Star	  Forma)on	
  Newsle0er	
  #262	
  54-­‐58本目	
麻生 有佑	THE	  CHANDRA	  CYGNUS	  OB2	
  LEGACY	
  SURBEY:	
  DESING	
  AND	  X-­‐RAY	
  POINT	
  SOURCE	  CATALOG	
  Nicholas	
  J.	
  Wright	
  et	
  al.	  (2014)	

Ø Cygnus	
  OB2	
  associa)on,	
  d=1.4	
  kpc,	
  65個のO型星と数100のOB星。	
  Ø Chandra	
  Cybnus	
  OB2	
  Legacy	  Surveyのから点源を決定し、カグを作った。	
  Ø 観測、可視、近赤外、OB2	  memberなど分担した論文がいくつかある。	
  Fig. 1.— Left: the “as designed” Cygnus OB2 Chandra Legacy Survey tiling for the 36 30 ks pointings. The thick bleft) represents one ACIS-I pointing,the thin boxes
all the pointings. Di↵erent colors show
areas with di↵erent numbØ 8.3’のCCD	
  chipが4つあり、8’グ
で6x6で、Fig	  1のように傾けてグ。	
  pointings: teal – 4 overlapping pointings; blue – 2 overlapping pointings; purple – 1 pointing. The black bars show roof one pointing (⇠160 ), of the deep inner area (⇠420 ), and of the total field (⇠560 ). Pointing 1-1 lies at the boØ )l)ngの効果。	
   dimensions
1-6 lies at the top right (NW), 6-1 lies at the bottom left (SE) and 6-6 lies at the top left (NE). Right: the “as executed”
the Cygnus OB2 Chandra Legacy Survey and complementary observations. The color bar gives the achieved e↵ect expos
まず重なる枚数をばらつかせるて内側で多くできる(図の緑、黄色)。	
  The deepest region of the survey with a non-negligible area has an exposure time of ⇠220 ks.
From the list of cross-matched sources, which typically少なくとも1枚は“on-­‐axis”に近くなる点を増やせる	  contained between ⇠20 and ⇠100 sources per ObsID, asßspacecrabの中心がCCDの中心とずれてるから(Garmire	
  trometric o↵sets were calculated in RA and
Dec. Theet	
  al.	
  2008)	
  spacecraft roll angle is known to a very high precision
based on guide stars spaced3 a degree apart or more on thesky, and consequently astrometric uncertainties arisingfrom roll angle uncertainties are negligible compared tothose arising from knowledge of absolute pointing. Rollangle changes were therefore set to zero for all ObsIDs
for the purposes of astrometric corrections. The o↵sets
between the X-ray and near-IR positions have mean values of R.A. = 0.0400 and Dec. = 0.0300 with all values
smaller than ⇠0.100 . The astrometric o↵sets were applied
to the L1 data products to create new aspect solutions foreach ObsID and the data were then re-reduced following
the procedure in Section 2.3.

Chandra Cygnus OB2 Legacy Survey Catalog

Fig. 1.— Left: the “as designed” Cygnus OB2 Chandra Legacy Survey tiling for the 36 30 ks pointings. The thick black box (bottom
left) represents one ACIS-I pointing, the thin boxes all the pointings. Di↵erent colors show areas with di↵erent numbers of overlapping
pointings: teal – 4 overlapping pointings; blue – 2 overlapping pointings; purple – 1 pointing. The black bars show roughly the relative
dimensions of one pointing (⇠160 ), of the deep inner area (⇠420 ), and of the total field (⇠560 ). Pointing 1-1 lies at the bottom right (SW),
1-6 lies at the top right (NW), 6-1 lies at the bottom left (SE) and 6-6 lies at the top left (NE). Right: the “as executed” exposure map for
the Cygnus OB2 Chandra Legacy Survey and complementary observations. The color bar gives the achieved e↵ect exposure in units of ks.
The deepest region of the survey with a non-negligible area has an exposure time of ⇠220 ks.

2.5. Exposure maps and survey sensitivity

From the list of cross-matched sources, which typicallycontained between ⇠20 and ⇠100 sources per ObsID, astrometric o↵sets were calculated in RA and Dec. The
spacecraft roll angle is known to a very high precision
based on guide stars spaced a degree apart or more on thesky, and consequently astrometric uncertainties arisingfrom roll angle uncertainties are negligible compared tothose arising from knowledge of absolute pointing. Rollangle changes were therefore set to zero for all ObsIDs
for the purposes of astrometric corrections. The o↵sets
between the X-ray and near-IR positions have mean values of R.A. = 0.0400 and Dec. = 0.0300 with all values
smaller than ⇠0.100 . The astrometric o↵sets were applied
to the L1 data products to create new aspect solutions foreach ObsID and the data were then re-reduced following
the procedure in Section 2.3.

Exposure maps were constructed for each ObsID on
a per-CCD basis and in multiple energy bands using
the standard CIAO tool sequence of asphist, mkinstmap, and mkexpmap. The exposure maps were calculated using a thermal plasma model spectrum withkT = 1.35 keV and NH = 1.25 ⇥ 1022 cm 2 (typicalof a stellar coronal X-ray source at the extinction ofFig. 2.— Top: Histogram of the exposure tim
tions used in this work. The narrow peaks lie atposure values generated from the overlapping gridPeaks are also visible at ⇠80 ks and ⇠140 ks w
3 and 5 grid pointings overlap due to variationsof neighboring ObsIDs, as well as at ⇠170 and ⇠existing 50 and 100 ks observations overlap with

THE	  CHANDRA	  CYGNUS	  OB2	
  LEGACY	
  SURBEY:	
  DESING	
  AND	  X-­‐RAY	
  POINT	
  SOURCE	  CATALOG	
  Nicholas	
  J.	
  Wright	
  et	
  al.	  (2014)	
  続き	
Ø 点源の検出方法は3種類	
  Ø wavdetected...	
  CCD	
  chipに対して、2,	
  4,	
  8,	
  16,	
  32pixelのwaveletを使い、閾値もS0=10–6,	
  
