星形成ゼミ 2014年6月13日
論文速報 SFN #257 (46–50本目)
東京大学 博士2年 小屋松 進"ESO-Ha 574 and Par-Lup3-4 jets: Exploring the spectral, kinematical and physicalproperties” E.T. Whelan et al.
"Multi-epoch Sub-arcsecond [Fe II] Spectroimaging of the DG Tau Outflows with NIFS. I.
First data epoch” Marc C White et al.
"Multi-epoch Sub-arcsecond [Fe II] Spectroimaging of the DG Tau Outflows with NIFS. II.
On the Nature of the Bipolar Outflow Asymmetry” Marc C White et al.
"A Parametric Modeling Approach to Measuring the Gas Masses of Circumstellar Disks”
Jonathan P. Williams and William M. J. Best"A spider-like outflow in Barnard 5 - IRS 1: The transition from a collimated jet to a wideangle outflow?” Luis A. Zapata et al.

ESO-Hα 574 and Par-Lup3-4 jets: Exploring the spectral,
kinematical and physical properties
E. T. Whelan, R. Bonito, S. Antoniucci, J. M. Alcalá, T. Giannini, B. Nisini, F. Bacciotti6, L. Podio, B.
Stelzer, and F. Comerón

ジェットを伴う2つの系のVLTのX-ShooterおよびUVESによる観測
2011年のレター(Bacciotti et al. 2011)の詳細
ooter Spectra of the ESO-HA 574 Jet【背景】
ジェットの形成シナリオ
B
Magneto-centrifugal jet launching model が有力だが観測的な証拠がほしい
A
多波長分光観測により正確な質量の放出率と降着率の比を求めてモデルを検証
A1【天体】
・ESO-Hα 574: K8, 古典的おうし座T型星(CTTS)
・Par-Lup 3-4: 超低質量(0.13E Msun)(VLMS), M5
X-Shooter Spectra of the Par-Lup3-4 Jet
ともにedge-onに近い降着円盤をもつ天体

Shooter observations of the ESO-H↵ 574 and Par-Lup 3-4 Jets

B(←) ESO-Hα 574のPV図

A

A&A 5

-

R

-

A1
B

E

B

(→) Par-Lup 3-4のPV図

Fig. 7. Position velocity diagram of the Par-Lup 3-4 jet in [S II] 67
with the velocity centroid measured at various positions along theTable A.1. continued.
Identification
Ion

