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のデータから点源を決定し、カタログを作った。	
  Ø 観測パラメーター、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’グリッドで6x6で、Fig	  1のように傾けてマッピング。	
  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
allPlanevalues
2. FocalDesignsmaller than ⇠0.100 . The astrometrico↵setswereappliedThe ACIS focal plane (Figure 2.1) is composed of two arrays with one array designed to optimize imaging and the otherarray forto the L1 data products to create new aspect solutions forspectroscopy
utilizing theobjective gratingseach ObsID and the data were then re-reduced followingthat are part of the
Chandra X-ray
the procedure in Section 2.3.
Observatory

Chandra Cygnus OB2 Legacy Survey Catalog

ACIS Focal Plane Array

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.

Dimensions are in inches [mm]at operating temperature.
Datum dimensions apply to the active area,other dimensions apply to the silicon.
Dimensions marked * apply to theouter edges of the array.
Spacecraft coordinates are used withthe origin of the on-axis focal pointof the HRMA when the SIM is at IP#1.

(CXO). In the caseof the imaging
array, the CCDs
are tipped to lie
tangent to the
optimum focal
surface of the
Wolter Type I
mirrors employed
for imaging. This
nearly spherical
focal surface lies
10.04 meters fromthe joint separating
the paraboloid andhyperboloid
mirrors and has aradius of curvatureof 85 mm. By
Figure 2.1
tilting the CCDs in
the imaging array more pixels encompass regions of high angular resolution (see Figure 2.3 for a comparison between a

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.

IP#1

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 of.078 [1.99].057 [1.46].000 [0.00]

Fig. 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
4 grid pointings. The broader bases correspond t
by slight variations in the roll angles of the ObsI

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...	
  異なるグリッドに対してoverlapがあるので、合わせて解析したい。グリッドが違うと1つのピクセルでの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メンバーも他から加えた。	
  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%コントアーとして抽出。	
  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。	
  Chandra Cygnus OB2 Legacy Survey Catalog
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を半減するようなデータを取り除く。)	
  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。	
  1Wright et al.

Chandra Cygnus OB2 Legacy Survey Catalog

ig. 7.— Distribution of X-ray sources in the Chandra Cygnus OB2 Legacy Survey colored as a function of the lowest o↵-axis angle thatsource is detected at. The observational tiling strategy is evident in the grid pattern seen in this image. Due to the tiling strategy usedvast majority of sources were observed at least once at an o↵-axis angle <40 and therefore with PSF sizes <200 in radius (at 1.49 keV).

is techniques were introduced to maximise the scienfic return of our unique set of observations, including
enhanced version of the CIAO tool wavdetect thatrforms source detection on multiple non-aligned X-ray

13

Fig. 8.— Distribution of various source quantities for the 7924 sources in the final X-ray catalog. Left: The distribution of net counts
all extracted sources, with quartiles at 7, 13, and 28 net counts. Center: The background-subtracted median photon energy of all extracte
the low- (Flaccomio et al. 2014b) and high-masssources,(Rauw
with quartiles at 1.7, 2.0, and 2.8 keV. Right: The distribution of model-independent source energy fluxes calculated using th
et al. 2014) members of the association is presented elseequation Fenergy = 1.602 ⇥ 10 9 Ẽphoton ⇥ Fphoton (see Section 5.2) with a median of 2.6 ⇥ 10 15 erg cm 2 s 1 .
where, along with analysis and discussion.

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バンド幅が狭かった(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”のエンベロープ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.

3

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 theTable
1. Keyparametersof the sources.
Ae/Be
discs
with and without
respectively,by theneutralcarbonatom coincides
with theThe
firstHerbig
sidelobeof theneighbouring
CO line,flaring,
but hassignificant are
fluxlistedof its first,
own followed
before background
T Tauridiscs. Within
the ordering
is by
spectral typefromearly to late.
subtraction(grey), eachwhilegroup,
its strength
is reduced
considerably
afterbackgroundsubtraction is applied (black; same spectrum as in Fig. 2).
6(b) Zoom of the 60 GHz section of the HD 163296 spectrum centred around the [N ii] line, again before and after background subtraction.
4+
This Object
shows that
N line towards
HD 163296
originates
extended [ref.]background2
namethe(alternative)
Spectral[ref.]
Cat.(a)from theDistance
Disc medium.
diameter [ref.](b)Flaring(c)
0
type
(pc)
(au)
(y/n)
HD 100546
B9 Ve[1]HD 179218
B9e[3]
HD 97048 (CU Cha)
B9/A0e[1]
AB Aur (HD 31293)
A0 Ve
[1]
HD 36112 (MWC 758)
A5 IVe[3]
HD RAS,
169142(MWC 000,925) 1–16 A5 Ve[3]c 2014MNRAS
HD 100453
A9 Ve[3]
HD 142527
F6 IIIe[11]
HD 50138 (MWC 158)
B6-7 III-Ve [14]HD 163296 (MWC 275)
A1-3 Ve[16]
HD 104237 (DX Cha)
A4 IVe[3]
HD 144432
A9/F0 Ve [18]RNO 90G5[19]
RY TauK1[19]
DR TauK5[23]
TW HyaK7[25]
FZ TauM0[28]
VW Cha
M0.5
[29]

HAeBe
HAeBe
HAeBe
HAeBe
HAeBe
HAeBe
HAeBe
HAeBe
HAeBe(d)
HAeBe
HAeBe
HAeBe
T Tauri
T Tauri
T Tauri
T Tauri
T Tauri
T Tauri103±6
240+7040
180+3020139±19205⇠145
122±10145±15500±150
119+1310
115±5160+36
25
125
140
140
51±4
140
180

[1]
[1]
[1]
[6]
[6]
[3]
[6]
[12]
[6]
[6]
[6]
[6]
[20]
[21]
[21]
[26]
[21]
[30]