10–5,	
  10–4	
  (6が厳しい)。検出数の比は1:1.4:2.6。S0の定義も含めFreeman	
  et	
  al.	  (2002)	
  Ø PWdetected...	
  Damiani	
  et	
  al.	  (1997)。閾値σ=4.1。PSFの半分以上の大きさのみwavelet変換。	
  Ø Enhanced	
  wavdetected...	
  異なるグに対してoverlapがあるので、合わせて解析したい。グが違うと1つの ク
でのPSFの値が異なるので、上手く工夫。	
  Ø さらにここでは未検出なCyg	
  OB2
も他から加えた。	
  Chandra Cygnus OB2 Legacy Survey Catalog

9

4

THE	
  CHANDRA	  CYGNUS	  OB2	
  LEGACY	
  SURBEY:	
  DESING	
  AND	  X-­‐RAY	
  POINT	
  SOURCE	  CATALOG	
  Nicholas	
  J.	
  Wright	
  et	
  al.	  (2014)	
  続き	
Ø 点源候補について、PSFの40%コ
ア として抽出。	
  Ø PSFの99%を含む半径から100	
  photonを含む半径までを背景として抽出。	
  Ø 背景のﬂuctua)onである確率が>0.01となる点源候補を棄却。	
  Ø 棄却されるものがいなくなるまでこれを繰り返す。à7924個検出。	
  Ø (測光、分光の測定量のS/Nを半減するようなを取り除く。)	
  Ø smallest	
  oﬀ-­‐axis	
  angleが小さい。位置の精度も良い。	
  ß)l)ngによって少なくとも1度はon-­‐axisで観測される範囲が増えた。	
  Wright et al.

Chandra Cygnus OB2 Legacy Survey Catalog13

CO 6–5
CO 7–6
CO 8–7
CO 9–8
CO 10–9
CO 11–10
CO 12–11
CO 13–12
13 CO 6–5
13 CO 7–6
13 CO 8–7
13 CO 9–8
13 CO 11–10[C i] 3 P2 –3 P1
HCN 11–10

691.47806.65921.801036.91
1151.99
1267.01
1382.00
1496.92
661.07771.18881.27991.331211.33
809.34974.49433.56371.65325.23289.12260.24236.61216.93200.27453.50388.74340.18302.41
247.49370.41
307.64

116.2
154.9
199.1
248.9
304.2
365.0
431.3
503.1
111.1
148.1
190.4
237.9
348.9
62.5
974.5

15.7 [2.3]
23.5 [2.7]33.1 [3.5]45.3 [4.1]
55.3 [3.7]51.3 [4.5]58.3 [3.3]60.7 [4.7]⇤ 5.9 [2.5]⇤ 5.8 [2.4]⇤ 6.3 [3.4]⇤ 13.9 [4.8]
16.8 [3.4]⇤ 7.6 [2.7]⇤ 8.9 [5.5]33.0 [2.6]a34.9 [2.8]43.0 [3.6]43.0 [4.3]41.1 [4.0]48.6 [4.5]38.3 [3.8]43.0 [4.7]- [2.3]⇤ 3.8 [2.5]- [3.1]⇤ 7.1 [4.0]- [3.0]⇤ 5.6 [2.8]⇤ 6.4 [4.2][C i] 3 P1 –3 P0
CH+ 1–0
H2 O 110 –101
H2 O 211 –202
H2 O 202 –111
H2 O 312 –303
H2 O 111 –000
H2 O 312 –221
H2 O 321 –312
H2 O 220 –211

492.16835.08556.94752.03987.931097.361113.34
1153.131162.911228.79

609.14359.00538.29398.64
303.46273.19269.27259.98257.79243.97

23.6
40.1
61.0
136.9
100.8
249.4
53.4
249.4
305.2
195.9

- [4.3]- [2.1]- [3.8]- [2.1]- [4.9]- [3.2]- [3.1]- [2.9]- [2.4]- [4.1]- [4.2]- [2.3]- [3.8]- [2.0]- [3.7]- [2.8]- [2.7]- [3.7]- [4.0]- [3.8]⇤ 6.3 [3.5]⇤ 7.4 [2.9]⇤ 18.4 [7.5]⇤ 12.4 [5.6]

23.8 [5.4]
34.6 [4.7]36.4 [6.8]27.9 [5.9]49.2 [8.7]- [2.2]- [2.7]- [5.6]- [4.3]- [4.2]- [2.3]- [4.9]⇤ 9.1 [3.8]- [4.0]- [5.3]- [4.2]- [3.2]- [4.3]- [2.9]- [2.6]- [3.2]- [3.0]- [6.2]⇤ 8.2 [3.7]11.7 [3.4]11.3 [2.9]11.7 [3.4]
15.2 [3.7]- [1.9]- [2.3]- [4.0]- [3.5]- [3.0]- [1.9]- [4.5]11.8 [1.4]
11.6 [1.5]
11.6 [2.3]12.8 [2.6]12.4 [2.2]
12.3 [2.4]
11.6 [2.1]
11.6 [2.3]- [1.4]- [1.3]- [1.8]- [2.7]- [1.9]⇤ 1.4 [1.5]- [2.5]- [1.2]
5.3 [1.7]- [2.4]⇤ 5.1 [2.0]⇤ 3.9 [1.8]
9.1 [2.1]
8.0 [2.0]
7.5 [2.3]- [1.3]- [1.1]- [2.0]- [2.4]- [1.7]- [1.2]- [2.5]- [3.7]- [3.2]- [3.1]- [3.0]- [3.5]- [2.4]- [2.9]- [3.1]- [3.1]- [3.5]- [2.5]- [1.4]- [2.0]- [1.5]- [2.7]- [1.5]- [1.5]- [1.8]- [2.2]- [2.1]- [2.4]- [1.6]- [2.4]- [1.5]- [2.3]- [1.4]- [1.5]- [1.6]- [1.7]- [1.9]⇤ 9.0 [3.0]Signatures	
  of	
  warm	  carbon	  monoxide	
  in	
  protoplanetary	  discs	
  observed	
  with	  Herschel	
  SPIRE	
  
M.	
  H.	
  D.	
  van	
  der	
  Wiel	  et	
  al.	  2014	Ø 原始惑星系円盤では密度や温度の幅が広いので多波長が必要。	
  Ø CO	
  Jup=4-­‐15のような遷移は大気の窓や一度に取れる
幅が狭かった(Hershel	
  HIFI)。	
  Ø Hershel	
  SPIREでCO	
  Jup=4-­‐13の分光、12個のHerbig,	
  6個のTTauri。	
  Ø 空間的に離れた(>1’)成分を引いてある(10-­‐30”のエの影響が既知のものもある)。	
  Ø 6天体で3σ以上を3輝線以上検出。	
  Notes. a Suspected contribution from a marginally extended envelope.