Knot
Flux (⇥10 17 erg/s/cm2 )obs (Å)A
10 286.6 20 ± 1
2
10 320.5[S II]
P3/2 –2 D5/2
24 571.5
10 320.4 64 ± 1
A 57410and
320.4
203-4
± 1Jets
E. T. Whelan et al.: X-Shooter observations of the ESO-H↵
Par-Lup2
2
10 336.4[S II]
P1/2 – D3/2
24 524.8
10 336.3 50 ± 1
A
10 336.3 14 ± 1
Appendix A: Tables of identified lines and line fluxes
2
10 370.5[S II]
P1/2 –2 D5/2
24 524.8
10 370.4 22 ± 1
Table A.1. ESO-H↵ 574 emission line fluxes.
A
10 370.4 7 ± 1
2
10 397.7a[N I]
P3/2 –2 D5/2
28 839.3
10 398.0 4 ±1
2
2
10 398.2a[N I]
28 838.9
Identification 2 P1/2 –2 D5/2
10 407.2a[N I]
P3/2 – D3/2
28 839.3
10 407.2 3 ± 1
Eu (cm 1 ) Knot obs (Å) Flux (⇥10 17 erg/s/cm2 )air (Å) a Ion
2Type 210 407.6[N I]
P1/2 – D3/2
28 838.9
a3721.9
H
I
109
119.2
± 0.5
314–23
10 830.3
He I
P–S
169 087.0
103721.8
830.1 3.5
132.7
± 0.8
1 1 3 13721.6a[S III]S0 – P1
27 161.00 A
10
830.3
56.2
± 0.3
2
4
3726.0
[O
II]
D
–
S
26
830.6
3726.0
76.8
±
3/2
3/2
10 938.0HI
0–3
106 632.2
10 937.8 14.5 ± 0.5
0.8
11 618.9[Ti II]
a2 G7/2 –a4 F9/2
8997.7 A
113726.0
618.9 71.7
6 ± 1± 0.5
46B3726.018.3±
0.5
12 485.4[Fe II]a D5/2 –a D7/2
8391.912 484.9 2.3 ± 0.4
612 566.8[Fe II]a2 4 D7/2 –aD9/2
7955.3 E
123726.6
566.6 9.2
82.7±±0.50.73728.8[O II]
D5/2 –4 S3/2
26 810.6 A
123728.8
566.4 38.0
16.7 ±± 0.5
0.2
12 610.9[Cr I]
a5 D2 –a7 S3
7927.5 A
123728.8
611.0 46.7
3 ± 1± 0.5
B3728.822.5±
0.5
12 787.8[Fe II]a4 D3/2 –a6 D3/2
8680.45
12 787.5 11 ± 2
12 818.1
HI
5–3
105 291.7 E
123729.0
817.7 13.3
13 ± ±1 0.5
3734.4
H
I
13–2
109
029.8
4
6
12 942.7[Fe II]a D5/2 –a D5/2
8391.9123734.3
942.5 1.3
13.8±±0.50.83750.2H[FeI II]
108680.4
8917.1
4
6
12
977.7
a12–2
123750.3
977.8 1.4
3.7 ±
± 0.5
0.8
a
3D3/2 –a
3 D1/23769.9
[Co
II]
a
27
469.1
1.7
±
0.5
4 H5 –a F
63
13 205.5[Fe II]a3 D7/2 –a7955.30
13 205.3 26.9 ± 0.8
a
3 D7/23770.2I II]
F4 –a–aG63D
418391.9
624.2
13277.8 Ti[Feas4 D
13 277.8 9.5 ± 0.9
5/2
3/2
3770.6
H
I
11–2
108
772.3
4
4
15 334.7[Fe II]a D5/2 –a F9/2
8391.9153771.3
334.4 1.713.9±±0.5
0.63797.9HI
10–2
108 582.0
3797.8 2.3 ± 0.5
15 994.7[Fe II]a4 D3/2 –a4 F7/2
8680.515 994.7 14.4 ± 0.5
中略3835.3HI
9–2
108 324.7
3835.3 4.9 ± 0.5
16 435.5[Fe II]a1 4 D7/23 –a4 F9/2
7955.316 435.3 72.5 ± 0.7
3869.1[Ne III]D – P2 4
25 838.7
3968.9 11.5 ± 0.5
16 637.7[Fe II]a4 D2 1/2 –aF5/2
8846.816 637.4 7.2 ± 0.6
3889.0HI