700
240±40
⇠750
900
⇠330
470±10
⇠100
⇠600
>30
⇠1000
120±40⇠140
180–380400
200–600.30

submm [2]model [4]mid-IR [5]submm [7]submm [8]mm [9]
optical [10]submm [13]mid-IR [15]submm [17]model [4]mm [22](sub)mm [24](sub)mm [27]mm [24]near-IR [31]

y
y
y
y
y
y
y
y
n
n
n
n
-

3
2
1
0
3
2
1
0
3
2
1
0

CO

1480

CO

1470

13CO

1460

CO

1450

13CO

1440

(GHz)
Warm carbon monoxide frequency
in protoplanetary
discs

13CO

830

CO

800
810
820
frequency (GHz)

13CO

790

AB Aur

CO

780[N II] 3P1-3P0

18

CO

1.0

20

[C I] 3P2-3P1

HCN
13CO

CO 7-613CO

1.5

HD 100546[C I] CO

22

6
4
2
0
6
4
2
14900

13CO

24

2.0

CO

2.5

M. H. D. van der Wiel et al.

13CO

6

26

CO

28

3.0

CO

flux density (Jy)before backgroundsubtraction
after backgroundsubtraction

30

13CO

32 (b) HD 163296(a) DR Tau

3.5

flux density (Jy)

4.0

HD 97048

HD 163296

DR TauFZ Tau

500

600

700

800

900
1000
1100
frequency (GHz)

1200

1300

1400

1500

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,
Notes. a Category based on spectral type of central star. b Disc diameter is taken from resolved imaging at millimetre or submillimetre
mainly visible near the stronger lines, e.g., the highest frequency CO lines in the top panel. The broader spectral features visible in thefor all observed transitions, (iii) all probed gas can be characterized by a single temperature, and (iv) the emission lines
are opticallythin. Forrotational linedtransitionsof CO with	
  Herschel	
  SPIRE	  Signatures	
  of	
  warm	  carbon	
  monoxide	  in	
  ppurerotoplanetary	
  iscs	
  observed	
  on accuracy
discs, typically only the first of these conH.	
  D.	
  vinan	
  
der	
  
Wiel	
  ein
t	
  aprotoplanetaryl.	
  (2014)	
  続き	
linesM.	
  detected
our
Herditions holds. Nevertheless, the rotation diagram provides ale10 3)M.have
also
been
H. D. vander Wiel
et al. obuseful tool to gauge variations between di↵erent disc sources.
amely CO 6–5, CO 7–6 andrespectively,
g is thestatistical weightof state‘u’, A isThe
verticalandhorizontal
axes12y (a function of the(a) HD 100546
24
the Einsteincoefficient for spontaneousemission, ⌫ is theØrota)onal	  diagramを用いてCOを解析。	
  APEX(Panić
et
al.
2010)
frequency related to the energy di↵erence between states ‘u’integrated
line
intensity
F
for
each
transition
from upper
ul
22
and
‘l’,
Q
is
the
partitionfunction,
and
d
is
the
distancechel20(FedelePACSet al. 2013b).
to theßLTE,	  source. If all theビabove
assumptions hold, a determiØlevelームより小さい、単一温度、光学的に薄い、を仮定(LTEし‘u’ to lower level ‘l’), and x (upper level energy of thatnation of a will yield the rotational temperature T . Likety fromthe
APEX
meawise,c
would
be a measure for the total number of emitting
18 R
transition)
can be related through:か成り立たないが便利)。	
  molecules
N , but only if the beam filling factor f in Eq. (1)GHz,16 Tmb dv = 17.7±0.9
is eitherassumedtobeuniform✓◆or can be constrained using✓◆
high angular resolution observations. Since this paper does
2
SPIRE 18
4⇡Finulthis section is restricted to1the slope,
Eup
fN
14
neither, the analysis±0.7)⇥10
W m , using
lnof the rotation diagrams.
=
+ ln
,
(1)
a,
0
500 1000 1500 2000 2500 3000 3500
Figures
diagrams of T
three
gu4a,b,c
Aulshowh⌫rotation
Q d2
his value
is consistent with
rotHerbigk
ul
Ae/Be sources with at least nine CO line detections in Her(b) AB Aur24|{z} | {z } |{z} | {z }schel SPIRE (this work) and Herschel PACS (Meeus et al.
.7±2.3)⇥10 18 W m 2 to2013) combined,yaugmented by ground-based a
datapoints inx
22
c
PACS
the (sub)millimetre gathered from the literature (see Apg into
account
theabsolute
20
pendix A). Note that some of the CO rotational lines probedØwherePkACSも合わせて、高温低温に分けてフィット。	
  here Hershel	  may be opticallythick.
Thearehigh opticalis corinstruments.
h andthedepthPlanck
and Boltzmann constants,
18
roborated by predictions from numerical models (e.g., BruduulTrot = 215 K

rot

Trot = 879 K

Trot = 78 K

16

ground-based

ln

4 Ful
gu Aul h ul

(m 2 )ln

4 Ful
gu Aul h ul

ground-based(m 2 )

ulTrot = 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 HD	
  100546はPACSだけで2温度とわかっていたが、SPIREの部分は
c 2014 RAS, MNRAS 000, 1–16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)。	
  erer 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 SPIRE

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 lines

in a model of the HD 100546 disc. Solid outlines represent areasØ 2天体について他の波長でフィットしたモデルと12CO	
  SLED	  (輝線のSED)を比べる。	
  delimited by the cumulative 15 and 85% line flux contribution

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 refinemof each transition, both vertically and radially. The grey-scaleØ HD	
  100546は概ね合っているが、モデル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 モデルはNIR,	
  mmの連続波をフィットしたもの。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バンダンスが減るなどのisotope	
  selec)ve	
  processが必要。	
  by SPIRE and PACS are optically thick, the shape of theØ HD	  163296のモデルは可視からmmの連続波と低励起COをフィット。	
  rotational SLED is sensitive to adjustments of the gas temperature in the relatively tenuous, irradiated upper layer ofØ 2−4倍モデルが弱い。100	
  AU以内がモデルより高温のはず。	
  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).