6

M. H. D. van der Wiel et al.
6
4
20

HD 100546

Figure 3. (a) Zoom of the 60 GHz section of the DR Tau spectrum centred around the CO 7–6 and [C i] 3 P2 –3 P1 lines. The line of theneutral carbon atom coincides with the first sidelobe of the neighbouring CO line, but has significant flux of its own before backgroundsubtraction (grey), while its strength is reduced considerably after background subtraction is applied (black; same spectrum as in Fig. 2).
(b) Zoom of the 60 GHz section of the HD 163296 spectrum centred around the [N ii] line, again before and after background subtraction.
This shows that the N+ line towards HD 163296 originates from the extended background medium.

c 2014 RAS, MNRAS 000, 1–16

Figure 2. SPIRE spectra of sources with more than three spectral lines detected at >3 . Spectra are shown here after subtraction of athird order polynomial fit to the continuum baseline. The high frequency end of SLW (>960 GHz) is not shown because it is noisier thanthe overlapping part of SSW. Laboratory frequencies (JPL database, Pickett et al. 1998) for species detected in our spectra are indicatedin the third panel. Dotted lines in the other panels mark those lines that are detected in each individual source. Note that each spectralline is unresolved and is dominated by the instrumental Sinc profile (discussed in Section 3.1 and see also Fig. 3), including sidelobes,
mainly visible near the stronger lines, e.g., the highest frequency CO lines in the top panel. The broader spectral features visible in theSignatures	
  of	
  warm	  carbon	  monoxide	
  in	
  protoplanetary	  discs	
  observed	
  with	  Herschel	
  SPIRE	  on accuracy

H.	
  D.	
  vinan	
  
der	
  
Wiel	
  et	
  al.	  (2014)	
  続き	
linesM.	
  detected
ourHerle10 3)M.have
also
been
H. D. vander Wiel
et al. obamely CO 6–5, CO 7–6 andrespectively, g is the statistical weight of state ‘u’, A is(a) HD 100546

22
20
18
16

PACS
ground-based

ln

4 Ful
gu Aul h ul

(m 2 )

24

Trot = 215 K

Trot = 879 K

SPIRE

14
0

500

1000 1500 2000 2500 3000 3500(b) AB Aur

20

16

Trot = 78 K

PACS

ground-based(m 2 )

18

ln

22

4 Ful
gu Aul h ul

24

Trot = 267 K

SPIRE

14
0

500

1000

(c) HD 97048(m 2 )

18

ln

4 Ful
gu Aul h ul

24
22

Trot = 22 K

PACS

20
Trot = 231 K

16
SPIRE

14
0

500

1000

4 Ful
gu Aul h ul

18

20

16

Trot = 65 K

ground-based(m 2 )

22

1500(d) DR Tau

24

ln

1500

Trot = 126 K

SPIRE

14
0

500

1000
Eup /k (K)1500

Figure 4. Rotation diagrams for the targets with at least nine
CO detections in the SPIRE and PACS bands combined. Notethe di↵erent horizontal scale limits for HD 100546. Vertical error bars include contributions from measurement uncertainties(cf. Table 3) and absolute calibration uncertainties. Upper lim-Ø rota)onal	
  diagramを用いて12COを解析。	
  Ø ßLTE,	  より小さい、単一温度、光学的に薄い、を仮定(LTEしか成り立たないが便利)。	
  
u

ul

the Einstein coefficient for spontaneous emission, ⌫ul is thefrequency related to the energy di↵erence between states ‘u’
and ‘l’, Q is the partition function, and d is the distanceto the source. If all the above assumptions hold, a determination of a will yield the rotational temperature Trot . Likewise, c would be a measure for the total number of emitting
molecules N , but only if the beam filling factor f in Eq. (1)is either assumed to be uniform or can be constrained using
high angular resolution observations. Since this paper does
neither, the analysis in this section is restricted to the slope,a, of the rotation diagrams.
Figures 4a,b,c show rotation diagrams of three Herbig
Ae/Be sources with at least nine CO line detections in Herschel SPIRE (this work) and Herschel PACS (Meeus et al.
2013) combined, augmented by ground-based datapoints inthe (sub)millimetre gathered from the literature (see Appendix A). Note that some of the CO rotational lines probed
here may be optically thick. The high optical depth is corroborated by predictions from numerical models (e.g., Bruderer et al. 2012) as well as by the low value of the observed
12
CO/13 CO integrated intensity ratio (in cases where 13 CO
is detected: ⇠3–10 for AB Aur and HD100546, see Table 3).
Thus, the derived values for rotational temperature should
not be interpreted as the kinetic temperatures of the gas:
Trot 6= Tgas .
To compare our results with rotational diagrams fromearlier work by Meeus et al. (2013), we derive rotationaltemperature values from linear regression fits to two separate energy ranges, chosen such that each of the two can be
represented as well as possible by a straight line. The Trot
values of the warm components found in Fig. 4a,b,c are consistent with those from fits to the PACS-only dataset from
Meeus et al. (2013), i.e., the warmest of the two components
for HD 100546 and the single-component fits for HD 97048
and AB Aur. For HD 100546, the break between the cold
and warm components is at the same energy of Eup /k =1200 K both here and in Meeus et al. (2013). For the othertwo sources, the highest few lines in the SPIRE range fallon the same linear relation in the rotation diagram as thatdefined by the PACS measurements.
We have chosen not to include the CO 23–22 line
(Eup =1524 K) in this analysis because there are indications
that this PACS detection may su↵er from instrumental effects that lead to a systematic overestimate of the line flux.
For example, Meeus et al. (2013) showed that CO 23–22 falls
above the trend set by the other lines in all six sources of
their sample of eight in which this line is detected, including
AB Aur and HD 97048.
When the lines from the SPIRE FTS measurements areadded to the rotation diagrams (Fig. 4a,b,c), it becomesapparent that additional, colder components need to be invoked for AB Aur and HD 97048, at ⇠80 and ⇠20 K, respectively. In the case of HD 100546, for which Meeus et al.
(2013) already introduce a two-component fit, the colder of
the two drops from ⇠280 to 215 K by adding the SPIREØ Hershel	
  PACSも合わせて、高温低温に分けて ィ 。	
  