8–2
107 965.1
3888.9 1.2 ± 0.5
16 768.8[Fe II]a2 4 D5/22–a4 F7/2
8391.916 768.1 16.7 ± 0.6
3933.7
Ca II
P3/2 – S1/2
25 414.4
3933.7 16.6 ± 0.3
17 475.0H
1-0S(7)
1.74746 9.3 ± 0.9
133967.8[Ne2 III]D
25 838.7
3967.4 4.7 ± 0.3
4 2 – P1 4
17 971.0
[Fe
II]
a
D
–a
F
8680.517970.7 4 ± 1
3/2
3/2
a
4
3968.2[Fe
G5/2–a–a46FD3/2
267955.3
055.4
18
093.9
[FeII]II]aa42D
183968.5
093.5 12.5
16 ± ±2 0.3
7/2
a
4 7/2
3968.2
[Ti
II]
c
D
–a
F
25
192.8
5/2
3/2
19 570.2
H2
1-0S(3)19 570.8 69.7 ± 0.2a22
3968.4
P1/2 – S1/2
25 191.5
21212.5 CaH2II
1-0S(1)21 212.5 39.2 ± 0.2
3970.1
7–2
107 440.5
0.3
22226.8 HHI 2
1-0S(0)223970.1
227.0 10.0
8.4 ±±0.3
A0.3
22 471.0
H2
2-1S(1)223970.1
470.0 3.35
4.8 ±±0.3244068.6[S II]
P3/2 – S3/2
24 571.5
4068.6 149.0 ± 0.4
A4068.5 41.7 ± 0.4
24076.3[S II]
P1/2 –4 S3/2
24 524.8
4076.3 50.1 ± 0.4
A4076.3 15.9 ± 0.4
4101.7HI
6–2
106 632.2
4101.8 13.4 ± 0.4
A4101.8 8.8 ± 0.4
4114.5[Fe II]b2 H11/2 –a4 F9/2
26 170.2
4114.5 2.7 ± 0.4
244211.1[Fe II]b H11/2 –a F7/2
26 170.2
4211.1 1.2 ± 0.4
4244.0[Fe II]
a4 G11/2 –a4 F9/2
25 428.84244.0 6.0 ± 0.5444276.8[Fe II]a G9/2 –a F7/2
25 805.3
4276.9 2.5 ± 0.5664287.4[Fe II]a S5/2 –a D9/2
23 317.6
4287.4 6.6 ± 0.4444319.6[Fe II]a G7/2 –a F5/2
25 981.6
4319.6 1.9 ± 0.3
4340.5HI
5–2
105 291.7
4340.4 25.0 ± 0.4
A4340.5 18 ± 0.4
4346.9[Fe II]
a4 G11/2 –a4 F7/2
25 428.84346.6 0.7 ± 0.3
4347.4[Fe II]b4 D5/2 –a4 D5/2
31 387.94347.6 1.0 ± 0.3
4352.8[Fe II]
a4 G9/2 –a4 F5/2
25 805.3
4352.8 1.0 ± 0.3
664359.3[Fe II]a S5/2 –a D7/2
23 317.6
4359.3 4.8 ± 0.3
334362.7CI
D3 – D3
87 002.3
4362.5 0.8 ± 0.3
114363.2[O III]S0 – D2
43 185.7
4363.4 3.0 ± 0.3
664413.8[Fe II]a S5/2 –a D5/2
23 317.6
4413.8 3.1 ± 0.3
464416.3[Fe II]b F9/2 –a D9/2
22 637.2
4416.3 4.4 ± 0.3
464432.4[Fe II]b F5/2 –a D7/2
22 939.4
4432.7 1.6 ± 0.3
444438.9[Fe II]b D1/2 –a D1/2
31 368.5
4439.2 1.0 ± 0.4664452.1[Fe II]a S5/2 –a D3/2
23 317.6
4452.2 2.8 ± 0.4464457.9[Fe II]b F7/2 –a D7/2
22 810.4
4457.9 2.3 ± 0.3
34471.7He I
D1 –3 P0
191 444.6
4471.5 4.0 ± 0.4
24571.1aKI
D3/2 –2 P3/2
34 913.3
4571.5 3.2 ± 0.4
a
3
Mg I]
P1 –1 S0
21 870.5
A80, page 20 of 4571.1
224580.8[Cr II]a4 P3/2 –a6 S5/2
21 824.1
4580.8 1.4 ± 0.3
34607.0[Fe III]F3 –5 D4
21 699.9
4606.8 1.0 ± 0.3
34658.1[Fe III]F4 –5 D4
21 462.2
4658.1 9.2 ± 0.3
air (Å)【結果】
輝線を同定してフラックステーブルを作成

