THE STAR FORMATIONNEWSLETTER
No. 269 (#1-5)
塚越崇(茨城大)5 Jun. 2015

1. Planet formation around binary stars: Tatooine made easy
2. The inner environment of ZCMa: High-Contrast Imaging Polarimetry
with NaCo
3. The infrared emission of young HII regions: A Herschel/Hi-GAL
study
4. Early evolution of embedded clusters
5. Dust trapping by spiral arms in gravitationally unstable protostellar
discsPlanet formation around binary stars: Tatooine made easy
B.C.Bromly and S.J.Kenyon

of PA = 8 .8 ± 1 .5 when analyzing previous observations cov1
ering a 11.2 yr Imaging
baseline, derivinga PA rate ofwith⇠ 0 .7 yr
for a
The inner environment of ZCMa: High-Contrast
Polarimetry
NaCo
circular orbit composed of a 1M and a 5M stars.

H.Canovas et al.

• ZCMa (d~930pc)•

Herbig star(P)とFUOri(S)の連星•

jet 2本(長短)•

変光現象あり•質量降着によるもの•

jetによってenvelopeに穴が開く

clu
bri
ha
ne
sta
an
pri
lar
the

• VLT/NaCoを使用•

H, Ksバンド

PDI観測• 両バンドで連星を分解

Fig. 1. Top and bottom-left: Intensity images in logarithmic scale. AO
artifacts are visible in the Z CMa images. HD 52841 is slightly elongated in the SE-NW direction. Bottom-right: The PI image of HD 52841
shows the noise that inevitably appears in the innermost 000. 1 (indicatedby the white circle) of the image as a result of imperfect speckle suppression.

• 2001年データと比べると星の位置が動いていた•~0.7deg/yr

pa
S/N
sq
aro
en
art
fro
en
po
lik
fac
aro
In

Our polarized images (Fig. 2) show three complex features

5.

Ou
me
is
(20
(w
int
co

A&A proofs: manuscript no. zcma_letter

Fig. 2. Right and center: PI images at H and K s bands. Gr
scale, using square root scaling. The pixels falling in the nof large and micro jets, respectively (Whelan et al. 2010).
extended feature, respectively. Center: Arrows “a” and “d”
contours the Keck J band image presented in Millan-Gabecross indicates the brightest point of feature “2” at K s band

Fig. 2. Right and center: PI images at H and K s bands. Green contours indicate the position of the two stars. The two images are plot in the same
scale, using square root scaling. The pixels falling in the non-linear regime are set to zero. Left: the dotted and dashed lines indicate the positionof large and micro jets, respectively (Whelan et al. 2010). Arrows “1”, “2” and “3” indicate the polarized halo, the bright clump, and the sharp
extended feature, respectively. Center: Arrows “a” and “d” shows an artifact and a local decrement in polarized flux, respectively. The white line
contours 星の周りのハロー
the Keck J band image presented in Millan-Gabet & Monnier (2002). Right: Relative polarized color computed as [H] - [K s ]. The whitecross indicates the brightest point of feature “2” at K s band.

• 3つのコンポーネント••

東西に広がったクランプ••

embedded very young star?

南西方向に伸びた構造

•

要因は不明

•

ダストの性質は星周と違うかも• Results show NaCo’s imaging capabilities

Fig. 3. Left: P✓ over polarized intensity at K s -band. The lehighlight faint features. AO artifacts are indicated by the aThe infrared emission of young HII regions: A Herschel/Hi-GAL
study
R. Cesaroni et al.: Infrared emission of young HII
R.Cesaron et al.