Ø HD	
  100546はPACSだけで2温度とわかっていたが、SPIREの部分は曲線で多温度のように見える。	
  Ø AB	  Aur,	
  HD	
  97048,	
  DR	
  Tauについて、PACSではわからなかった低温成分を検出。	
  Ø しかし、先の条件は成り立たないのでsophis)cated	
  modelも使う(ProDiMo)。	
  rotoplanetary disc models

m to provide an exhaustive exs using complex protoplanetary
esent in this section a compariof three sources for which dedels have been published in thethe discs around HD 100546,
are generated by the thermoc model code ProDiMo (Woitke,
et al. 2010) and are based on aobservational constraints. The
tions (SLED) for the first two ofFig. 5. In any targets where the
CS studies (Meeus et al. 2013),
pper limit (see Section 4.1).

3 5

15

Jup
20

25

30

35

0.6

HD 100546

0.5

1000
Eup /k (K)

Jup
15
20

10 16 HD 163296

2000

25

gas

0
=2
s

30

35

SPIRE observations
other observations
ProDiMo (Tilling et al. 2012)

10 17

K

T ga

3500

0.3

0

CO J=18-17
CO J=12-11

75 K

10

500

0.4

s

3 5

200

5

Tga =

10 50

4

T

10 18

1 2 3
log nCO [cm-3]
K

SPIRE observations
other observations
ProDiMo (Thi et al. 2011)
Bruderer et al. model

0

50

10 17

-1

=7

10 16

10

z/r

ation diagram of the T Tauri obmponents are required to fit the
spectrally resolved CO 3–2 mea(2001) to separate contributions
ed envelope. The point shown in
nts purely the disc contribution.
y unresolved, 4–3 line observedsu↵er from contributions fromwe plot it as an upper limit inderived rotational temperature
DR Tau is lower than for thee other three panels. Confirmad come from (non-)detections ofobservations (Blevins et al., inh energy range unaccessible toE observations presented here.

integrated line intensity (W m 2 )of CO gas is revealed that was
tudy. As discussed above, mostthe rotation diagram method is
cal conditions prevailing in proy, molecular line transitions of(radially) di↵erent locations inmodels must be applied to takeysical and chemical structure of

integrated line intensity (W m 2 )Signatures	
  of	
  warm	  carbon	  monoxide	
  in	
  protoplanetary	  discs	
  observed	
  with	  Herschel	
  SPIRE	
  
M.	
  H.	
  D.	
  vWarman	
  carbonder	
  monoxideWiel	
  inet	  al.	
  (2014)	
  protoplanetary
discs 続11 き	
12 M. H. D. van der Wiel et al.

0.2

CO J=7-6
CO J=3-2
0.1

CO J=30-29
10 18
10 50

200

500

1000
Eup /k (K)

2000

3500

Figure 5. Observed CO spectral line energy distributions of
the two targets with at least three CO lines detected in the
SPIRE bands, for which a ProDiMo model is also available (black
crosses). In addition, for HD 100546 we also show in green plus
symbols the model provided by Bruderer (private communication), calculated directly from the model as published in Fedele
et al. (2013b). Observed line detections are shown as (blue)
squares, and 3- upper limits are shown as triangles for lines thatare either not detected or detected at 63 . Error bars represent
the combination of line intensity errors (see Table 3 for SPIRE)
and the absolute calibration uncertainty, 6% for SPIRE, 30% for
PACS (see PACS Observers’ Manual 2013, section 4.10.2). Data
points with open markers are either ground-based (Appendix A),
i.e., at lower energies than the SPIRE data points, or from Herschel PACS measurements at higher energies (Meeus et al. 2013
for HD 100546; Meeus et al. 2012 for HD 163296).

0.0
10

100
r [AU]

Figure 6. Illustration of the origin of several CO rotational linesØ 2天体について他の波長で ィ した
delimited by the cumulative 15 and 85% line flux contributionof each transition, both vertically and radially. The grey-scaleØ HD	
  100546は概ね合っているが、よりも積分強度が下がり始める温度が低い。	
  indicates number density of CO molecules. Dashed contours markgas temperature. The inner disc, at <10 au, is not shown, since itØ
はNIR,	
  mmの連続波を ィ したもの。COは上層を見ており、PAH加熱などが必要。	
  does not contribute significantly to the total amount of line flux
12COが合わなくなる。	
  from pure rotationaltransitions of CO.
Ø 13COは10-­‐30倍弱い(図は無し)。全質量を大きくすると
Ø 13COは減光が小さい領域でアが減るなどのisotope	
  selec)ve	
  processが必要。	
  by SPIRE and PACS are optically thick, the shape of theØ HD	  163296のは可視からmmの連続波と低励起COをィ 。	
  rotational SLED is sensitive to adjustments of the gas temperature in the relatively tenuous, irradiated upper layer ofØ 2−4倍が弱い。100	
  AU以内がより高温のはず。	
  the disc, where PAHs are the dominant heating agent.