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、またフレ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やフレアしていないHerbigは強かった。	
  Ø 12CO	
  10-­‐9は12CO	
  18-­‐17,	
  [OI]3P1-­‐3P2と相関。サンプルは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 curve

are the minim
Eﬀects	  grain	
  growth	
  on	
  the	
  interstellar	
  polariza)on	
  
urve	
  
the extinction
Theof	
  wavelength
dependenceof polarization
P c(λ)
in the visible
Nikolai	  .	
  Voshchinnikov	
  &	
  Hiroyuki	
  byHirashita	  (2014)	
   formula suggestedon the particle
partVofspectrum is describedan empiricalradiation (Bohby 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。	
   through ϕ, ω,P (λ)/Pmax = exp[−K ln2 (λmax /λ)].
(1)
between them,Ø 特に天の川銀河ではKとλmaxに経験則がある。	
  the distribution
This formula has three parameters: the maximum degree of polar-

ization Pmax , the wavelength corresponding to it λmax and the coØ 部分的に整列、非球体、減光断面積、を考えたモデル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
2.2.1 Size dis
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.
HV14 calculatØ 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).
accretion and(1)モデルでは>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。	
  The parameter K determines the half-width of the normalized(2012). Silicat
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
grain species.
縮しない。	
  The(2)観測ではRvが増えるとAv/NHが急に減るが、モデルでは減らないàsilicateが~0.03	
  µmgrow
W = λmax /λ− − λmax /λ+ ,(2)
mined by the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,(assumed to be(nH ) density,
for accretion (Z = Z⊙ , nH
HV14 comotion as a fu

rV + δ0, j

rV, cut = 0.20 µm, λmax = 0.628

whereEﬀects	  of	
  grain	  growth	
  on	
  the	
  interstellar	
  polariza)on	
  curve	
  Nikolai	  V.	
  Voshchinnikov	
  &2	
  Hiroyuki	  irashita	  (2014)	
  続き	  PolarizationH
curve
and grain growth
3
κB
23
IDG
δ0, j = 8.23 10

1/2

nH Tgas Td

µm.

(9)

3

RESULTSWe made calculations for p
spheroids consisting of silicate
In the IDG mechanism, smaller grains which are powerful pocles of 59 sizes in the range fromlarizers are aligned better than larger grains (see Eq. (8)). In order
1 µm were utilized. The extinc
to reduce the polarizing efficiency below the observed values, it isbeen calculated for 77 waveleng
possible to increase the minimum cut-off in grain size distribution
to λ = 5 µm. We next determ
rV,min,j (e.g. Das et al. 2010) or to assume that small particles are
curve Pmax , λmax , and K. In t
randomly oriented (e.g. Draine & Fraisse 2009).
ized curves P (λ)/Pmax and the
Because the grain size distribution is fixed in our model (seepolarization curve and the posi
r ,µm
r ,µm
Section 2.2.1)
we suggest the following
modified IDG alignment
acteristics of the Serkowski cu
function with reduced alignmentof small grainsalignmentを、小さい粒子の整列を減らすよう(see also Mathis
Ø 整列の仕方はimperfect	  Davis-­‐Greenstein	  size distribution nSi (rV ) and th1986): v,cutというパラメータを導入)。	
  に修正(rcate particles and weakly depen
The evolution of silicate grain size distribution for the tuned
Using the size distributions obtained for silicate and carbonaceous dust, HV14 calculated the evolution of extinction! curve and model is shown in Fig. 1 (right panel).
" Compared with the dust dis3
tributions in the original model without tuning (Fig. 1, left panel),found that the prominent variation of extinction caused by increastion (Ω) of the particle orientatio
fj (ξ,
β, ...) = 1 − exp(−r/rfurther,) × fIDG,j(ξ, β) , (10)
Vgrow
V,cut
the grains
even beyond 1 µm at > 100
Myr. As showning accretion occurs
at < 30 Myr. After that, the evolution of
IDG
Ilater, silicate grains in the large-size tail of size distribution are ofextinction
curve is driven only by coagulation, which flattens theum. Notethat theparameters:
δ
=
0.5
µm,
δ
0,
Si
0
crucial importance in explanation of observed polarization.
extinction curve and makes R larger. Observations indicate that
However, this facteven for large whereR ≈ 4.4, the carbon
bump aroundis clear
We use size distributionTheof silicatealignmentand carbonaceous grainsfunction (10)r1V,cut
is0.22aµmcut-off
parameter.
y spheroidal
and the ones
UV extinction rises with a positive curvature. However, theobtained by HV14 but replace spherical particles by spheroids ofcurves,e.g., thetheoreticalpredictions show that at large R (i.e. after significant
the same mass. The uncertainties of grain shape to spheroidal has
a smoothswitchnon-alignedalignedgrains in comcoagulation atgives
T > 100 Myr),the carbon bump disappears.
This fromonly a minor
influence on accretionto
and coagulationcompared
implies that coagulation is in reality not so efficient as assumed inwith the change of S and S. Therefore, that we simply use
3.1 Wavelength dependence
abruptjump
suggestedby Draine& Fraisse (2009).
the model for parison
carbonaceous dust.with
There is anotherdiscrepancy:
theHV14’sresults with an assumptionof a = r .
observed A /N tends to decrease drastically as R increases,
while it does not decreaseThoughsignificantly after pronounced
coagulathe silicate
and carbonaceous particles comparablytion at T > 100 Myr.This is because silicate stops to coagulate at
3.1.1 Linear polarization
2.2.2 Materials
r that∼ 0.03µm, which is still too small to affect the UV opacity.
gned sotheir
0.5
contributetoextinction,thepolarizationisassumedtobeproduced
In summary, the observational data indicate (i) that carbonaJones (2012a,b,c) published a comprehensive analysis of the comotation) andtheir
ceous dust should be more inefficient in coagulation than assumed,
positional properties of hydrogenated amorphous carbons. It intion ofandthe(ii)magFirst of all, we calculated the ponlyshouldbygrowsilicate
Such
an assumption
hascludesthe considerationof the evolution of carbon
solids beenfrom hy- earlier donethat silicate
beyond r ∼particles.
0.03 µm (thening angle
growthofisthe
limited by the coagulation threshold velocity). Theredrogen rich aliphatic grains (a-C:H, sp bonded) to hydrogen poor
fore, Davis–
for a better
observational&
data,Krügel
HV14 also tried models
aromatic
grains(a-C, sp bonded),
calculations ofandthe complex
remperfect
the initial size distribution of
byfit theChini(1983)fractive
andMathis(1986),
recently
has gotindexes as a function of the material band gap, and inwith S< 1 for carbonaceous dust and no coagulation threshthe distribution
for silicate dust as a tuned model. Such a tuning is acceptablesize-dependent properties. Later, Jones et al. (2013)
we assume that grains of all si
additional support in the vestigation
workofof
Voshchinnikov et al. (2012) who
ntationold
parameter
if we consider uncertainties in grain properties. The removal offound refractive indexes of amorphous silicates using the mixture
ong & theGreenbergof amorphous forsterite with different amount of metallic iron (acoagulation threshold may be justified if the grains have fluffy
mechanism (rV, cut = 0 in Eq.
a 0correlationinterstellar
polarization
structures thatfoundabsorb the collisionenergy efficiently, and/orbetweenareSil ) andtheKöhler etobservedal. (2014) added iron sulfide
in this mixture.
coated with sticky materials such as water ice (Ormel et al. 2009;
In modelling, we choose the refractive index ’silicaHirashita (7)& Lidegree2013).
for amorphouswith grains.
10% volume fraction
UV (λmax < 0.2 µm). In ord
and the abundancete1 FoFe10.RFI’of silicon
insilicates
dust
Polarization in
In the tuned model, the coagulation efficiency was decreasedof Fe (equivalent to ∼70% of the cosmic Fe) and the data for an aλ,µmby a factor ofIR
2 by adoptingS= 0.5
for carbonaceous
dust,
C(:H)
with aof
band gap
E = 2.5 eV. Thepopulations
optical constants
Serkowski curve (λ= 0.55featuresalso
supports
thematerial
idea
separateof polar1

