Star Formation Newsletter #289 No. 21-25
名古屋市立大学 楠根貴成

#21
#22

Fragmentation of Kozai-Lidov DisksBAYYAN. IX. The Initial Mass Function and Planetary-mass
Object Space Density of the TW HY Association

5p
51p

#23

Star Formation In Nearby Clouds (SFiNCs): X-ray And
Infrared Source Catalogs And Membership

86 p

#24

Atomic Gas in Debris Discs

15 p

#25

Deuteration of ammonia in the starless core Ophiuchus/
H-MM1

18 p

#21 Fragmentation of Kozai-Lidov Disks
Wen Fu, Stephen H. Lubow, and Rebecca Martin
Accepted by ApJL
背景
巨大惑星の形成理論：ディスクの重力不安定性(GI)disk GIを用いたこれまでの理論研究の多くは、単独星での星周円盤を仮定
観測結果から、40-50%が連星系の中にあると推察されている連星系でdisk GIの研究： ディスク面と連星軌道面が一緒連星の軌道面に対して傾いている（>45 ）の若い連星系が報告されている
臨界Kozai-Lidov(KL)角度（40 ）を超えるとKL振動が起きる

The Astrophysical Journal Letters, 835:L29 (5pp), 2017 February 1

KL振動： ディスクの離心率と傾斜角が交換される

step further and included the effects of disk self-gravity. We
found that the disk KL mechanism may be suppressed by disk
However, the suppression “window” in terms of
内容 self-gravity.
the disk mass and the disk inclination is quite narrow. The disk
KL mechanism may still operate even when the disk in nearly
連星軌道に対して傾いたディスクのSPHシミュレー
gravitational unstable, if the inclination is large.
For disk tilts greater than about 60o, the KL effect attempts
ションを行った
to induce a large disk eccentricity that cannot be fully
accommodated by a smooth continuous disk. Instead, the disk
undergoes strong shocks during its KL oscillations. Suchshocks could have4年）するまで続けた
important consequences for disk連星が約10周（約10fragmentation.
In this Letter, we report on fragmentation triggered by the
KL mechanism in the outer disk regions. In the current
analysis, we assume that the disk is locally isothermal, as is
more likely the case in outer disk regions, and ignore thecomplexities of non-isothermal effects. Our emphasis is on theFu, Lubow, & Martin

Table 1SPH Simulation Parameters for an Equal Mass Binary Star Systemwith Total MassM and Separation abBinary and Disk Parameters
Mass of each binary component
Binary orbital eccentricity
Initial number of particles
Initial disk mass
Initial disk outer radius
Initial disk inner radius
Mass accretion radius
Disk viscosity parameter
Disk aspect ratio
Initial disk surface density, 4 r r H
Initial disk inclination

SymbolValue

Mp M  Mc M
ebN
Mdisk M
rout abr in abracc abαH r (r  r in )
γ
i0

0.5
0
106
0.035
0.25
0.025
0.025
0.01
0.1
1.5
70°

#21 Fragmentation of Kozai-Lidov Disksの続き

The Astrophysical Journal Letters, 835:L29 (5pp), 2017 February 1

Fu, Lubow, & Martin

・4.6周期：離心率↑、アーム構造
・6.28周期：ﬁrst clump
・6.3周期：second clump
クランプ質量：3 x 10-5 M（M=1Msunの場合、10Mj）

離心率が約0.4になるとFragmentation
Figure 1. Evolution of the self-gravitating disk with initial mass Md  0.035 M , disk aspect ratio H R  0.1 (at the inner disk radius), viscosity parameter B  0.01,
and an initial inclination of 70◦. Each panel corresponds to a time (in units of the binary orbital period P b ) and each panel has the same length scale. The disk isprecessing and also changing tilt angle in time, but we show the face-on view of the disk. The binary separation is ab . The central binary component is at the center of
each panel (denoted by a black circle), while the perturber star is out of the view on this scale. Newly formed clumps/fragments are denoted by black dots. The color
coding is for the logarithm base 10 of the column density (i.e., the density integrated along the line of sight) in units of M ab2 . The disk initially is modeled with 106
SPH particles.
The Astrophysical Journal Letters, 835:L29 (5pp), 2017 February 1