100546 shown in the top panel
k by Thi et al. (2011), and fits
etre continuum spectral energy
ed interferometric data of thisTatulli et al. 2011). The modeli] fine structure lines, the CH+011), low-J CO emission, as well
nal lines in Hein Bertelsen et al.

ally separated origin of the suite

Although the ProDiMo model for HD 100546 was not
tailored to fit the intermediate- and high-J range of the CO
rotational ladder, the model points are always within a factor 3 of the observed CO line fluxes (Fig. 5, top panel).

in a model of the HD 100546 disc. Solid outlines represent areasと12CO	
  SLED	  (輝線のSED)を比べる。	
  could play a r
are implemenis needed to dinant in the 1
the measured
total disc mamore opticall
For comp
CO SLED frveloped indep
recently upda
have focused
Herschel PAC
presented inserved SLED
while also reenergy distri
taken in eac
some aspects
efficiency of smation rates
codes are pro
ature (see thedeviations in
istry are exp
the gas temp
and thus thegas-to-dust rtimes higher
Given o13
CO line int
selective chemthe disc of H
implemented
erer et al., w
line intensitie
and assumpt
guide refinem

Apart from the shape of the optically thick 12 CO SLED,Signatures	
  of	
  warm	  carbon	  monoxide	
  in	
  protoplanetary	  discs	
  observed	
  with	  Herschel	
  SPIRE	
  
M.	
  H.	
  D.	
  van	
  der	
  Wiel	  et	
  al.	  (2014)	
  続き	
Ø HerbigはTTauriよりUVが強く、またアしているので上層が高温。X線はTTauriの方が強いが影響は少ないと考えられている。	
  ß今回の観測でもJ=10-­‐9を比べて、TTauriやアしていないHerbigは強かった。	
  Ø 12CO	
  10-­‐9は12CO	
  18-­‐17,	
  [OI]3P1-­‐3P2と相関。ゴは4,	  5天体だが同じ場所から来ているかも。	
  Ø HD	
  169142とHD	
  100546で初めて12CO	
  J=11-­‐10を検出。	
  2.1 Observations: Serkowski curveEﬀects	  grain	
  growth	
  on	
  the	
  interstellar	
  polariza)on	
  
urve	
  
Theof	
  wavelength
dependenceof polarization
P c(λ)
in the visible
Nikolai	  .	
  Voshchinnikov	
  &	
  Hiroyuki	
  byHirashita	  (2014)	
   formula suggested
partVofspectrum is describedan empiricalby Serkowski (1973)Ø 星間粒子の直線偏光度の波長依存性P(λ)はSerkowski	
  curveで表される。	
  P (λ)/Pmax = exp[−K ln2 (λmax /λ)].

(1)Ø 特に天の川銀河ではKとλmaxに経験則がある。	
  This formula has three parameters: the maximum degree of polarization Pmax , the wavelength corresponding to it λmax
and the coØ 部分的に整列、非球体、減光断面積、を考えた
で偏光の時間変化を調べる。	
  Ø Hirashita	  &	
  Vcharacterizing
oshchinnikov	
  (2014)ではaccre)onとcoagula)onによるgrain	
  efficient K
the width of the Serkowski curve. The sizeの進化を計算した。	
  values of Pmax in the diffuse interstellar medium usually do not ex3	  cm–3,	
  T =10	
  K,	
  S =0.3を固定し、silicate	
  dustとcarbonaceous	
  dustの混合。	
  Ø Z=Zo,	
  nH=10andgas
ceed 10%,the meanvalueaccof λmax is 0.55 µm (Serkowski et al.
Ø しかし、観測(後述のTaurusなど)と比べると相違点があり、tuned	
  modelを作った。	
  1975).
(1)
では>30	  Myrではcoagula)onだけになって、減光曲線を平らに、Rvを大きくする。	
  Theparameter K determines the half-width of the normalized
観測ではRv~4.4と大きくても、carbon	
  bumpが0.22	
  µmにある。àcarbonaceousはあまり凝interstellar linear polarization curve
縮しない。	
  (2)観測ではRvが増えるとAv/NHが急に減るが、では減らないàsilicateが~0.03	
  µm
W = λmax /λ− − λmax /λ+ ,(2)
より成長する(coagula)on	
  threshold	
  velocityがなく、衝突すれば必ず凝縮する)。	
  where λ− < λmax < λ+ and P (λ+ ) = P (λ− ) = Pmax /2. The
relation between W and K is as follows:
W = exp[(ln 2/K)1/2 ] − exp[−(ln 2/K)1/2 ].

(3)

Treating K as a third free parameter of the Serkowski curve,

Eﬀects	
  of	
  grain	  growth	
  on	
  the	
  interstellar	
  polariza)on	
  curve	
  Nikolai	
  V.	
  Voshchinnikov	
  &	  Hiroyuki	
  Hirashita	  (2014)	
  続き	
  Ø 整列の仕方はimperfect	  Davis-­‐Greenstein	  alignmentを、小さい粒子の整列を減らすように修正(rv,cutというを導入)。	
  

Eﬀects	
  of	
  grain	  growth	
  on	
  the	
  interstellar	
  polariza)on	
  curve	
  Nikolai	
  V.	
  Voshchinnikov	
  &	  Hiroyuki	
  Hirashita	  (2014)	
  続き	
  Ø ゴイ は0.001	
  µm-­‐1	
  µmに59グ、波長は0.2	
  µm-­‐5	
  µmに77グ調べた。	
  Ø 観測としてTaurus,	
  Chamaeleon,	  ρ	
  Oph,	
  R	
  CrAから57個。	
  Ø original	
  modelはほとんどP(λ)が時間変化せず、tuned	
  modelは時間変化する。	
  Ø また、original	
  modelではrv,cutを変えても、λmax>0.7	
  µmを説明できない。	
  àtuned	  modelが良い。	
  Ø さらに、prolateにしたりoblateにしたりa/bを変えたりするとK-­‐λmaxの関係が変化し、説明できる範囲が増える(上ではprolate,	
  a/b=3)。	
  Ø 観測を説明するために必要な時間<30	
  Myrは化学進化の3-­‐6	
  Myrや力学進化の5-­‐25	  Myrという イケに矛盾しない。	
  Ø 固定したnH=103	
  cm–3を104	
  cm–3にすれば時間は0.1倍。	
  Polarization curve and grain growth1