• Herschel+VLAを使ってHII領域を調べ大質量星形成を探る• 使用データ/解析4•

Herschel Hi-GAL survey•

VLA CORNISH 5GHz survey•

SEDを作成し光度を計算•

Spitzer, WISE, MSXなども使用•

2温度成分のmodified blackbodyでフィット• ~200個くらいのHII領域のサンプル•

クランプ質量 >~100Msun

R. Cesaroni et al.: Infrared emission of young HIITable 1. Steps of the source selection process.
CORNISH
HIIs
281

→

HIIs separation<11.′′ 5244→with Hi-GAL
counterparts
230→

Hi-GAL cwith ≥
2

Fig. 4. Distribution of the clump mass associated with CORNISH HII
regions. The solid and dotted histograms are obtained by choosing, respectively, the far and near distances for the 11 targets for which the
KDA could not be solved.

The distribution of the clump masses, shown in Fig. 4,
demonstrates that the vast majority of the clumps are above several 100 M⊙ , as expected for high-mass star forming regions.
It is thus very likely that our measurements refer to rich stellar
clusters, tightly associated with the early-type stars ionizing the
CORNISH HII regions.

4. Discussion
4.1. Lyman continuum excess•

ルミノシティ >~104Lsun

The large luminosities and masses of the objects under study arestrongly suggestive of the presence of multiple stars. In order toestablish the contribution of low-mass stars and investigate theproperties of the early-type stars, we have estimated their Lyman

Fig. 5
gions
Hi-G
matic
for u
no K
tance
NLy –
90%
ber o
curve

main
cont
fecti
ertie
Thom
The
Lym
a ZA
it is

• ZAMSから期待されるLyman continuumより超過したものが存在• Lyman excess天体は光度が小さめ=> 比較的若く埋もれている？

• 解釈• 円盤への降着によるショック

ces
and

Fig. 10. Number of HII regions per unit surface as a function of
Galactocentric distance for sources with Lyman excess (dashed histogram, numbers on right axis) and without (solid histogram, numberson left axis).
8[Hosokawa&Omukai2009]• 星のhot spotへの降着[Smith2014]

R. Cesaroni et al.: Infrared emission of young HII regions: A Herschel/Hi-GAL study

Fig. 7. Analogous to Fig. 5, where the bolometric luminosity computedfrom the Hi-GAL data has been replaced by that estimated from the
IRAS fluxes.

Of course, we cannot rule out the possibility that a few objects in the forbidden zone are aﬀected by an inappropriate estimate of L. In particular, some of the sources above the blackbody curve appear too extreme not to be misplaced. In Fig. 8 wecompare the full width at half power (FWHP) at 70 µm of thesesources with that of the other objects in the forbidden region
lying below the blackbody curve. Clearly, the former are more
extended than the latter and this could make it more diﬃcult toestimate their integrated flux with CuTEx (see Sect. 3.1), which
was conceived to identify compact sources. Consequently, it isplausible that the luminosity might have been underestimated forsome of these sources.
In conclusion, we believe that Fig. 5 proves that the• 星形成効率は5%程度であり、Lyman
excessの有無で違いはない
10
R. Cesaroni et al.: Infrared emission of young HII regions: A Herschel/Hi-GAL study
Fig. 15. Ratio between the HII region size (from the CORNISH catalogue) and the FWHP at 500 µm of the Hi-GAL counterpart, versus
the ratio between the Lyman continuum photon rate and the bolometricluminosity. The solid curve was obtained by rebinning the data points.

• ほとんどの天体がsupervirial

star is relatively compact, which generates high free-fall speeds,
while the free-fall onto the limited area significantly raises thetemperature and hence the Lyman flux. In fact, a comparison
between Figs. 15 and 16 of Smith (2014) and our Fig. 5 demonstrates the ability of the model to explain most of the Lymanexcess sources.
More detailed numerical calculations are needed in order tocome to a firm conclusion on this issue, since a substantial improvement is necessary to predict the exact amount of Lyman
continuum photons emitted in the accretion process.
4.3. Star formation efficiency

Fig. 16. Ratio between the Lyman continuum photon rate and the bolometric luminosity, versus the ratio between the luminosity and theclump mass for the CORNISH HII regions. The red points indicate
sources with Lyman excess. The hatched area contains 90% of the simulated clusters assuming a star formation eﬃciency of 5%.