S ilicate

T=0 M yr (initial)
T=3 M yr
T=10 M yr
T=30 M yr
T=100 M yr
T=200 M yr
T=300 M yr

0.01n(r V )/n m ax

0.01

n(r V )/n m ax

1

S ilicate0.0001

1E-006

1E-0080.0001

1E-006

1E-0081E-0101E-010

0.001

0.01

0.1

1

0.001

0.01

0.1

V

1

V

Figure 1. Evolution of grain size distribution for silicate dust. The grain size distribution is normalized to the maximum. Effects of evolution of grain size
distribution are shown for original (left panel) and tuned (right panel) models at T = 3, 10, 30, 100, 200 and 300 Myr. The initial condition is shown by thedotted line.

V

V

V

acc

V

H

v

P(λ)/Pmax

V

coag

V

V

3

rcut=0.05 µm

2

rcut=0.08 µm

coag

rcut=0.13 µm
rcut=0.20 µm

average Serkowski curve

Fe

0.2

coag

0.5

2.0

g

Eﬀects	
  of	
  grain	  growth	
  on	
  the	
  interstellar	
  polariza)on	
  curve	
  Nikolai	
  V.	
  Voshchinnikov	
  &	  Hiroyuki	
  Hirashita	  (2014)	
  続き	
  Ø サイズは0.001	
  µm-­‐1	
  µmに59グリッド、波長は0.2	
  µm-­‐5	
  µmに77グリッド調べた。	
  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というタイムスケールに矛盾しない。	
  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 growth6

Voshchinnikov and Hirashita
Taurus: cloud 1 [1] (14)Taurus: cloud 2 [2] (5)10.5

T=0 Myr
T=300 Myr
average Serkowski curve0KP(λ)/Pmax

P(λ)/Pmax

1.5

0.5

0.5

1

λ, µm

2.0

0.5

0.3

T=0 Myr
T=10 My
T=30 Myr
T=50 Myr
average Serkowski curve

0.5

1

0.12

1.0

0.1
0.2
0.08
0.15

0.12

0.06

0.06
0.1

0.05

0.05

T=0 Myr
T=300 Myr

0.5
0.2

0.5

0.2
0.15

0.08

0
0.2

0.3

T=30 Myr
T=40 Myr

0.2
0.15

0.1

CygOB2 A

T=20 Myr

1.5

0.12

1.0

T=0 Myr
T=10 Myr

Elias 22

ρ Oph cloud [1] (8)Cha I dark cloud [1] (22)
R CrA cloud [1] (13)mainly CygOB2 [3] (6)all sky [4] (30)all sky [5] (12)all sky [6] (5)Musca, Coalsack [7] (10)all sky [8] (13)stars with |b|>200 [9] (22)

K

1

5

0.08

0.5

2.0

λ, µm

Figure 3. Normalized wavelength dependence of interstellar linear polarization. Effects of evolution of grain size distribution on curves P (λ) are shown forthe original model without tuning (left) and the tuned (right) model. The average observational Serkowski curve (λmax = 0.55 µm, K = 0.92) is plottedusing the open circles. The theoretical curves were calculated for the model: prolate spheroids, a/b = 3, rV, cut = 0.13 µm. The parameter rV, cut for initial
curve was chosen in such a manner in order to be closer to observational curve. The characteristics of theoretical curves are: T = 0 Myr, λmax = 0.538 µm,
K = 1.07; original model: T = 300 Myr, λmax = 0.519 µm, K = 1.04; tuned model: T = 10 Myr, λmax = 0.588 µm, K = 1.05; T = 30 Myr,
λmax = 0.689 µm, K = 0.95; T = 50 Myr, λmax = 0.812 µm, K = 0.81.