0.5

T=0 Myr
T=300 Myr
average Serkowski curve

0
0.2

0.5

1

2.0

5

Forma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  
F
Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  

C

C

C

F

10 5

Ø 惑星形成はmm-­‐cmのpebble(小石)に
ice粒子が積もることで始まる。	
  
10 6
10 4
10 2
100
102
10 2
100
102
Ø
重力的に束縛されたpebbleたちが衝突してエギを放出し、収縮する様子を計算
Collision speed v (m s 1 )Collision speed v (m s 1 )した。	
  omesØ
as 収縮にかかる時間と最終的なゴイ
a function of projectile radius, a, collision speed,v, and relative particle size, f ⌘ m2 /m1 . Green regions i
分布を求めたい。	
  ions bouncing and red regions fragmentation. Outcome regions taken from Güttler et al. (2010).
Ø 粒子の衝突には衝突速度と粒子の大きさに応じて、coagula)on,	
  bounce,	
  fragmenta)onがある。	
  Ø 衝突速度は全体の質量で決まり、軽いと遅く、fragmenta)onせずpebbleから微惑星がi) Two similar-sized particles colliding or a larger partできる。	
  colliding with a small particle, 1: f 0.1
Ø fragmenta)onの度合いで、できる微惑星のporosityと密度が決まる。	
  Rii) A small
particle, 2, ‘running’ into (colliding with) a
porousでrockが多い。	
  .Ø 密度はiceが多いと低くなり、つまり重いとless	
  (6)particle, 1: f < 0.1
Karl Wahlberg Jansson and Anders Johansen: Formation of pebble-pile planetesimals

0 4ange in ⌘ after a dissipative collision is< | ⌘max | the new value of ⌘ is equal toe, otherwise the new value is the old plus
not enough time for the system to virialize
0.1

Projectile radius a (m)

B

10 3

We do a simple interpretation of Fig. 11 in Güttler(2010), for collisions between compact particles:

Large target: f < 0.1

vstick

10 1

B

10 3

vstickProjectile radius a (m)Large projectile or similar sized particles: f

10 1

v (m s 1 )v < vstick
updating ⌘K we need to take into accountvstick  v < 1
s both dissipated in the collisions and re1  v < 25
d contracts, so using Eqs. (1), (3) and (4)Fig. 1. Collisional outcomes as a function of projectile radius, a, collision speed, v, and relative particle size, f ⌘ m /m . Green regions indicate
coagulation, yellow regions bouncing and red regions fragmentation. Outcome regions taken from Güttler et al. (2010).
v 25
C

F

C

C

10 6

F

10 5

10 5

10 4

10 2

Collision speed

100v (m s 1 )

102

10 6

10 4

10 2

Collision speed

100

102

v (m s 1 )21i) Two similar-sized particles colliding or a larger particle, 2,

f

0.1
C
B
F
F

f < 0.1
C
B
CF

Forma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  続き	
  
Ø
で非弾性衝突による収縮時間とゴイ 分布を調べる。	
  Ø 初期のpebblesはmm,	
  cm程度のmonodisperseなゴイ 分布。	
  
Ø 最終的に微惑星から冥王星ゴイ Rsolid=1-­‐1000	
  kmができる質量で、カイと太陽を想定し、半径に広がるpebbleの集まり(雲)から始める。	
  Ø pebblesはSiO2の密度の一様球でµmゴイ の粒子からなる。	
  Ø 速度は系全体の運動エ
ギ から決まるMaxwell分布。	
  Ø 衝突は0次元、回転無し、representa)ve	
  par)cle	
  approach	
  (Zsom	
  &	
  Dullemond	
  2008)をカで統計的に計算。	
  à質量が変わらず粒子数が変わる”群れ”を考え、ある群れを代表する粒子iと粒子kが代表する群れの中にいる粒子との衝突率を時間毎に計算する。	
  Karl Wahlberg Jansson and Anders Johansen: Formation of pebble-pile planetesimals
次の衝突までの時間は全ての衝突率の和で決まりるPoission	
  process。	
  BeforeE = Ebind + Ecoll

After
E

′

′
= Ebind= E > Ebind

E

Nmon = 3EcollrollForma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  続き	
  Ø 時間変化するとして3つ。(左式)	
  Ø エ
ギ が抜けて–∞になったら終わり。	
  Ø 最終的な半径Rsolidと収縮にかかる時間tcollの
関係を
。	
  
Ø pebblesが初めcmゴイ のとき3種類ある。	
  (1)Rsolid<50	
  km	
  bouncingが支配的でtcoll∝Rsolid–1	
  (2)50	
  km<Rsolid<70	
  km	
  fragmentaionが効く。重いとfragmenta)onが早く、頻繁に起こる。	
  
解析的にはpebblesのゴイ をaに固定するとtcoll∝a/Rsolid–1。	
  (3)Rsolid>70	
  km	
  衝突がとても頻繁に起こり、自由落下時間が収縮。密度は同じなので自由落下時間も一定。	
  A&A proofs: manuscript no. KWJ_AJ_Pebble_arXiv

Collapse time as a function of solid radius of
the planetesimal.
αPower-law fits: tcoll ∝ Rsolid
10000

Simulations (cm)Collapse time tcoll (yrs)1000Ø 初めmmのときはfragmenta)onが効く前に、自由落下で決まる ェ
になる。	
  

100

10

1
0.1

1

10

100
1000
Solid radius Rsolid (km)10000

100000Forma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  続き	
  Ø ηたちを見ると、重いとfragmentが起こり、軽いときより早く収縮が進むとわかる。	
  Ø さらに、重いときはη=0.5のあたりで、次の衝突までにア 平衡(η=ηeq=ηK)にならなくなり、自由落下速度で収縮する(エギ がどんどん抜けて速度が落ちる)。	
  Karl Wahlberg Jansson and Anders Johansen: Formation of pebble-pile planetesimals
Collapse parameter η as function of time.
Planetesimal has a solid radius of Rsolid = 5 km.
Simulated η
Simulated ηK
Simulated ηeqFree-fall collapse
Analytic solution

0.8

Simulated η
Simulated ηK
Simulated ηeqFree-fall collapse
Analytic solution

1
Collapse parameter η (R/R0)

1
Collapse parameter η (R/R0)

Collapse parameter η as function of time.
Planetesimal has a solid radius of Rsolid = 50 km.