この段階：KLサイクルの初期なのでディスクの傾きは変わらない

Fu, Lubow, & Martin

38 Mab3) and an accretion radius of 0.001 ab is imposed
Figure 1 where disk self-gravity is included. The right column,
around a newly formed sink particle. If ab  100 au andin contrast, shows a simulation without disk self-gravity (andthus no disk fragmentation) while all the other simulation
M  1 M:, then the critical density is 3×10−10 g cm3 andthe accretion radius is 0.1 au.
parameters are the same. The left column demonstrates that
Figure 1 shows the evolution of the disk column density up
disk fragmentation occurs when the disk eccentricity has grown
to a time of about 6.3 binary orbits. The disk orientation
up to almost 0.4. At this stage, the disk is still fairly early in its
changes due to both orbital precession and the KL oscillations
KL cycle, so the disk inclination has only declined slightly. If
but we always show the face-on view of the disk. Because of
there is no fragmentation (see right column), then we see an
the spatial scale of the ﬁgure, the perturber star is out of view,interrupted disk KL oscillation during which disk eccentricity
while the primary component is at the center of each panel. The
can grow up to about 0.7 while the disk inclination falls to
top left panel shows that the disk is initially circular. Afterabout 30◦. Note that in both simulations the disk remains quite
ﬂat throughout the simulation (see the similarity of disk
about three binary orbits, the disk remains circular but the
inclination at the three different radii plotted). The KL cycle inperturber star has driven spiral waves in the disk. In the top
the run that includes self-gravity starts at an earlier time thanright panel, which is at t x 4.6 Pb , the disk starts to show someeccentricity and a relatively strong one-arm spiral shock. At
the one without self-gravity. This is because the initial drivingt x 6 Pb (bottom left panel), the disk has clearly become moreof disk KL oscillation is very sensitive to initial model setup.
eccentric and the spiral shock has evolved into an arc-like
Even slightly different initial conditions could lead toshocked region with material being even more concentrated
signiﬁcantly different lengths of the pre-KL stage (as discussed
along that arc. Shortly after, at a time of t x 6.28 Pb (bottom
in Fu et al. 2015b). Also note that the disk eccentricity in the
middle panel), the ﬁrst clump forms (denoted by the black dot).
beginning is slightly higher than zero even though we set the
A second clump forms shortly after this, as shown in the last
disk up to be circular. This is due to the fact that we assume thepanel at time t x 6.3 Pb . Both clumps have a mass around
SPH particles are on Keplerian orbits when computing their5
3 q 10 M . This makes them about 10 M objects if we takeorbital eccentricities, whereas they are actually on slightly subFigure 2. Evolution of the eccentricity (upper panels) and inclination (lower panels) of the disk. We have averaged quantities over one binary orbit to smooth out
form at a radius ofthesmall-scalebinary oscillations.
mass to be
Keplerian orbits because of the disk pressure (FuM  1 M:. They
The left panel shows
the simulation that include disk self-gravity (the run presented in Figure 1), whereas the right panel shows the simulation
rx
0.2adisk
the Different
local disk
value three
is now
about
et al.star2015a,
2015b).
b , whereab (blue), r  0.2 ab (red), and r  0.3 ab (green). Lines inwithout
self-gravity.
linesQrepresent
different
disk1.5.
radii By
from the centralof r  0.1t  6.4
comparison,
disk
value
onPtheotherof the disk
etcolumnal. (2015b)
wepresented
a cyclesimulationthatthedisksidefragmentation
occurs. Lines inIntheFu
right
exhibit thecompletedisk KL
when there
is noisselfthe left columnthe
endlocalat a time
of Q
about
b whengravity
andthereforenofragmentation.
(also x0.2ab ) is almost 10 times higher (Q x 15).
same as the one in Figure 1, except that the disk initial tilt was
In Figure 2 we show the time evolution of the eccentricity
50◦, as opposed to 70◦ here. A disk with an initial tilt of 50◦