If our sample is indeed homogeneous, the observed distributions of mass, luminosity, and Lyman continuum emission
should mirror the variation of the stellar cluster characteristics
rather than a large spread in age. This can be verified by studying the ratios L/Mgas and NLy /L. Both should increase with time
during the process of high-mass star formation because stars
gain mass to the detriment of the surrounding envelope andthus increase their luminosities and Lyman continuum fluxes.
Consequently, if our sources spanned a large age interval, one
should observe a correlation between NLy /L and L/Mgas .
In Fig. 16 we plot these two ratios against each other for allof the 204 objects of our sample. We note that we have also included the four sources without VLSR information because theL/Mgas and NLy /L do not depend on the distance. This feature

Early evolution of embedded clusters
J.E.Dale et al.
ABSTRACT

We examine the combined eﬀects of winds and photoionizing radiation from O–typestars on embedded stellar clusters formed in model turbulent molecular clouds
covering a range of masses and radii. We find that feedback is able to increase thequantities of dense gas present, but decreases the rate and eﬃciency of the conversionof gas to stars relative to control simulations in which feedback is absent. Star
formation in these calculations often proceeds at a rate substantially slower than
the freefall rate in the dense gas. This decoupling is due to the weakening of, andexpulsion of gas from, the deepest parts of the clouds’ potential wells where most of
the star formation occurs in the control simulations. This results in large fractionsof the stellar populations in the feedback simulation becoming dissociated fromdense gas. However, where star formation does occur in both control and feedback
simulations, it does so in dense gas, so the correlation between star formation activity
and dense gas is preserved.
The overall dynamical eﬀects of feedback on the clusters are minimal, with
only small fraction of stars becoming unbound, despite large quantities of gas being
expelled from some clouds. This owes to the settling of the stars into virialised andstellar–dominated configurations before the onset of feedback. By contrast, the eﬀects
of feedback on the observable properties of the clusters – their U–, B– and V–band
magnitudes – are strong and sudden. The timescales on which the clusters become
visible and unobscured are short compared with the timescales which the clouds are
actually destroyed.

(a) Run I

RunMass (M⊙ )R0 (pc)⟨n(H2 )⟩ (cm−3 )vRMS,0 (km s−1 )⟨T0 ⟩(K)⟨M0 ⟩tﬀ,0 (Myr)

A

106

180

2.9

5.0

143

6.5

19.6

B

106

95

16

6.9

92

11.2

7.50

X

106

45

149

9.6

52

20.7

2.44

D

105

45

15

3.0

92

4.9

7.70

E

105

21

147

4.6

52

4.9

2.46

F

105

10

1439

6.7

30

19.0

0.81

I

104

10

136

2.1

52

4.6

2.56

J

104

5

1135

3.0

30

8.5

0.90

UZ

106

45

149

18.2

52

39.3

2.9

UB

3 × 105

45

45

10.0

68

18.9

6.0

UC

3 × 105

21

443

14.6

36

37.9

1.9

UV

105

21

148

12.2

52

26.3

3.3

UU

105

10 UP
(d) Run

1371

8.4

26
(e) Run 25.6
UQ

1.1(b) Run J(c) Run UF

4

3 × 10
10
6.7
19.7
2.0
Figure 2. Dense gas mass UF
depletion
times inunits 410of the freefall timeat ρcrit (≈280.35 Myr)in thecontrol (blue) and dual–feedback
UP UQ
104simulations.
2.5 The 9096
27.9 time
0.4 of tﬀ (ρcrit ), which would indicate(black) Runs I, J, UF, UP and
red dashed line7.6corresponds to18a depletionstar formation proceeding atUQ
the freefall
rate 5.0
in gas of1137
this density. 5.4
104
30
14.9
1.2
Table 1. Initial properties of clouds listed in descending order by mass. Columns are the run name, cloud mass, initial radius, initial
RMS turbulent velocity, the initial mean gas temperature, the initial mean turbulent Mach number, and the initial cloud freefall time.