0.6

0.4

0.2

0.8

0.6

0.4

0.2

0

0
0

0.2

0.4

0.6
Time t (tcoll)

0.8

1

1.2

0

0.2

0.4

0.6
Time t (tcoll)

0.8

1

1.2

A&A proofs: manuscript no. KWJ_AJ_Pebble_arXiv

Speed (m s-1)

Fig. 4. The total energy, parametrized through ⌘ ⌘ R(t)/R0 , as function of time for two clouds with, initially, cm-sized pebbles. The left plot
displays the collapse of a cloud with Rsolid = 5 km and the right plot a cloud with Rsolid = 50 km. The red lines show the actual size of the cloud,actualof theenergy
particle cloudof time.
blue lines the equilibrium values (⌘eq ⌘ E0 /E) and Virial
greenandlines
thespeedkinetic(⌘K ⌘asTfunction
pink line shows the free-fall collapse of a cloud
0 /T ). The
Values for three different planetesimal sizes: 5, 50 and 500 km
and the teal lines the analytic solution from Appendix
A. Note that time is measured in units of the collapse time. Therefore the free-fall line isdi↵erent for the two clouds. In the case100of the low-mass cloudActual
the collapse
speed, 5 kmis slow and the only outcome of collisions is bouncing. Therefore it isin virial equilibrium until the very end of the collapse and follows the analytical solution (⌘ = ⌘K = ⌘eq ). The massive cloud, on the other hand, has
initially higher collision speeds which lead to pebble fragmentation and rapid energy dissipation. At about ⌘ = 0.5 the energy dissipation becomes
faster than free-fall, so the cloud does not
10 get into virial equilibrium before the next pebble collision (the first term in Eq. (7) dominates over thesecond). This results in a transition to cold collapse (⌘ is parallel to the free-fall curve) with sub-virial speeds (⌘K > ⌘eq ).
1

are lower and collisions will only result in bouncing and the primordial pebbles survive the collapse.

between the primordial pebbles. Collision speeds only reachvalues high enough for fragmentation to occur in the very
0.1
end of the collapse. The power-law fit shows us that tcoll /1
Rsolidwhich is the same dependence as we get from the an3.1. Initial conditions and simulation setup
alytic derivations when we assume that the particle size is
0.01a self-gravitating cloud ofconstant and bouncing is the only collisional outcome (ApIn our simulations we start out with
0
0.2
0.4
0.6
0.8
1
1.2
A). We findthat thecollapse time increases linearlysilicate pebbles with a monodisperse size distibution. The peb-Time t (t pendix
coll)
with
pebble
size,
which
also
agrees with the analytic model
bles are assumed to be homogeneous spheres (density ⇢SiO 2 ⇠2.5
in Appendix A.
g cm 3 ) built up of µm-sized monomers. The reason for this den-Forma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  続き	
  Ø 最終的にcm	
  pebbleだけだとporousで低密度。	
  Ø mm以上のものもできると高密度。	
  àカイ
では低質量ほど低密度なので、合う。	
  Ø ここで扱った ェが終わると自己重力で密度は増える(Kataoka	
  et	
  al.	  2013)。	
  Ø Rsolid=100	
  kmのときのゴイ 分布を。	
  Ø 初めはcmでmonodisperse、fragmentがあるとµmゴイ ができる。	
  Karl Wahlberg Jansson and Anders Johansen: FormatiØ sub	
  virial	
  speedの衝突でmmゴイ に成長する。	
  anuscript no. KWJ_AJ_Pebble_arXiv

Particle size distribution for a cloud collapsing into a planetesimalwith a solid radius of 100 km at Pluto’s distance from the Sun.
1.2

1000 rep. part.
200 rep. part.

1

Pebble mass fraction

Size distribution (fraction of total mass)

Fraction of total mass in pebbles as a function of solid radius of the planetesimal.
For simulations with, initially, cm-sized pebbles.

0.8

0.6

0.4

0.2

0
1.2

time forblue andbles. The
sions aref E. The

0.1

1

10
Solid radius Rsolid (km)

100

1000

t = 0 yrst = 5.1 yrs
t = 23.6 yrs
t = 25.2 yrs

1

0.8

0.6

0.4

0.2

0
10-5

10-4

10-3

10-2
10-1
Particle size (cm)

100

101

102

>1	
  mmのpebblesとそれ以下の割合。	
 Fig. 7. The particle size distribution in a collapsing pebble cloud with

Fig. 6. Pebble mass fraction in the resulting planetesimals for simulations with, initially, cm-sized pebbles, as a function of the planetesimal
radius, Rsolid . Particles are defined as pebbles if they have radii a > 1
mm. For more massive planetesimals the collision speeds are larger andthe collisions can then result in fragmentation, decreasing the pebble