KL振動がディスクのFragmentationをトリガー連星系での巨大惑星の効率的な形成メカニズム

#22 BANYAN. IX. The Initial Mass Function and Planetary-mass
Object Space Density of the TW HYA Association
Jonathan Gange, Jacqueline K. Faherty, Eric I. Mamajek, Lison Malo, Rane Doyon et al.
Accepted by ApJS
背景
若い移動アソシエーションの研究が近年注目されている
とりわけ、太陽系近傍の明るい極めて低質量の準恒星の含んでいる可能性に注目
100pc以内で最も若いアソシエーションがTW Hya (TWA) (~10 Myr)
IMFの見積もりはLooper (2011)のみ
近年、L7型星（5-7Mj）が発見
TWAのIMFの改訂、惑星質量メンバーの空間密度の見積もりが求められている
内容○TWAの新しいメンバーの同定
・BANYAN All-Sky Survey (BASS)：
太陽系近傍の若い移動集団（~M5-L5）の候補天体を同定するサーベイ
2MASSとALLWISEのcross matchに基づき、固有運動を測る
・BASS-Ultracool Survey：L5以上のメンバーを同定
・SUPERBLINK proper motion survey：○観測
FIRE：近赤スペクトルを取得

Flamingos-2 at GEMINI-S：低分散近赤スペクトル

GMOS：可視スペクトルを取得

ESpaDonS at CFHT：可視高分散スペクトル

SpeX at IRTE：低分散近赤スペクトル#22 BANYAN. IX. The Initial Mass Function and Planetary-mass
Object Space Density of the TW HYA Association の続き
The Astrophysical Journal Supplement Series, 228:18 (51pp), 2017 February

Gagné et al.

メンバーの可能性が高い星：・対数正規分布（中心：0.21Msun、
幅：0.8dex in 12Mj-2Msun）に従う
・一方、0.1-2MsunではIMFの傾きはSalpeter冪則α=2.2で良く表現される

Figure 12. Fitted log-normal cumulative IMFs (thick red line) to the observed unbinned TWA cumulative IMFs (blue dashed–dotted line). The thin blue dotted lines
represent the ±1σ range due to small number statistics based on a Poisson distribution as well as errors on the TWA effective volume estimate. The yellow region
represents the mass
range used
to performthe ﬁt,228:18
and the(51pp),
red shaded
indicate 1σ–3σ random draws from the MCMC solutions. See Section 8Gagné
for moredetails.
The AstrophysicalJournalSupplement
Series,2017regions
February
et al.

due to TWA being denser than the ﬁeld: this estimate corresponds
3.1
to one expected ∼5–7 MJup TWA object for every 1.90.6
currently known main-sequence ( 75 MJup) member of TWA.
This estimated fraction of planetary-mass to stellar TWA
members is larger than recent estimates based on the lowest-mass
candidate members of the Tucana–Horologium Association(THA). Gagné et al. (2015b) estimated this fraction at one
6.6
main-sequence member of
12.5–14.0 MJup object per 17.55.0
THA. Using the Bayesian formalism presented here, we revise this
4.8
q 103 objects pc−3 mainto one low-mass object per 16.63.2
sequence THA member. This estimate remains high in comparison
to predictions based on a log-normal IMF of TWA. Such an
estimation for TWA, based on the primary + companions IMF ofhigh-likelihood and bona ﬁde members (Section 8), would predict
0.4_ 5–7 MJup objects in the whole TWA
a total of only 0.50.2
association.