(a) Run I(a) Run I(b) Run J(b) Run J(c) Run UF(c) Run UF(d) Run UP(e) Run UQ

Dust trapping by spiral arms in gravitationally unstable protostellar discs
G.Dipierro et al.

• 重力不安定性円盤でダストの運動にどういう影響があるか？
それらは観測可能か？

• Disk model•

準静的な重力不安定ガス円盤でのダスト運動と進化を考える•

gas disk••

Loadto&Rice(2004)のSPHモデルのスナップショットを使用•

R=150AU, M*=0.25Msun, Md/M*=q=0.25

Dust

dust evolution model (初期状態)••

coagulation-fragmentation steady state
power-law size distribution of dust: n(a)∝a-q•

amin=0.1µm• RADMC-3Dによる放射輸送計算
Figure 1. Gas surface density structure of a 0.25M

disc around•a 1M star that has reached a long-lived, self-gravitating state.
静水圧平衡、ダスト沈殿とvertical turbulentを考慮して計算結果を3Dに•

GAS AND DUST DISC MODELS
H-band (convolution with PSF), ALMA band 6, 7,8, and 9 2(intoCASA)

The simulations presented here calculate the evolution

F
a

in

St = ↵

cool . The solid line shows the fragmentation barrier, i.e. the maximum value of the initial grain size distribution (see Eq. 7).

• St=1となるサイズのダストはarmに溜まる• interarmでのdust-to-gas比はほぼ初期値(1E-2)、armでは1/100程度になっている• fragmentation threshold velocity

Figure 5. Dust density distribution at the beginning (left) and at the end (right) of simulation (t = 10 yrs) along a slice across theを減少させた場合
disc. The simulation starts from an initial size distribution with exponent q = 3.5 and a fragmentation velocity equal to 10 m s . The
31dashedpline represents the particle size corresponding to a Stokes number of unity while the dotted line show the particle size where
St = ↵ cool . The solid line shows the fragmentation barrier, i.e. the maximum value of the initial grain size distribution (see Eq. 7).

• 30->10m/s• ダストが成長せず円盤全体でSt<1
As expected, the dust density distribution for particles with

Figure 4. The upper panels show the surface density at the end of simulation (t = 103 yrs) integrated in size in the range [1µm, 100µm](left) and [1cm, 100cm] (right). The lower panels show the total dust density distribution (left) and the dust-to-gas ratio (right) at theend of simulation.

Stokes number close to unity (see the upper right panel of=>trappingが非効率的になりなまされ
Fig. 4) show the highest concentration in spiral arms. In thiscase, the spiral arms are very thin and display a strong© 0000conRAS, MNRAS 000, 000–000

た分布に
trast between arm and interarm regions, while the structure

shown by smaller particles is broader and essentially follows
the gas structure (see Fig. 1). Note that beyond ⇠ 50 au thedensity of large particles drops because we have crossed thefragmentation threshold. The lower panels of Fig. 4 show thedensity distribution (left) and the dust-to-gas raq=3total->dusttrappingが働きやすいtio (right). It can be noticed that the dust transport process
in a self-gravitating accretion discs quickly modify the radial
distribution of the dust-to-gas ratio. Starting
q=4and->azimuthal
trappingが働きにくいfrom an initial uniform dust-to-gas ratio (10 2 ), dust particles
migrate quickly towards pressure maxima leading to a=>よりなまされた分布になる
creation of a spiral structure in the dust-to-gas ratio. The
local density of particles in interarm regions reaches levels• ダストサイズ分布を変えた場合••

