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 yrfor a
The inner environment of ZCMa: High-Contrast
Polarimetry
NaCocircular 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clubri
hanestaanprilarthe• 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/Nsq
aro
en
art
fro
en
polik
fac
aro
In

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

5.

Oumeis(20(wintcoA&A proofs: manuscript no. zcma_letter

Fig. 2. Right and center: PI images at H and K s bands. Grscale, 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

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 ≥2Fig. 4. Distribution of the clump mass associated with CORNISH HIIregions. The solid and dotted histograms are obtained by choosing, respectively, the far and near distances for the 11 targets for which theKDA 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 theCORNISH 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 LymanFig. 5gionsHi-G
maticfor u
no Ktance
NLy –90%
ber ocurve

maincontfectiertie
ThomTheLyma 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

cesandFig. 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 theIRAS 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、Lymanexcessの有無で違いはない• ほとんどの天体がsupervirial

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 areactually destroyed.

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

166.9

92

11.2

7.50

X

106

45

149

9.6

52

20.7

2.44

D

105

45

15

3.092

4.9

7.70

E

105

21147

4.6

52

4.9

2.46

F

105

101439

6.7

30

19.0

0.81

I

104

10136

2.152

4.6

2.56

J

104

51135

3.0

30

8.5

0.90

UZ

106

45149

18.2

52

39.3

2.9

UB

3 × 105

4545

10.0

68

18.9

6.0

UC

3 × 105

21443

14.6

36

37.9

1.9

UV

105

21148

12.25226.3

3.3

UU

105

101371

8.4

26

25.6

1.1

UF

3 × 104

10410

6.7

28

19.7

2.0

UP

104

2.59096

7.6

18

27.9

0.4

UQ

104

5.0

1137

5.4

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.

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

Fain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
31

dashedpline 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<1As 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->dust
trapping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 radialdistribution of the dust-to-gas ratio. Starting
q=4and->azimuthaltrapping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 fromthe lower
left cornerindicates
the size of theFigure
images of rms
the noise.
self-gravitating
model
obtainedstarting
an initialdust grain
size distributionwith
half-power
the synthesizedvelocity
beam (seetable
exponentq = contour
3.5 and of
a fragmentation
equalto 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 a3 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 value for ISM grains is ism = 1.7, several observations(e.g. Testi et al. 2001, 2003, Natta et al. 2004, Ricci et al.
2010) have found that disc < ism , which is naturally interpreted in terms of grain growth (e.g. Draine 2006). In thecontext of the present work, we calculate the value of theparameter by computing the ratio of the ALMA simulated
images at two diﬀerent frequencies, with the aim to infer thediﬀerent level of dust concentration of large grains in arm
and interarm regions. The parameter can be obtained as: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 ofcan 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, theparameter is relatively higher in the interarm regions, wheln F⌫1 ln F⌫2=2,(15)large grains are depleted due to the dust trapping mechln ⌫1 ln ⌫2nism. Thus, taking into account all the assumptions abo
where F⌫1 and F⌫2 are the fluxes at frequencies ⌫1 and ⌫2 ,dust grainproperties, theunstable
dust trapping
Dust thetrappingin gravitationallydiscs induced
13 by G
respectively. In Fig. 10 we show the map ofcomputed
produces a migration of large dust grains into spiral arm
from the model data (left) and the parameter calculated
region that can be potentially investigated through mulfrom two simulated ALMA observations at 650 µm and 3
wavelenght observations.

•

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