It is possible to make a prediction for the ﬁeld space density of∼5–7 MJup objects by assuming that the low-mass end of the ﬁeld
IMF is identical to that of TWA. This was done by assuming thatthe ﬁeld is well-represented by a ﬁducial log-normal IMF anchoredon the stellar space density of Chabrier (2005), and anchoring it inturn on the stellar population of TWA. Since TWA has a notablelack of massive, early-type stars (no A3–G9 high-likelihood
candidates or members are currently known), the ﬁeld IMF wasanchored on K0–K9 members of TWA, corresponding to a mass
range of ∼0.7–1.5 Me. This calculation based on TWA yields a29
q 103 objects pc−3 forpredicted space density of 2615∼5–7 MJup ﬁeld objects, most of which would be faint, Y-type
dwarfs (300 K; Luhman 2014). This estimate should be seen as
an upper limit because it is likely that there are still K dwarfs thatare missing in the current TWA census. We avoided anchoring thisestimation on the two A-type members of TWA (A0 and A2)
because these spectral types span a small range of masses

特徴的な幅：他のアソシエーションよりも大きい
TWAの低質量星の調査が不十分であることに起因しているかも

孤立した惑星質量メンバーが多いことを確認

Figure 12. Fitted log-normal cumulative IMFs (thick red line) to the observed unbinned TWA cumulative IMFs (blue dashed–dotted line). The thin blue dotted lines
represent the ±1σ range due to small number statistics based on a Poisson distribution as well as errors on the TWA effective volume estimate. The yellow region
40 random draws from the MCMC solutions. See Section 8 for more details.
represents the mass range used to perform the ﬁt, and the red shaded regions indicate 1σ–3σ
Figure 13. Fitted Salpeter cumulative IMFs to observed unbinned TWA cumulative IMFs. Colors and curves are similar to those of Figure 12. The cumulative integralof the IMFs used for the Salpeter ﬁts have a lower bound of 0.4 Me. See Section 8 for more details.

due to TWA being denser than the ﬁeld: this estimate corresponds
3.1
to one expected ∼5–7 MJup TWA object for every 1.90.6
currently known main-sequence ( 75 MJup) member of TWA.
This estimated fraction of planetary-mass to stellar TWA
members is larger than recent estimates based on the lowest-mass
candidate members of the Tucana–Horologium Association(THA). Gagné et al. (2015b) estimated this fraction at one
6.6
main-sequence member of
12.5–14.0 MJup object per 17.55.0
THA. Using the Bayesian formalism presented here, we revise this
4.8
q 103 objects pc−3 mainto one low-mass object per 16.63.2

It is possible to make a prediction for the ﬁeld space density of∼5–7 M objects by assuming that the low-mass end of the ﬁeld
視線速度の測定によって
IMF is identical to that of TWA. This was done by assuming thatthe ﬁeld is well-represented by a ﬁducial log-normal IMF anchoredon the stellar space density of Chabrier (2005), and anchoring it in・2MASS J11472421がメンバーであることを裏付けた
turn on the stellar population of TWA. Since TWA has a notablelack of massive, early-type stars (no A3–G9 high-likelihood
・TWA 28 (M8.5γ)・TWA 29(M9.5γ)・TWA
33(M4.5e)をメンバーとして同定
candidates
or members are currently known), theﬁeld IMF
wasanchored on K0–K9 members of TWA, corresponding to a mass
range of ∼0.7–1.5 M . This calculation based on TWA yields a
TWAとは関係のない若いL7型星
26q 10
predicted2MASSspace density ofJ09553336-0208403を発見した
objects pc
for
Jup

e

29

3

−3

#23 Star Formation In Nearby Clouds (SFiNCs): X-ray And Infrared
Source Catalogs And Membership
Konstantin V. Getman, Patrick S. Broos, Michael A. Kuen, et al.
Accepted by ApJS

モチベーション
観測による星形成の特徴付けが乏しく、星団形成過程の詳細やその初期進化の理解が不十分
主な原因：様々な星形成環境のYSOに関する信頼できる調査研究がない
近年、X線と近赤外線のデータを使ってのYSOセレクションが進んでいる（MYStIX）
SFiNCsプロジェクト：MYStIX（Massive Young Star-Forming Complex Study in Infrared and Xray）の拡大
MYStIX：20のmassive SFRs

SFiNCs：近傍の22のより小さなSFRs

近傍の22の星形成領域でのYSOや星団形成の詳細な研究成果の提供を目的とする
内容
X線：Chandra-ACIS
赤外：2MASS、Spitzer-IRAC のアーカイブデータを使用してYSOのクラス分け、空間分布図を作成
これらデータから導かれる科学的問題は、別の論文で議論する予定