• ALMAではどの波長でもarmが分解される•

ν大でarm間相対強度が大•

光学特性のせい• qを変えた場合•

q=4でinterarmが光学的に厚くなり星の光が吸収され外側の放射が見えなくなる•

Figure 7. ALMA simulated images of the self-gravitating disc model obtained starting from an initial dust grain size distribution withexponent q = 3.5 and a fragmentation velocity equal to 30 m s 1 . Observing frequencies are reported at the top of each image. Contours
are 4,7.6,ALMA
8, 10 ⇥simulated
the corresponding
The whitedisccolourin thefilled ellipse
in from
the lower
left cornerindicates
the size of the
Figureimages of rms
the noise.
self-gravitating
model
obtainedstarting
an initialdust grain
size distribution
withhalf-power
the synthesizedvelocity
beam (seetableexponentq = contour
3.5 and of
a fragmentationequal
to 130form details).
s 1 . Observing frequencies are reported at the top of each image. Contours
are 4, 6, 8, 10 ⇥ the corresponding rms noise. The white colour in the filled ellipse in the lower left corner indicates the size of thehalf-power contour of the synthesized beam (see table 1 for details).

実際τを調べると、q=4では
τ~1, q=3でτ<1がdominantFigure 8. Comparison of ALMA simulated images at 460 GHz of disc models obtained from simulations with an initial dust grain size
distribution with power-law exponent q = 3 (left) and q = 4 (right) and a fragmentation velocity equal to 30 m s 1 . Contours are 4, 6,
8, 10 ⇥ the corresponding rms noise: 0.78 mJy/beam (left) and 0.72 mJy/beam (right). The white colour in the filled ellipse in the lower

distribution with power-law exponent q = 3 (left) and q = 4 (right).

• ダスト成長の指標ßをモデルとシミュレーションで比較=>観測から見積もる事が可能

Figure 10. The parameter computed from (left) the model data and from (right) two simulated ALMA observations at 650 µm a
3 mm of the disc model obtained starting from an initial dust grain size distribution with exponent q = 3 and a fragmentation veloci
equal to 30 m s 1 . The white colour in the filled ellipse in the lower left corner of the right panel indicates the size of the half-pow
contour of the synthesized beam: 0.054 arcsec ⇥ 0.052 arcsec.

12• HiCIAO H-bandの観測

Giovanni Dipierro, Paola Pinilla, Giuseppe Lodato and Leonardo Testi

work (seeEq. Polarized6), a good
approximation
the dust opacity
3.2.2
H-bandscattered lightfor
imagesis a power law: ⌫ / ⌫ , where the parameter is a veryIn this subsection we show synthetic HiCIAO H-band
good indicator
the levelof grain
growth.
thestarttypipolarizedof
intensity
imageof the discmodelWhileobtained
cal valueingfor
ISM
grainsdustis grain= 1.7,
severalwith
observations
ism sizefroman initial
distribution
exponent(e.g. Testi
2003, Nattaet al.
2004,
q = et3.5al.
and2001,
a fragmentation
velocityequal
to 30Ricci
m s 1et. Asal.
previously
intensity
images essen2010) havefound mentioned,
that disc the< polarizedis naturallyinterism , whichtiallytraceofthegas disc
structuresmall
dust grains
preted in
terms
grain
growth(e.g. since
Draine2006).
In theare less subject to the radial drift. The calculated scatteredcontext of the present work, we calculate the value of thelight image convolved with the measured HiCIAO point
parameter
by function
computingthe inratio
spreadis shown
Fig. of11.theTheALMA
H-band simulated
image reimages at
twoandiﬀerent
frequencies,
the theaimfarto(top)inferandtheveals
asymmetryin brightnesswith
between(low)
This is dueof
to large
forwardscatteringofdiﬀerentnear
levelof discdustregions.
concentration
grainsin armstarlight
by largeThe
dust particles
in a non-face-on
disc (incliand interarm
regions.
parameter
can be obtainedas:

mm (right). We consider the disc model obtained startin
from an initial dust grain size distribution with expone
q = 3 and a fragmentation velocity equal to 30 m s 1 . A
wavelengths mentioned above we verify that the disc is op
cally thin and the dust temperature is high enough to ensu
that the disc is in the Rayleigh-Jeans limit. Therefore, tvariations of
can be interpreted as a sign of dust traping regions. From Fig. 10 it can be noticed that the beindex map has a shape that resemble the spiral structu
of the disc model. The spiral arms, where most of the cetimetre sized grains are concentrated due to the radial dri
are characterized by a lower value of . By contrast, thenation of 30°). In addition, the outer ring-like structure isparameter is relatively higher in the interarm regions, wheF⌫1 ln Fvery bright since lnof⌫2smaller particles increases= the density
2,(15)
large grains are depleted due to the dust trapping mechwith radius due tolnthe
lower
value
of
the
fragmentation⌫1 ln ⌫2nism. Thus, taking into account all the assumptions abo
threshold size in outer disc regions (see Fig. 3). Moreover,where FtheF⌫2 arebrightness
the fluxes
at outerfrequencies⌫1 and⌫2 ,⌫1 anddust grainproperties, theunstable
dust trapping
significativeof thering is also
causedDust thetrapping
in gravitationallydiscs induced
13 by G
respectively.
we showthemapdeterminedofcomputedby the In
factFig.
that 10
scatteredlight ismostlyby theproducesa migrationof large
dust intensity
grains image
into spiral arm
Figure11. Synthetic
HiCIAO H-bandpolarizedof the (left)
disc surface
that, dueto the flaringof thefrom thecurvaturemodel data
and theparameter
calculatedof thenon-face-on(inclination of 30°)
obtained starting
regionthat disccanmodelbe potentiallyinvestigatedthrough mul
disc,
is very enhanced
the outer regions.
to 3from an initial dust grain size distribution with exponent q = 3.5
from twosimulated
ALMAin observationsat However,
650 µmdue
andwavelenght
observations.
and a fragmentation velocity equal to 30 m s 1 .
the higher brightness of the inner regions and the instrument•

分解能は厳しい•

円盤サイズを倍にすると見える•

小さいダストのコントラストが弱いの© 0000 RAS, MNRAS 000, 000–000
で、ALMAよりコントラスト弱め

resolution, the spiral structure is not readily detected. The
spiral arms and cavities clearly detected in (sub-)millimetre
ALMA images (see Fig. 7) are smeared out and are thus
marginally visible in the HiCIAO H-band image. It is worth
remarking that the shape of the HiCIAO PSF is characterized by large sidelobes which produce a decrease of the
dynamic range in the convolved images and, consequently, adecrease of the detected contrast between arm and interarm
regions.
We also consider a new set of simulations with a gas
disc model two times bigger. In Fig. 12 we present the synthetic HiCIAO H-band polarized intensity image of this specific disc model with two diﬀerent inclinations: 10° (left) and30° (right). A remarkable improvement of the detection of
spiral arms can be noticed with comparison to the model
previously considered (see Fig. 11), due to the fact that the
HiCIAO PSF is small enough to preserve spiral features ingrains dynamics under the action of systematic and random motion induced by the gas-dust aerodynamical coupling. Once having computed the resulting 2D dust density
distribution, we perform 3D Monte Carlo radiative transfer
simulation in order to make predictions of self-gravitating
discs for future observations.
The spiral density waves induced by GI produce a remarkable eﬀect on the dust dynamics leading to the formation of significant dust overdensities in the spiral arms
regions. The dynamical eﬀect produced by the occurrenceof pressure bumps in discs closely depends on the size of theparticles. We found that, assuming that particles are able
to grow through collisions with velocities up to 10 and 30
m s 1 , the dust density structure of centimetre sized particles is aﬀected by the presence of the spiral overdensities.