0.06
0.05

0.2

0.4

0.6

λmax, µm

0.8

0.2

0.4

0.6

0.8

λmax, µm

Figure 5. Parameter K of Serkowski curve versus the wavelength λmax of the maximum linear polarization. The symbols indicate the observational datafor 160 stars in the Milky Way. The sources of data and number of stars are given in the legend in square brackets and parenthesis, respectively ([1] –Voshchinnikov et al. 2013; [2] – Whittet et al. 2001; [3] – Martin et al. 1992; [4] – Whittet et al. 1992; [5] – Wilking et al. 1980; [6] – Wilking et al. 1982; [7]– Andersson & Potter 2007; [8] – Clayton et al. 1995; [9] – Larson 1999). Stars (⋆) connected with solid and dashed lines illustrate the results of modelling fororiginal model without tuning (left panel) and tuned model (right panel). Model: prolate spheroids with a/b = 3 are adopted. The values of rV, cut are shownfor the initial size distribution (T = 0 Myr) and in the case of tuned model (T = 40 Myr). Horizontal and vertical dashed lines correspond to the average2GM R0 R
ii) A small particle, 2, ‘running’ into (co.
(6)
particle, 1: f < 0.1
R0R
Forma)on	
  
o
f	
  
p
ebble-­‐pile	
  planetesimals	
  
C
F
C
C
F
Karl	
  Wahlberg	  &	
  Aanders	  Johansen	  (2014)	
  is
We do a simple interpretation of 10
Fig5‘desired’change Jansson	
  in ⌘ after
dissipative
collision⌘. If | ⌘| < | ⌘max | the new value of ⌘ is equal to (2010), for collisions between compact pareqilibriumvalue, otherwise the new value is the old plusØ
惑星形成はmm-­‐cmのpebble(小石)にダスト・ice粒子が積もることで始まる。	
  2
10to6 virialize 10 4
10 2
100
102
0 4
100 for the system
102
since
there is10notenough time1Ø
重力的に束縛されたpebbleたちが衝突してエネルギーを放出し、収縮する様子を計算
v s(m1 ) s )f 0.1
Collision
Collision speed v (mme step
t.speed v (m s 1 )v < vstick
C
した。	
  
en
it
comes
to
updating
⌘
we
need
to
take
into
account
KomesØ
as 収縮にかかる時間と最終的なサイズ分布を求めたい。	
  a function of projectile radius, a, collision speed, v, and relative particle size,
vstickf ⌘ mv2 /m< 11. GreenBregions i
neticenergyisbothdissipatedinthecollisions
andreions bouncing and red regions fragmentation. Outcome regions taken from Güttler et1 al.
(2010).
v < 25F
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,	
  when the cloud contracts, so using Eqs. (1), (3) and (4)v 25Ffragmenta)onがある。	
  ainØ 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Here C denotes
bouncingi) Two similar-sizedparticlescoagulation,
colliding or B
a larger
part
できる。	
  
The athreshold
speed forT0
U
colliding withsmall particle,
1: fsticking,
0.1 vstick , is tak
2 T
2 EØT fragmenta)onの度合いで、できる微惑星のporosityと密度が決まる。	
  
=
⌘
=
⌘
K
K
(2010),2,as‘running’ into (colliding with) a
RT 2
ii)
A
small
particle,TT00porousでrockが多い。	
  (6)0.Ø 密度はiceが多いと低くなり、つまり重いとless	
  1particle,
1: f < 0.1
BBB
CCC Jansson and Anders Johansen: Formation of pebble-pile planetesimals
Karl Wahlberg
s
BBB
CCC
E
2
2 B
BB in ⌘+ after⌘ aCCCC ,dissipative collision is5⇡a0 Froll of Fig. 11 in Güttler
We dointerpretation⌘ange(7) a simple
KB
BBB E0C
2v=,
CCC value of ⌘ is equal to (2010), for collisions
stick
⌘ new
betweencompact
particles:<B@|{z}| ⌘max | the
µ
|{z}A
>0
e, otherwise
the<0 new value is the old plus
not enough time for the system to virialize
where µ =1 m1 m2 /(m1 + m2 ) is the reduv(ms)
f 0.1 f < 0.1
E is the energy lost in the collision (see Section 2.2), U monomer radiusand Froll is the rolling for< vstickCCravitational
releasedis the previously cal- Thev coagulation
updating ⌘Kenergy
we needto takeand
into⌘accountoutcome of collisions wi
BB
1  v<1in ⌘ (could
be collisions⌘max ). Theandfirstre-term in Eq. (7) svsticks changeboth dissipatedin the v < 25 m s 1 is explained by mass t
1  v < 25FCes
the kinetic(energydissipationd contracts,
so energy
using Eqs.
(1), (3)and (4) in collisions) (Wurmet al. 2005). A map of collision outc
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 andred regions fragmentation.
regions takenenergy
from Güttler et al.
(2010).
v 25
Fhesecondterm
increases
theOutcome
kinetic(releaseof is shownin Fig. 1. F
ional potential in the collapse). If the virialization is too
Next we determine the amount of energi) Two similar-sized particles colliding or a larger particle, 2,Large projectile or similar sized particles: f