#23 Star Formation In Nearby Clouds (SFiNCs): X-ray And Infrared
Source Catalogs And Membership の続き
YSOの検出（Gutermuth et al. 2009、Getman et al. 2012を参照）
1. IRACの4バンドを使って、星形成銀河を除く（Gutermuth et al. 2009）
2. 2MASSとIRACを使ってSEDを書く
3. 暗いMIRカウンターパートを持つX線源を除く（X線AGNと考えられるため；Getman et al. 2012）
60

Table 1. SFiNCs Star Forming Regions
Name

R.A.(J2000)

Decl.(J2000)

l

b

D

:m : s:0 :00

deg

deg

pc

(1)

(2)

(3)

(4)

(5)

(6)

NGC 7822 (Be 59)

00:02:15.27+67:25:59.7

118.2262

5.0120IRAS 00013+6817 (SFO 2)

00:04:03.85+68:33:33.7

118.6050

6.0865hDist. Ref.

OB

Spitzer Ref.

Chandra Ref.

(7)

(8)

(9)

(10)

900

Ma08;Pa08

O5V

···

Ma08;Pa08

···

···

900

Gu09

···

Wi10

NGC 1333

03:29:04.76+31:20:10.7

158.3063-20.4982

235

Hi08

B5V

Gu09

IC 34803:44:30.30+32:07:44.8

160.5025-17.8368

300

He08

B5V

Lu16

St12

LkH↵ 101

04:30:09.86+35:16:09.2

165.3491-9.0203

510

Wo10

B4

Gu09

Wo10

NGC 2068-2071

05:46:43.54+00:07:34.2

205.2813-14.2981

414

Me07

B2/3

Me12

Gr04;Sk09

ONC Flank S

05:35:05.79-05:39:07.9

209.2379-19.5421

414

Me07

B9V

Me12

Ra04

ONC Flank N

05:35:14.38-04:49:49.5

208.4786-19.1380

414

Me07

B1V

Me12

Ra04

OMC 2-3

05:35:23.44-05:07:03.1

208.7676-19.2349

414

Me07

B0.5V

Me12

Ts02

Mon R2

06:07:46.15-06:22:53.6

213.7009-12.6029

830

Ra68

B0

Gu09

Na03

GGD 12-15

06:10:49.51-06:11:45.7

213.8772-11.8404

830

Ra68

B8/A0

Gu09

···

RCW 120

17:12:23.25-38:28:45.4

348.2381

0.4639

1350

Za07

O9:

···

18:29:56.30+01:13:11.5

31.5647

5.3394

415

Dz10

···

···

Serpens Main

Gu09

Gi07;Wi07

Serpens South

18:30:02.56-02:03:42.2

28.6501

3.8133

415

Dz10

···

Po13

···

IRAS 20050+2720

20:07:06.56+27:30:02.3

65.7962-2.6014

700

Wi89

B3

Gu09

Gu12

Sh 2-106

20:27:24.36+37:22:54.9

76.3802-0.6137

1400

No05;Sc07

B0V

Gu09

Gi04

IC 5146

21:53:29.82+47:15:22.4

94.3896-5.5133

800

Ga14

B2V

Gu09

···

NGC 7160

21:53:49.98+62:35:10.2

104.0159

6.4318

870

Co02

B0.5V

Si06

22:38:48.05+75:11:50.4

114.6324

14.5068

300

Ku08

···

···

LDN 1251B

Ev03

Si09

Cep OB3b

22:55:39.32+62:38:11.2

110.0643

2.6936700

Mo09;Dz11

O7V

Al12

Ge06;Al12

Cep A

22:56:14.13+62:02:17.8

109.8676

2.1246700

Mo09;Dz11

O9V

Gu09

Pr09

Cep C

23:05:49.42+62:30:27.2

111.0778

2.0919700

Mo09;Dz11

B9:

Gu09

···

Note—Column 1: SFR name. Columns 2-5: Coordinates: right ascension, declination, Galactic longitude, and Galactic latitude. Column 6:
Distance from the Sun, in parsecs. Column 7: Literature reference to the distance value. Column 8: The earliest OB star within the Chandra
field from the catalogs of Ski↵ (2009) and SIMBAD. Columns 9-10: Literature references to the previous Spitzer and Chandra studies of theregion. Reference code in Columns 7, 9, and 10: Al12 (Allen et al. 2012), Co02 (Contreras et al. 2002), Dz10 (Dzib et al. 2010), Dz11 (Dzib etal. 2011), Ev03 (Evans et al. 2003), Ga14 (Garcı́a-Rojas et al. 2014), Ge06 (Getman et al. 2006), Gi04 (Giardino et al. 2004), Gi07 (Giardino etal. 2007), Gr04 (Grosso et al. 2004), Gu09 (Gutermuth et al. 2009), Gu12 (Günther et al. 2012), He08 (Herbst 2008), Hi08 (Hirota et al. 2008),
Ku08 (Kun et al. 2008), Lu16 (Luhman et al. 2016), Ma08 (Majaess et al. 2008), Me07 (Menten et al. 2007), Me12 (Megeath et al. 2012), Mo09
(Moscadelli et al. 2009), Na03 (Nakajima et al. 2003), No05 (Noel et al. 2005), Pa08 (Pandey et al. 2008), Po13 (Povich et al. 2013), Pr09
(Pravdo et al. 2009), Ra68 (Racine 1968), Ra04 (Ramı́rez et al. 2004), Sc07 (Schneider et al. 2007), Si06 (Sicilia-Aguilar et al. 2006), Si09
(Simon 2009), Sk09 (Skinner et al. 2009), St12 (Stelzer et al. 2012), Ts02 (Tsujimoto et al. 2002), Wi89 (Wilking et al. 1989), Wi07 (Winston etal. 2007), Wi10 (Winston et al. 2010), Wo10 (Wolk et al. 2010), Za07 (Zavagno et al. 2007).

上記クラス分けによって、8492のSFiNCs probable cluster member（SPCM）
を同定
過去のSpitzerやChandraのYSOカタログと比べて：
それぞれのSFCsで6-200%ほどYSOメンバーが増えた
SPCMのpropertyを掲載
また、各領域ごとのSPCM分布図を作成・掲載

#24 Atomic Gas in Debris Discs
Antonio S. Hales, M. J. Barlow, I. A. Crawford and Casassus
Accepted by MNRAS

背景
Debrisディスク： 惑星形成過程の最終段階
いくつかのデブリディスクは、わずかに星周ガスの存在を示す
初期ディスクの残りなのか、seconday起源なのかは不明※吸収線を使ってガスを検出するのは難しい（edge-onでなければならないから）ガスを持つβPicのような系を捜すために、16のデブリディスクの可視スペクトルを調べた
Atomic Gas in Debris Di
Star

Other
Name

Previous
Spectral TypeV

(B-V) Distance(pc)

LIR /L?

HD 99211
Crt
HD 110058
HD 117360 S Cha
HD 123247
HD 166841
HD 176638 ⇣ CrA
HD 181327

F5V
B8V
B8V
A0V
A1V/A2V
A1V
A9V
A0V
F5V/F6V
B9.5V/B9V
B8V/B9V
B9V/A0V
F5V/F6V

7.46
6.89
5.46
5.87
3.84
4.72
4.08
7.98
6.48
6.43
6.32
4.75
7.04

0.460
-0.028-0.094-0.016
0.063
0.042
0.216
0.148
0.480
0.000
-0.030-0.027
0.480

57.3
362.3
96.9
80.8
31.5
79.4
25.7
107
35.1
101.1
213.7
56.3
50.6

2.6⇥101.4⇥101.8⇥106.6⇥107.5⇥101.3⇥101.1⇥102.0⇥101.5⇥102.4⇥101.1⇥105.1⇥102.4⇥103.95 -0.083
7.17 0.480

