple T
Tauri systems
Gergely Csépány1, 2 , Mario van den Ancker1 , Péter Ábrahám2 , Rainer Köhler4, 5 , Wolfgang Brandner3 ,Felix Hormuth3 , and Hector Hiss3
1

European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germany
Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, H-1121 Budapest,
Konkoly Thege Miklós út 15-17, Hungary
3Max-Planck-Institut für Astronomie, Königstuhl 17, 69117 Heidelberg, Germany
4
Institut für Astro- und Teilchenphysik, Universität Innsbruck, Technikerstr. 25/8, 6020 Innsbruck, Austria
5
University of Vienna, Department of Astrophysics, Türkenschanzstr. 17 (Sternwarte), 1180 Vienna, Austria

arXiv:1702.08583v1 [astro-ph.SR] 28 Feb 20172

Received / Accepted
ABSTRACT
Context. In multiple pre-main-sequence systems the lifetime of circumstellar disks appears to be shorter than around single stars, andthe actual dissipation process may depend on the binary parameters of the systems.
Aims. We report high spatial resolution observations of multiple T Tauri systems at optical and infrared wavelengths. We determine if
the components are gravitationally bound and orbital motion is visible, derive orbital parameters and investigate possible correlations
between the binary parameters and disk states.
Methods. We selected 18 T Tau multiple systems (16 binary and two triple systems, yielding 16 + 2 ⇥ 2 = 20 binary pairs) in the
Taurus-Auriga star forming region from the survey by Leinert et al. (1993), with spectral types from K1 to M5 and separations from
0.2200 (31 AU) to 5.800 (814 AU). We analysed data acquired in 2006–07 at Calar Alto using the AstraLux lucky imaging system,
along with data from SPHERE and NACO at the VLT, and from the literature.
Results. We found ten pairs to orbit each other, five pairs that may show orbital motion and five likely common proper motion pairs.
We found no obvious correlation between the stellar parameters and binary configuration. The 10 µm infra-red excess varies between
0.1 and 7.2 magnitudes (similar to the distribution in single stars, where it is between 1.7 and 9.1), implying that the presence of thebinary star does not greatly influence the emission from the inner disk.
Conclusions. We have detected orbital motion in young T Tauri systems over a timescale of ⇡ 20 years. Further observations witheven longer temporal baseline will provide crucial information on the dynamics of these young stellar systems.

PMS多重星のまわりの円盤	  • 重力的に束縛されているか？ 軌道運動は見えるか？ 軌
道パラメタ → 円盤の状態との相関は？	
  1. Introduction
multiple systems) the sub-millimeter flux densities (and thus thedisk masses) are lower for binaries with a projected semi-major• 可視・赤外の観測	  
The influence of the multiplicity on the evolution of the proto- axes < 100 AU. Cieza et al. (2009) also found that circumstelreduced in binaries with separation lessplanetary disk is •an openquestion of1
astronomy(see e.g. Bouw- lar disk lifetimes-­‐are
Tau-­‐Aur	
  8天体、K1-­‐K5、0.22”(31AU)	
  	
  5.8”(814AU)	
  man et al. 2006; Kraus et al. 2012). Multiple systems are present than 10–100 AU. Although the disks of the components are usuabundantly amongstars, Aaslto	  a largefractionstars form ally assumed to be coeval, they may evolve with di↵erent pace.
• youngCalar	
  
2.2m	
  
+	
  ofVLT	  in binary or multiple systems, and, for example, Lada (2006) has This di↵erential disk dispersion is in the focus of recent studies• the結果	
  found thatsingle star fraction is only ⇡40% for G-type stars, (Patience et al. 2008; Kraus & Hillenbrand 2009; Cieza et al.
rising to 70% for late M-type stars. This shows that multiplicity 2009; Kraus et al. 2012; Daemgen et al. 2012a,b, 2013). How• We	
  found	  no	
  obvious	
  correlamon	
  between	
  the	
  stellar	
  might be a significant factor in the stellar evolution. The young ever, although these studies derived some correlations between
the presence of a	
   disk and binarity, they did not perform any
multiple systems are of
special interest, because
low-massparameters	
  and	
  btheinary	
  conﬁguramon.	
  evolutionary models at the early phases of stellar evolution are detailed analysis of disk content and structure. Simulations comand X-ray iphotoevaporation• constrainedthe	
  presence	
  of	
  the	
  inary	
  sbiningtar	
  dviscousoes	
  drainingnot	
  greatly	
  nﬂuence	
   suggest thatcurrently only poorly
by observations,andbobserving such systems can help to refine and calibrate those models in systems with separation below 100 AU it is the secondary disk
disk.	
  