Eu (cm 1 )Type

Table A.1. ̶ ESO Hα 574で検出された輝線とそのフラックス

ESO-Hα 574
・jet lineは多いがaccretion lineは少なく弱い・H, O, Fe, S, N, He, Ca, H2など数十の輝線・弱いながらもたくさんのBalmer線, Paschen
ラインはγ, βのみ, Brackett線はなし
・Ca IIの三重線はCTTSよりも弱い
CTTSの典型ではないがedge-onに近い円盤をもったCTTSの特徴(TWA30と似ている)
Par-Lup 3-4
・通常の降着トレーサー
・jetではいくらかの弱い輝線([OI]λ6300, [SII]λ6731など)
典型的なVLMSのスペクトル

Notes. Fluxes refer to knot A1, are in units of erg/s/cm2 ⇥ 10 17 and are not extinction corrected. Where the identified lines are also found inknots A, B or E the observed wavelengths and fluxes are also given. The regions over which the spectra for A1, A, B and E are extracted are shown(a)approximately 150 and 25 for ESO-H↵ 574 and Par-Lup 3-4, reWhelan
et obtainedal.: X-Shooter
observationsof the ESO-H↵
574 and
Par-Lup
3-4 Jets
spectively. BacciottiE.etT.al.
(2011)
valuescomE. T.these
Whelanet al.:by
X-Shooter
observationsof theESO-H↵
574 and Par-Lup 3-4 Jeparing luminosities corrected for any relative extinction to thebserved forarcsecond
jets average
centredonthe starofinthe caseof
3-4,
is conluminosities
young
stars
ofstarsimilarmass
insame 3-4, is conarcsecond
centredonthePar-Lup
in thecaseof thePar-Lup
ature of a magsidered.
Previously,
theemissionextracted
over the whole
spa-over the
減光量の評価(質量降着率の見積もりに必要)
starformingsidered.
regions.
Regardingthe forbiddenlines generated
in whole spaPreviously,
theemissionextracted
tial range
the however,spectrum
wasof
considered.
Thirdly,
theofnewes-Thirdly,
theofjets,
we
note
that
from
analyses
infraredlines the new es同じスペクトル型の減光なしClass
IIIテンプレートとスペクトルを比較
tial
range
the
spectrumwasconsidered.
+timates of Fe
relative
extinctionin thatthe aboveparagraph, Giannini
et ofal.described(2013) findtheseregionsnotaf- paragraphtimates
relative
extinction
described
inare
the
above
(i)
750
nmのフラックスでテンプレートを規格化
E.
T.
Whelanetal.:
X-Shooter
observationsof the ESO-H↵ 574 and Par-Lup 3-4 Jets
have been
appliedusing
curvesWeingartner
fectedby extinction.
Therefore
AV =the0ofis
assumedcurves
for&theof
diaghave thebeenextinction
applied using
extinction
Weingartner &Av=0–4.0
magの人工的な赤化を加えて観測の傾きを再現
Draine(ii)(2001).
Finally,
for conditions
the calculation
Lthe
nosticsofphysicalin theofjets
offromboth L
sources.
acc
line theDraine(2001).
Finally,forcalculationof Lisaccconfrom Lline thearcsecond
centred
on
the
star
in
the
case
of
Par-Lup
3-4,
newrelationships
described
inAlcalá
etal.
(2014)
areadopted.
ary for measur- Relative extinctionはESO-Hα
574extracted
で1.5
3-4ではなし
new relationships
described
inmag、Par-Lup
Alcalá
et al.
(2014)
are adopted.
sidered.
Previously,
theemission
over
the
whole
sparelationships These
consideracombination
of
all
the
accretionators), andThese
thus ・光度は上の減光補正を加えても同質量の典型的なものより小さい
relationships
considera combination
allesthe accretion
tialcalibrated
range
of on
theof
spectrumconsidered.
Thirdly,
the ofnew
4.5.
EstimatesṀacc forwas
indicatorssourceswhich
theBalmerjump
has
sing an edge-on
indicators
calibrated
on sources
for above
which paragraphthe Balmer jump has
timatessimultaneously.
of relative
extinction
described
in the・円盤による掩蔽のためのgrey
extinction(波長に依らない減光?)の影響
been measured
can be complibeen measured
havebeen appliedusing
thesimultaneously.
extinction
curvesof Weingartner&As
discussedintheIntroduction,
previous
estimatesof Ṁacc forThe
obscuration
e↵ect
of
the
edge-on
disks
of
both
sources
・ESO-Hα574では150倍、Par-Lup
3-4では25倍暗い
which can pro- Draine (2001). The
obscuratione↵ect of theofedge-on
disks
ofthe
both sources
Finally,
the calculation
Lacc fromL
ESO-H↵ 574
and for
Par-Lup3-4
were
re-analysed.
Ṁlineacc iscausing bothgrey
extinctioniscorrected
for
in
the
followingway.
At
endent of wavecausingdescribedgrey extinction
is corrected
for inaretheadopted.
following way. At
new relationships
in Alcaláet al. (2014)
0.5
calculated
fromtheṀequation,
fixed massEq.
(1)
implies
/
L
·
L
,
as
the
stellarradius0.5
Luhman (2007)
accEq. (1)accaimplies⇤
fixed massṀacc / ofLacc
as the stellar radius
These relationships
considercombination
all· Lthe⇤ ,accretion
質量降着率の見積もり
goes
as
the
squarerootoftheluminosity.