52.1
53.5

7.9⇥10 6
1.8⇥10 3

161.4
3.5

0.494

9.2

1.0⇥10 5

2.4

103

4

5.32 0.044
4.06 0.100
5.12 -0.104

75.8
29.6
90.4

-

-

-

-

ターゲット
HD 53842
61950
Mannings & Barlow (1998)のカタログから、HD
HD 75416 ⌘ Cha
HD 80950
88955
q Vel
ダストのFractional luminosityが低い天体を HD
HD 91375
デブリディスク天体としてターゲットにした

結果

HD 181869
HD 191089

↵ Sgr

B8V
F5V

HD 203608

Pav

F7V/F6V

HR 5558
HR 5670
HR 6045

Cir

A0V
A3V
B8V

✓ Nor

4.22

L?
(L )
4
4
5
5
6
5
5
3
4
4
4
5
3

5.0
1121.8
206.8
48.8
41.7
123.4
18.8
33.0
5.4
52.3
450.8
64.4
4.1

Age [Myr] References
for Age
8
80
300
265
600
17
12+8
4
1103
12+8
4

0
1
1
1
2
5

3
1
3

輝線の証拠なし
Table 1. Debris disc host stars observed along with spectral standards used for spectral classification. Spectral types are fromSpectral Catalogue (Houk & Smith-Moore 1988) and distances are derived from parallax measures listed in the Hipparcos cataより進化していて、不活発な進化ステージにあることを裏付ける
Leeuwen 2007). B-V colours are taken from the Simbad database. The LIR /L? values were derived using the 25 µm and 60
excess fluxes. L? is the luminosity of the Castelli & Kurucz (2004) stellar model, normalised to the 12 µm flux. LIR is theluminosity of the black-body that fits the infrared excesses.

quoted for V

and v sin i correspond to the dispersion obtained when deriving the quantities using the di↵erent lines.

#24 Atomic Gas in Debris Discs の続き４つ（HD61950, HD75416, HD110058, 166841）は
光球のCaII輝線とNaID輝線に近いところで、狭い吸収を示す
さらに、同様の吸収がスペクトル標準星のスペクトル中でも同定された
そのうち２つは17Myrよりも若い
Welsh & Montgomery (2013)の成果と一致『星周ガスの活発性を示す星は不活発な系よりも若い』
Atomic Gas in Debris Discs

13

Figure 2. Heliocentric velocity plots of the narrow Ca ii H (dotted line) and K (solid line) absorption components present in the
spectra of HD 61950 and of HD 75416. The vertical lines denote
the radial velocity of the star. The H line spectra have been shiftedvertically for plotting purposes.

Figure 3. Heliocentric velocity plots of narrow absorption features seen towards HD 110058 in both the Ca ii H & K lines (toppanel) and in the Na i D lines (bottom panel). The solid lines correspond to the Ca ii K and Na i D2 lines. The Na i region of thespectra was divided by the spectrum of the B8 spectral standard
HR 6045 in order to remove telluric absorption features.