0.1

Large target: f < 0.1

10 1

B

10 3

B

vstick

10 3

vstick

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 )2

1

Projectile radius a (m)Projectile radius a (m)10 1

Forma)on	
  of	
  pebble-­‐pile	
  planetesimals	
  Karl	
  Wahlberg	
  Jansson	
  &	
  Anders	
  Johansen	
  (2014)	
  続き	
  Ø モデルで非弾性衝突による収縮時間とサイズ分布を調べる。	
  Ø 初期のpebblesはmm,	
  cm程度のmonodisperseなサイズ分布。	
  Ø 最終的に微惑星から冥王星サイズRsolid=1-­‐1000	
  kmができる質量で、カイパーベルトと太陽を想定し、ヒル半径に広がるpebbleの集まり(雲)から始める。	
  Ø pebblesはSiO2の密度の一様球でµmサイズの粒子からなるモノマー。	
  Ø 速度は系全体の運動エネルギーから決まるMaxwell分布。	
  Ø 衝突は0次元、回転無し、representa)ve	
  par)cle	
  approach	
  (Zsom	
  &	
  Dullemond	
  2008)をモンテカルロで統計的に計算。	
  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 = 3EcollrollFig. 2. In our model pebbles are composed of µm-sized monomer particles. In collisions resulting in partial fragmentation monomers are removed0
one by one to simulate erosion. The binding energy of the fragment, Ebind, is larger (closer to zero) than the binding energy of the original particle,
Ebind , since the fragment consists of fewer monomers. The number of monomers released, Nmon , is equal the amount of contact surfaces thecollision energy can break divided by three contact surfaces per monomer, Nmon = Ecoll / (3Eroll ).

0s

course not completely physical but it is a first step towards a
Hereis the initial
total kinetic
energy(the sum
kinetic
moreT 0complete
model.
We alsoassumethat ofthethecloudis in, orof
each
individual
particle
with
mass
m
and
speed
v
),
M
i
i
e energy
strivesto get into,
equilibrium.pThe
cloud is initially set inForma)on	
  of	
  virial
pebble-­‐pile	
  lanetesimals	  the total
mass of thecloud,
vvir,0is thevirial speed. isvirial
equilibrium
sopebblethat theinitialtotalenergy,E0 , iscalculated
fromtheinitialkinetic
energy,Uistheinitialpotential (2014)	
  続き	
  
Karl	
  Wahlberg	
  Jansson	
  &	  A
0 nders	
  Johansen	
  .
energy, G is the gravitational constant and R0 is the initial radius
l of the cloud. During the collapseNwe can follow
the cloud’s evo1X- lution through three parameters:
2 ⌘ .1The parameter
2
⌘
,
⌘
and
⌘eq
eq mi vi = K Mvvir,0E0 = T 0 + U0 = T 0 =s is defined as the initial energy2divided by the current
2
total energy Ø 時間変化するパラメーターとして3つ。(左式)	
  i=1
p E(t),
Ø エネルギーが抜けて–∞になったら終わり。	
  1
3 GM 2h= U0 =
.
(1)
2
10
R
g
Ø
最終的な半径Rsolidと収縮にかかる時間tcollの
0
E0
.
(2)
- ⌘eq (t) ⌘E(t)関係をプロット。	
  e Here T 0 is the initial total kinetic energy (the sum of the kineticof each individual
particlemass
mi R(t),and speedi ), M
The parameter⌘ is definedas thewith
current
size,of the vØ) energy
pebblesが初めcmサイズのとき3種類ある。	
  is thedivided
total mass
the pebble
by theofinitialsize, cloud, vvir,0 is the virial speed cal- cloud(1)Rsolid<50	
  km	
  f culated from the initial kinetic energy, U0 is the initial potential
G is the gravitational constant and R0 is the initial radius
bouncingが支配的でtcoll∝Rsolid–1	
  - energy,R(t)U0of the⌘ cloud.
= During. the collapse we can follow the cloud’s(3) evo- ⌘(t)RU(t)0(2)50	  lutionthroughthreeparameters:
⌘
,
⌘
and
⌘
.
The
parameter⌘eq km<Rsolid<70	
  km	
  
eq
K
r
initialdivided by
the current
Finally ⌘as
which
we energy
use to generate
relative
speeds total
whenenergy
K ,thefragmentaionが効く。重いとfragmenta)onが早く、
n is definedthecloud is not in virial equilibrium, is defined as the initial kiE(t),頻繁に起こる。	
  
解析的にはpebblesのサイズをaに固定すると
E0⌘eq (t) ⌘T 0.
(2)
–1
⌘K (t) ⌘ E(t).
(4) tcoll∝a/Rsolid 。	
  T (t)The parameter ⌘ is defined as the current size, R(t), of the
The cloud is bound and contracting, so ⌘eq and ⌘ have val- (3)Rsolid>70	
  km	
  
cloud
dividedby1,theinitialues
between
0 andwhile⌘K size,can have any value since kinetic 衝突がとても頻繁に起こり、自由落下時間が収energy is both dissipated in collisions and released in the gravi縮。密度は同じなので自由落下時間も一定。	
  tational collapse.
These
R(t)
U0 three parameters do not need to have thenetic energy divided by the current kinetic energy, T (t),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

1000

Simulations (cm)
Simulations (mm)Simulations (dm)Power-law fit (cm), α = -1.05
Power-law fit (mm), α = -0.97
Free-fall time of initial particle cloud

Collapse time tcoll (yrs)

f
n
o
,
w
t

⌘(t) ⌘= the cloud
.
(3)
same
value sincedoes not virialize immediately afterRU(t)0a collision.
We
assume
that
the
cloud
virializes
onatimescaleof100
the free-fallthe cloud.
At atocertain
size relativeof the cloud,⌘, whenØ 初めmmのときはfragmenta)onが効く前に、自Finallytime⌘K , ofwhich
we usegenerate
speedsf we canget the maximum possible change in ⌘, ⌘max , by assumthethat
cloud
is not in virial equilibrium, is defined as the initial ki-由落下で決まるフェーズになる。	
  