Assuming
that
the
same
isual extinction indicators calibrated
goes as theonsquarerootforof whichthe luminosity.
Assumingthat the same
sourcesthe Balmer
jump hasṀ=
1.25(L
R
)/(GM
)
(1)
により質量降着率を計算
obscuration
factorsuppresses
both
L
and
L
,
Ṁ
is
corrected
acc
acc
⇤
⇤
r-Lup 3-4 range⇤
acc
obscuration
factoraccsuppresses
bothLacc and L⇤ , Ṁacc is correctedbeen measured
simultaneously.
using:
used to estimate 以前に比べてThe obscuration
using: e↵ect of the edge-on disks of both sourcesLiscomputedLline
al. 1998).way.
Lline Atisn used in other whereacc
causing grey
extinctionfrom
iscorrected
for in theetfollowing
1.5(Gullbring・フラックスキャリブレーションが改善
1.5
Ṁ
(corrected)=(obscurationfactor)
⇤
Ṁ
(2)
acc
acc
Ṁacc(1)(corrected)=acc(obscuration⇤after
Ṁacc
(2)
line
luminosity
derived Ṁ
fromcorl. (2011), using the
fixed
mass
Eq.
implies/theLaccmeasured· L⇤0.5 ,f actor)
as fluxes
the stellar
radius→Lを計算
rectionextinction,and, M150the25
stellarradiusline
acc
d AV = 1.7where
mag ・天体に近い成分のみでL
⇤be
⇤ are
goes
asfor
the relativesquarerootof theluminosity.
Assuming
theobscuration
factor
taken
toR
for
where
the wasobscuration
factorwasandtakenthatto thebe same150 and 25 forandbothsuppresses
sourcesLboth
estimatedfromthe line
・新たな減光量を適用
line wasobscuration
factor
L
and
L
,
Ṁ
is
corrected
ESO-H↵
574mass.
and For
Par-Lup
3-4,
respectively,asdescribed
in
acc
⇤
acc
ESO-H↵ 574 and Par-Lup 3-4, respectively, as described influxesreported
TablesA.1 and A.2.
So forESO-H↵ 574 it isver more accuusing:
Section
4.4.
In
Fig.
8, in
Ṁacc
for InESO-H↵
・LlineからLaccを求める際に新たな関係を適用
Section
4.4.
Fig. 8, 574
Ṁaccandfor Par-Lup
ESO-H↵3-4
574 and Par-Lup 3-4
thewithluminosityof thelines
extracted
from1.5theregionof knot A1.
This was is
done
shown,
the blackandredpointsresultingfrom
thecalis shown,with the blackand redpoints resulting
from the calṀ⇤accand(corrected)= (obscuration
fandactor)
R⇤ aretaken
at 0.5 M
0.13 ⇤MṀ,accand
atにより掩蔽の影響を補正
0.12 R(2)pectra of young
culationMbefore
and
after
correction
fortheobscuration
factor.
culation before and after correction for the obscuration factor.
and0.18
R obscuration, for ESO-H↵
574wasandput
Par-Lup
3-4,
respectively.
al. 2013) ofConsideringthe 掩蔽の補正後の降着率は
where
factor
taken
be
150
and
thethecorrected
values,thetomean
accreConsidering
theresultscorrected
values,results
put 25
theformean accre1
The
values
of
L
and
T
used
in
the
computation
of
R
are
taken
by applyinge↵± 0.453-4,
⇤
574
and
Par-Luprespectively,asdescribed
intiontorate ESO-H↵
at log ( Ṁacc
)
=
9.15
M
yr
for
ESO-H↵
574
1
ESO-Hαtion rate
at log ( Ṁacc ) = 9.15 ±0.45 M 574
yr for ESO-H↵ 574
1 8, Ṁ
from
Merín
et
al.
(2008)
andLuhman(2007).
he flux andand(ii) 9.30
Section
4.4.
In
Fig.
for
ESO-H↵
574
Par-Lup3-4± 0.27 Mandyr 9.30
for ±Par-Lup
3-4.
are
the
1The errorsandacc
0.27 M yr for Par-Lup
Par-Lup 3-4.
3-4 The errors are thening and reproThe new
estimatesfrom
those presented
in Bacciotti
shown,
the
black
and
red points
resultingfromthe cal1 valueismarked
inwith8. Ṁ
forESO-H↵
is forat thelower
accdi↵er
1Fig.value
markedin Fig. 8.574
Ṁacc
ESO-H↵
574 is at the lower
culation
beforeandafter
correction
fortheobscuration
factor.
latter correctionetal.
(2011)
due
to thefollowing
factors.
Firstly,
improvements
end of the
range
ofend
valuesofṀmeasured
for
other
K
type
acc
of the range of values of Ṁacc measured for other K type
Considering
the1998),
corrected
put thanthe meanaccretainties in specin
the flux etcalibration
mean
that
lineresults
fluxesandtherefore
YSOs (Gullbringal.
and values,a etfactorten lower
the values
YSOs(Gullbringal.
1998),
and1 a factor ten lower than the
tion
rate
at log9.15±(2011)0.45ofM
for ESO-H↵
574
d globally to
an of
Lline
have
changed.
Early
versions
theyrobscuration
X-Shooterpipelineacc ) =
approximate
estimatein( Ṁ
Bacciotti
et
al.
with
approximate
estimatein
Bacciottietal.
(2011)
with1
andnot9.30±a 0.27
yr calibrationfor
The previerrors
theobscuration
specially strongdido↵er
reliableflux
but 3-4.
thishasaresinceconsidered.
The
resultfor MPar-Lup
3-4Par-Lup
is closer
to situation
ourA&A 565, A80 (2014)
Par−Lup3−4