	
  
which
disappears
first, while for wider systems this trend cannot(Stassun et al. 2009). the	  emission	
  from	  the	
  inner	  •Key words. Stars: variables: T Tauri, Herbig Ae/Be - stars: pre-main sequence - stars: evolution - binaries: visual - techniques: highangular resolution - instrumentation: adaptive optics

There is evidence that the age at which the star still has adisk, is shorter in multiple systems than around single stars (seebe seen (Reipurth et al. 2014, and references therein). The actual disk removal process that causes the relatively short disk
lifespan and the di↵erential disk dispersion is unclear. The disk

A&A proofs: manuscript no. ttbs_rev70
Table 3. Parameters of the fitted orbits. Eccentricity is zero for all andthe 2⌫ is listed in Table 5.

Binary paira (00 ) a (AU) period (yr) ⌦ ( )i( )LkCa 35.8
817> 10,000
161
95
7.4
1041
> 10,000
84
94
DD Tau
LkCa 7
1.1
157
575
27
94
FV Tau3.6
505> 10,000
4982
FV Tau/c
0.7
103
316
113
86
UX Tau AC
2.7
383
692
2
90
UX Tau AB 11.0
1542
> 10,000
100
84
FX Tau
0.9
132
417
114
100
DK Tau4.9
683> 10,000
10995
XZ Tau
0.3
42
182
143
116
HK Tau8.8
1233> 10,000
88
105
V710 Tau
4.1
573> 10,00038130UZ Tau AB
0.8
110
1202
87
62
UZ Tau AC
5.7796> 10,000
165
53GH Tau
0.3
48
575
145
138
HN Tau11.5
1604> 10,000
5187
HV Tau
4.0
565
501
44
92
V999 Tau
13.0
1826> 10,000
1889
V1000 Tau
0.2
34
138
98
60
RW Aur
1.5
206
955
74
94
a: semi-major axis (using the distance to Taurus star forming
region, see Sec. 2.2), ⌦: position angle of node, i: inclination
SED can tell us about the extra or extended emission in the systems, which can be signs of the dust in the disks. Since for our
purposes the most interesting part of the spectrum is between
the optical and millimeter wavelength, we looked for photometric measurement in this region. We employed many photometricsurveys: the measurements from Kenyon & Hartmann (1995);
Howard et al. (2013) and the IRAS, Spitzer, 2MASS, WISE allsky surveys, a few ALMA and HST observations along with our
lucky imaging data.
The magnitude data were converted to F using the conversion formulas from Johnson (1966); Bessell (1979); Glasset al. (1982); Berrilli et al. (1992); Gehrz et al. (1974); Helou& Walker (1988), while the data from the surveys were either already given in Jy or had their own conversion formulas to obtain
F (e.g. WISE).
We fitted a SED curve on all observations, using the
NextGen2 atmospheric models from Hauschildt et al. (1999)7
and the extinction formulae from Cardelli et al. (1989). We

using a grid search. We have derived the uncertainties of the extinction magnitude using a Monte-Carlo approach by randomly
varying the measured flux values within 1 and recording theextinction magnitudes of the di↵erent runs, repeated a few hundred times. We employed a “combined flux” fitting procedure
where we took into account both the individual flux of the components and the combined flux of the whole system at the sametime.
The obtained extinction magnitudes are listed in Table 4, andplots of the fluxes and SED curves are shown in Fig. 2. The fitted
extinctions agree with the literature data within uncertainties formany systems, but there are exceptions. In a few systems, only
the extinction of the primary is in agreement with the literature(FV Tau, FV Tau/c, HK Tau, V710 Tau, GH Tau and HV Tau),in these cases we anticipate that the discrepancy is attributed to
the low number of resolved observations to fit (FV Tau/c, HK
Tau, V710 Tau, GH Tau) or that our optical measurement showsa higher extinction for the secondary (in the case of FV Tau)or that we know that the secondary probably has a high extinction magnitude due to an edge-on disk (in the case of HV Tau),
where the companion may be seen through their disks. There
are five systems where the obtained extinction magnitudes do
not agree with the literature (LkCa 3, XZ Tau, HN Tau, V999
Tau, RW Aur), where we know that two of them show significant variability in time (XZ Tau and RW Aur, see the notes inAppendix A), one system that is a known quadruple system consisting of two spectroscopic binary systems (LkCa 3, Torres et al.
2013), and there are two systems where the optical fluxes may
be the cause of discrepancy (HN Tau and V999 Tau). The latter
case, the V999 Tau simply misses reliable optical measurements(although it was observed in the SDSS, it is flagged as a not clean
observation), and HN Tau features a fairly flat SED in the opticalthat we cannot achieve a reliable fit using the NextGen2 models.
The NextGen2 models only provide the photospheric SED
curve up to 2.5 µm, therefore to measure the infra-red excess at
higher wavelengths, we attached a black body radiation curve tothe K-band section of the model curve (with the same e↵ective
temperature that we used in the NextGen2 model), making the
SED available in the whole wavelength range.

5. Results
5.1. Astrometry

The individual epochs of the system are shown in Figure 3 in