- ing
the
10 cloud has been free-falling from R0 . This givesefe

netic energy divided by the current kinetic energy, T (t),rr
R
1
2GM R0 R
1 dR
1 max
T0 = 1
10t =
100
1000
10000
⌘
= 0.1
t 100000⌘Kmax(t) ⌘ R0 . R0 dt ↵ Solid radius
R0 Rsolid (km)R0
R(4)

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、次の衝突までにビリアル平衡(η=ηeq=ηK)にならなくなり、自由落下速度で収縮する(エネルギー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

1

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

Pebble mass fraction

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.
Fractionof totalin pebbles as a function of solid radius of the plane
blue lines the equilibrium values (⌘eq ⌘ E0 /E) and Virial
greenandlines
thespeedkinetic(⌘K ⌘asTfunction
pink line shows the free-fallcollapse
ofmassa cloud
0 /T ). The
Values for three different planetesimal sizes: 5, 50 and 500 km
simulations with, initially, cm-sized pebbles.
and the teal lines the analytic solution from Appendix
A. Note that time is measured in units of the collapse time. Therefore the free-fallForlineisdi↵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 is1000 rep. part.
Virialthe
speed,
5 km
200 rep. part.
1 on the other hand, hasin virial equilibrium until the very end of the collapse and follows
analyticalsolution (⌘ = ⌘K = ⌘eq ). The massive cloud,Actual speed, 50 km
initially higher collision speeds which lead to pebble fragmentationand50rapid
Virial speed,
km energy dissipation. At about ⌘ = 0.5 the energy dissipation becomes
Actual speed,before
500 km the next pebble collision (the first term in Eq. (7) dominates over the
faster than free-fall, so the cloud does not
10 get into virial equilibrium
0.8
Virial speed, 500 km
second). This results in a transition to cold collapse (⌘ is parallelto the free-fall curve) with sub-virial speeds (⌘K > ⌘eq ).
0.6

between the primordial pebbles. Collision
speeds only reach
0.4
values high enough for fragmentation to occur in the very
0.1
end of the collapse. The power-law fit shows us that tcoll /0.2
1Rsolidwhich 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 of0
constant 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
0.1
1
A). We findthat thecollapse timeincreases linearlysilicate pebbles with a monodisperse size distibution. The peb-Time t (t pendix
coll)
with
pebble
size,
which
also
agreeswiththeanalytic
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-

10
Solid radius Rsolid (km)100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カイパーベルトでは低質量ほど低密度なので、合う。	
  Ø ここで扱ったフェーズが終わると自己重力で密度は増える(Kataoka	
  et	
  al.	  2013)。	
  Ø Rsolid=100	
  kmのときのサイズ分布をプロット。	
  Ø 初めはcmでmonodisperse、fragmentがあるとµmサイズができる。	
  Karl Wahlberg Jansson and Anders Johansen: FormatiØ sub	
  virial	
  speedの衝突でmmサイズに成長する。	
  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)	
  
H9

Hβ

[OI]λ6300

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モデル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、観測バイアスがある。	
  Ø jet	  emsisionが最も強いChaI	
  217でフォローアップ。	
  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。	
  Wavelength (Angstroms)Ø absorp)on	
  doublet	
  NIとKIをBT-­‐Se0l	
  modelと比べて	
  Whelan, E.T. et al.: X-Shooter Spectroscopy of ISO-ChaI 217
Fig. 4. Spectrum as shown in Figure 1 with a zoom on key accretion and outflow tracers.
Teﬀ=2940	
  K,	
  v	
  sini=20	
  km/s,	
  log	
  g=3.5-­‐4.0ß質量と半径に一致。	
  Flux (erg/s/cm2/A)

H13

H12

H11

H8

H10

[NII]λ6583

Hα[SII]λ6731[FeII]λ7155[SII]λ6716

CaIIλ8498

CaIIλ8542

CaIIλ8662

HeI 10830

Paγ

Paβ

Ca II Ca II +H7

Ηγ
range in which the outflow is detected is shown. The
spectra[S II] ΗδH8were corrected for this velocityH9di↵erence.
All velocities givenin this paper are also corrected for the stellar rest velocity of
ISO-ChaI 217 measured with respect to the local standard of rest
(LSR). In Whelan et al. (2009b) the average velocity of YSOs in
Hα
Cha I of 2 km s 1 was adopted for[Οtherest velocity of ISOI] stellar
II]
ChaI 217. Joergens et al. (2012) measured[Stherest velocity of
1
ISO-ChaI 217 to be ⇠6.6 km s , from the Li I absorption feature
at ⇠ 6708 Å. Here we adopt the value to be ⇠ 6.4 km s 1 as measured from the same Li I feature. For the position-velocity (PV)diagrams presented here, bothemission and any
Paγ the continuumPaβ
He I
sky lines were removed. This was done using the continuum routine within the Image Reduction and Analysis Facility (IRAF)1 .
Each continuum row or background column was plotted separately, the extent of the continuum or sky line fitted and thenthe fit was subtracted. For Figures 1, 4 and 5 Gaussian smooth-

Hβ
He I

Ca II

Accre)on-­‐ejec)on	
  connec)on	
  in	
  the	
  young	
  brown	
  dwarf	
  candidate	  ISO-­‐ChaI	
  217	
  Whelan	
  E.	
  T.	
  et	
  al.	  2014	
  (続き)	
  Whelan, E.T. et al.