Figure 9. MIKE Ca ii spectra towards HD 110058 and reference targets for interstellar absorption characterization. The parallactic
distance to the reference stars and their angular separationWhen
from HDmultiple
110058 are epoch
annotatedin the figure.
observations
are available, one canreturns the equivalent width and centre of these unsaturated
also study variability of the absorption features, as circumcomponents. Table 5 presents the derived equivalent widths
stellar lines often vary with time (Beust et al. 1998; RedA
strongnarrowabsorption
featureiseasily
seen
5.3.5
HR
5558
and radial velocities of the narrow Ca ii and Na i absorption
星周ガスの最も説得力のある証拠：
HD110058のスペクトルに見える吸収の２つの速度成分の一つ
et al.
right at the centre of the photospheric Ca iifieldand Nai 2007; Welsh & Montgomery 2013). However notfeatures present in the spectra of our targets.
HR 5558 is an A0V spectral standard located 76 pc away
lines. Both the calcium and sodium absorption
all features
circumstellarlines are variable, as for instance in Pic
from the Sun. We serendipitously detected a narrow ab1
are foundtobelocated
at
V
=0±1
km
s
,
coinciding
・星の速度と一致、the strongestsorption
component
is stable
andNacorresponds
thefeature noticeable
in thei D lines (EW ofto
20 m
extremely well with the radial velocity of the star. The
and 33 m in(asthe opposite
D1 and D2 tolinestherespectively),
and
less
main
disc
absorption
variablecompoderived equivalent widths of the di↵erent component
5.1 Cloud modelingevident in the Ca ii K line. The feature was not detected insuggest ・CaII/NaIの柱密度比that the lines are not saturated and cloudmodnents~1、
which areinterpreted as evaporating comets; Beust et
the Ca ii H line.
eling indicates that Ca ii/Na i column density ratio⇠ 0.3.
al.is1998).
The Voigt Absorption Profile/Interstellar Dabbler software
・その方向に星間雲やシェルがない
There
is
no
correspondenceoftheradialvelocityoftheWe measured the equivalent widths and radial velocities
for modeling interstellar absorption lines (VAPID; Howarth
Although the velocity projections of known
feature with the radial velocity of the star, sugthe narrowabsorption
absorption
featuresonthenormalized
spectraet al. 2002), was used to estimate the column density (N ),clouds/shells cannot be responsible for theof observed
gesting that the absorption is of interstellar origin. This is・ディスクをedge-onで見ている
absorption feature, we investigated the presenceby
of using
similar the supported
Emission
Starlink
radial velocity (v) and the velocity dispersion (b) of the maby Linethe lowFitting
Ca ii/Na (ELF)
i ratio derived.
Givenroutine,
the diabsorption components in the spectra of stars in similaragnostics above,
plustothethefact thatHR 5558 isfeatures
not known and
to
which
fits
Gaussian
profilesabsorption
terial causing the observed absorption in the Ca ii K and
directions. The LSR velocity of the component seen toward
be a debris disc host star, we conclude that the absorption

吸収はHD110058のディスク内の原子ガスで起きているfeature seen toward this star is interstellar in origin.

HD 166841 (l=326o , b=-22o , d=214 pc) is -1.8 km s 1 .
Hunter et al. (2006) reports +2.4±0.1 km s 1 Ca ii Koo

MNRAS 000, 1–15 (2016)

#25 Deuteration of Ammonia in the starless core Ophiuchus/H-MM1
J. Harju, F. Daniel, O. Sipila, P. Caselli, J.E. Pineda, R.K. Friesen, A. Punanova, et al.
Accepted by A&A

背景
アンモニアと重水素化アンモニア： 分子雲コアの物理状態を探る
重水素の分別の時間依存・核スピンの相対存在量： 分子雲コアの進化ステージを決定する手法
目的
不活発なdense coreでのNH2D、NHD2、ND3の存在量とスピン比を決定することで、重水素化された分子のスピン状態化学を検証する

観測
対象：Ophiuchusの星なしコア
望遠鏡：APEX, GBT, IRAM 30-m 望遠鏡
スペクトル線：NH3, NH2D, NHD2, ND3, N2D+
放射輸送計算と組合わさったガス­粒子化学モデルを用いて、観測を解釈
化学モデル：軽水素分子の核スピン状態と、アンモニアとその重水素化物を区別する

#25 Deuteration of Ammonia in the starless core Ophiuchus/H-MM1 の続き
結果
高い重水素分別比が検出された
NH2D/NH3 ~0.4、NHD2/NH2D ~0.2、ND3/NHD2 ~0.06
NH2DとNHD2のオルソ・パラ比： 統計的な核スピン比に近い
化学モデルは観測された存在量をほぼ再現できる
一方、NH2Dのオルソ・パラ比をかなり低く予測し、NHD2のオルソ・パラ比をかなり高く予測する
N2H+とNHD2の長い寿命（in高密度ガス）： 粒子表面のN2の化学的不活発が原因
結論
化学モデルと観測の不一致： アンモニアの気相重水素反応における完全なスクランブルを
モデルは仮定していることが原因の可能性が高い
全ての原子核は反応性衝突で混ぜられることを意味する
今回の結果からわかること
アンモニアの重水素化： 物理条件や時間で変化する
アンモニア同位体置換体の核スピン比： 分子雲の進化ステージは探れない