solid radius Rsolid = 100 km at four di↵erent times. The collision speeds
are initially high enough for fragmentation to occur. Later in the collapse the particles move with sub-virial speeds (see Fig. 5) and theamount of fragmentation decreases. The low collision speed in the coldon the fr(Rsolid ⇠kof the pefor smallcollapsethe pebb
2009).
The
after a dinew (smamass clo
lision sp
tation. Fof the cohas timeof the cothe partic
that theyon the ot
beginnincollapse
the amouWhelan, E.T. et al.: X-Shooter
Spectroscopyof ISO-ChaI
217 dwarf	
  candidate	  ISO-­‐ChaI	
  217	
  Accre)on-­‐ejec)on	
  connec)on	  in	
  the	  young	
  brown	
  Whelan	
  E.	
  T.	
  et	
  al.	  (2014)	
  Ø Brown	
  dwarfのou•lowが10例ほど見つかってきたので、よくわかっているYSO	
  ou•lowを比べる。	  Ø kno0y	
  jet,	
  molecular	  &	  op)cal	
  ou•lows,	
  high	  &	
  low	
  velocity,	
  asymmetricなど共通点がある。	
  Ø YSO	
  ou•lowと比べるには空間分解が必要だが、BD	
  ou•lowは暗いので難しい。	
  Ø Mdot,out/Mdot,accを比べるとYSOより10%大きい(2009a,	
  2014a)が、観測 イア がある。	
  Ø jet	  emsisionが最も強いChaI	
  217で ォア。	
  Ø Chamaeleon	
  I,	
  d~140	
  pc,	
  M6.25,	
  80	
  Mjup,	
  0.64	
  Ro,	
  Lbol=0.023-­‐0.028	
  Lo,	
  5-­‐6	
  Myr	  Ø ou•lowはPA=20°,	
  red側が明るく、速く、電子密度が高い。	
  Ø i=45°,	
  4×10–6	
  Moの円盤やPA=238°	
  ,	
  5	
  AUにcompanionàou•lowの非対称を説明。	
  Ø VLT,	
  INAF	  X-­‐Shooter	
  GTOの一環。seeing	
  0.8”,	
  slit	  width~1.0”,	
  R=5000-­‐9000。	
  Ø absorp)on	
  doublet	
  NIとKIをBT-­‐Se0l	
  modelと比べて	
  Teﬀ=2940	
  K,	
  v	
  sini=20	
  km/s,	
  log	
  g=3.5-­‐4.0ß質量と半径に一致。	
  Accre)on-­‐ejec)on	
  connec)on	
  in	
  the	
  young	
  brown	
  dwarf	
  candidate	  ISO-­‐ChaI	
  217	
  Whelan	
  E.	
  T.	
  et	
  al.	  2014	
  (続き)	
  Ø [SII]のPV図では長さ1.6”のblue,	
  長さ1.0”にredを検出し、ou•lowを空間分解できている。	
  Ø blue側は3成分ある。2つの[SII]で検出されているので雑音ではない。	
  Ø 1つを除いて速度的にbackgroundでもない(青破線)。	
  Ø blue側は遠ざかるほど遅くなっている。YSOでもambient	
  gasとの相互作用で遅くなる。	
  Ø [OI]は臨界密度が高く、中心部だけを見ていて、空間分解はできていないが、Spectro-­‐astrometric	
  analysisからou•low起因とわかる( )。	
  Ø [NII]はYSOでは高速度
ゴ として知られており、blue側でのみ
Whelan, E.T. et al.: X-Shooter Spectroscopy of ISO-ChaI 217
検出。red側はイオ 化率が小さいということ。	
  -5 +/- 5-20 +/- 5-15 +/- 5-24 +/- 5-20 +/- 5
33 +/- 5-25 +/- 5

32 +/- 5

Accre)on-­‐ejec)on	
  connec)on	
  in	
  the	
  young	
  brown	
  dwarf	
  candidate	  ISO-­‐ChaI	
  217	
   Whelan, E.T. et al.: X-Shoote
Whelan	
  E.	
  T.	
  et	
  al.	  2014	
  (続き)	
  Ø 13個のaccre)on	
  tracerからLaccを求め、log	
  Mdot,acc=–10.6	
  Ø [SII]からne,	
  [OI]/[SII]からTe,	
  [NII]からxeを求めた。	
  Ø Mdot,outには[FeII]を使った減光補正が必要だが未検出なので、 .
Avに幅を持たせて推定。	
  Ø Mdot,out/Mdot,acc=0.05-­‐0.2はCTTSやそのmagneto-­‐centrifugal	
  jet	
  launching	
  modelsと合う。	
  Ø 2009aでは>1だったが、不定性の大きい臨界密度を仮定した Fig. 7. Mass accretion rate estimated from various accretion tracers. The solid line gives the mean value of log ( Ṁ ) and the
方法を使ったからで、今回は臨界密度を使っていない。	
  dashed line is the ± 1 uncertainty. The errors on the individual
mainly come from the error in the line fluxes andØ HIも検出され、その起源を調べるためにBalmer	
  decrements measurementsin the value of the on-source extinction. The value of Ṁ calculated by Muzerolle et al. (2005) from modelling of the H↵ line
and from the UVES spectra presented by Whelan et al. (2009b)を計算した。	
  are included here for comparison. No error estimates were givenØ 複数のnupとHβのク の比をTe,	
  neから計算される値と比in Muzerolle et al. (2005).
べ、Te=104	
  K,	
  ne=1010	
  cm–3となり、accre)on	
  ﬂow起源と思われ(2009b) the method used assumed a value for the critical density (n ) which was highly uncertain. Here an estimate of nる。	
  is not required. Estimates of Ṁ are more accurate due to the
ISO-ChaI 217

acc

acc

cr

cr

acc

use of several accretion tracers. Previous to this work it has only
been derived from the H↵ line. However, this change in the derived value of Ṁacc does not have a significant e↵ect on the final
values of Ṁout / Ṁacc . Similarily extinction does not have a largee↵ect on our conclusion in this case. As for all values of Av up
to the value measured for the driving source Ṁout / Ṁacc remains
within the limits of predictions by leading models.
4.2. Investigating the origin of the HI emission

Several HI emission lines are detected in the spectrum of ISOChaI 217 (see Figure 4). To investigate the origin of this emission the Balmer decrements were computed with respect to Hline and are plotted as a function of upper quantum number (nup )in Figure 8. While the spectral range of X-Shooter covers mostof the Balmer lines any lines with nup > 13 were found to be too
noisy for inclusion in our analysis. Additionally, at the interme-