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では高速度トレーサー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側はイオン化率が小さいということ。	
  B-5 +/- 5-15 +/- 5

Spectro-­‐astrometric	  Analysis	Fig. 6. Spectro-astrometric analysis of the [O i] 6300
lines. O↵sets are measured using Gaussian fitting. The
teriks are the o↵sets measured after continuum subtrac
accuracy in the measurement of the o↵sets is ⇠ 20 m
analysis confims that the [O i] 6300 line is tracing thwhile no outflow component is detected in the H↵ line

-20 +/- 5-24 +/- 5

to the lack of [N ii] 6583 emission from the ISO-C
red-shifted jet.

-20 +/- 5
33 +/- 5-25 +/- 5

32 +/- 5

R

4. Discussion
4.1. Estimating Ṁout / Ṁacc

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、 .
Whelan, E.T. et al.: X-Shooter Spectroscopy of ISO-ChaI 217
Avに幅を持たせて推定。	
  no outflow component is detected. It is likely that the oØ Mdot,out/Mdot,acc=0.05-­‐0.2はCTTSやそのmagneto-­‐centrifugal	
   whichflow can be traced by the H↵ line but the emission is many tim
fainter than the accretion component. Note that no Brackett lijet	
  launching	
  modelsと合う。	
  and only two Paschen lines are detected.
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 tracISO−Cha
I 217 Case
ers. The solid line gives
the meanvalueB 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
100.0 mainly come from the error in the line fluxes andmeasurements
Ø HIも検出され、その起源を調べるためにBalmer	
  decrements
in the value of the on-source extinction. The value of Ṁ cal.
culated by Muzerolle et al. (2005) from modelling of the H↵ line
and from10.0
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のフラックスの比をTe,	
  neから計算される値と比in Muzerolle et al. (2005).
1.0
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。	
   Whelan, E.T. et al.: X-Shooter Spectroscopyof ISO-ChaI 217
0.1
Fig. 7. Mass accretion rate estimated from various accretion tracis not required.
Estimates of Ṁ are more accurate due to the
ISO-ChaI 217

ISO-ChaI 217

acc

2000K

1010 cm−3, 1000 K
accK
1010 cm−3, 10000
108 cm−3, 1000 K
108 cm−3, 10000 K

2500K
3000K

Fn / Fβ

5000K

20000K

LTE

cr

Method B
LS II Blue (10 8 L )LS II Red (10 8 L )Ṁout (10 12 M yr 1 ) BlueṀout (10 12 M yr 1 ) Red( Ṁout Blue + Ṁout Red) / Ṁacc

0.0
4610
5490
2.15
1.63
0.078
0.045
6.0 ± 0.8
12.2 ± 4.4
1.1 ± 0.3
1.4 ± 0.3
0.7 ± 0.2
1.2 ± 0.5
0.05 (+0.07)(-0.02)

ers.
The solid line gives
the mean value of log ( Ṁacc ) and the
1.0
2.5
dashed
line
is
the
±
1
uncertainty.
The errors on the individual
4700
4920
measurements
mainly
come
from
the
error in the line fluxes and
5630
5750
in
the value of the on-source
extinction. The value of Ṁacc cal2.24
2.34
culatedet al. (2005) from modelling of the H↵ line
1.71 by Muzerolle1.81
andfrom the UVES spectra
0.063
0.048 presented by Whelan et al. (2009b)
areincluded here for 0.034
comparison. No error estimates were given
0.040
7.5Muzerolle± 1.0
10.3 ± 1.4
in
et al. (2005).
H12
14.0 ± 5.0

H1317.0 ± 6.2(2009b)the method usedassumed a value for the critical den2.3 ± 0.4
5.6 ± 1.0
sity
(n
)
which
was
highlycr
2.8 ± 0.5
6.9 ± 1.2uncertain. Here an estimate of ncr
is1.4notrequired.
Estimatesof Ṁacc are more accurate due to the± 0.3
3.3 ± 0.7
useof0.9
several accretion
2.3 ±
5.3tracers.
± 2.1 Previous to this work it has only
been
derived from the0.20
H↵(+0.30)(-0.09)
line. However, this change
in the deHα0.09 (+0.14)(-0.04)
rived value of Ṁacc does not have a significant e↵ect on the final
m2/A)

Av (mag)ne Blue (cm 3 )ne Red (cm 3
T e Blue (104 K)
T e Red (104 K)
xe Bluexe RednH Blue (104 cm 3 )nH Red (104 cm 3 )opt.thincracc

use of several accretion tracers. Previous to this work it has only
been derived from the H↵ line. However, this change in the de0 Ṁacc 2does not
4 have6a significant
810
12 the final
14
rived value ofe↵ect onWhelan,etaccal.:X-Shooter
Spectroscopy
ISO-ChaInup does notofhave
valuesof ṀE.T.
. Similarily
extinction
a largeout / Ṁe↵ect on our conclusion in this case. As for all values of Av up
to 8.
theBalmer
value measured
for theResults
driving seem
sourcetoṀbeout / Ṁacc remains
Fig.
decrements.
consistentw
the limitsof predictions
leading models.
an withinaccretion
scenario,
similar tobyPar-Lup
3-4 .
H9

H11

H8

Hβ

4.2. Investigating the origin of the HI emission
H10

Several HI emission lines are detected in the spectrum of ISO5. ChaI
Summary
217 (see Figure 4). To investigate the origin of this emission the Balmer decrements were computed with respect to H
Herea follow-up
X-Shooterstudy ofnumber
the BD(ncanlineweandpresent
are plotted
as a functionof upper quantum
up )
date
ISO-ChaI
217
and
its
outflow.
Previous
to
this
work
IS
in Figure 8. While the spectral range of X-Shooter covers
most
ChaI
217Balmer
was foundto lines
have witha bipolar
outflowwhichto exhibi
of thelines any
nup > 13were foundbe too[NII]λ6583
variousasymmetries
had been estimatedto b
noisy forinclusion inand
our Ṁ
analysis.
at[SII]λ6731
the intermeout / ṀaccAdditionally,