ESO−Ha 574 A1

100.001012 cm−3, 1000 K
1012 cm−3, 10000 K
108 cm−3, 1000 K
108 cm−3, 10000 K
104 cm−3, 1000 K
104 cm−3, 10000 K

2500K
3000K
5000K

Fn / Fβ

10.00

20000K

1.00

100.0

2000K
2500K
3000K

10.0

Fn / Fβガスの物理性質と質量放出率
Balmer decrementsを計算
輝線比を上準位の量子数の関数としてプロット
・Case B (赤・黒)・光学的に薄いLTE (青)と比較 (光学的に厚い場合も試したが合わなかった)

2000K

5000K
20000K

1.0

LTE
opt.thin

0.10

LTE

0.1
0.01
0

2

4

6

8

10

12

14

opt.thin

0

nup

2

ESO−Ha 574 A

Fig. 10. Balmer decremen
ESO-Hα 574 (A1, A)
the source spectrum. The
100.0010 cm , 1000 K
nation (black and red line
10 cm , 10000 K
The best fit is o↵ered by・Case B (104 cm-3, ~10000 K) がよく合う (A1はTe=20000
10 cm , 1000 K
10 000 K. This is consiste
10 cm , 10000 K
10.00
KのLTEでもよい)
10 cm , 1000 K
10 cm , 10000 K
The optically thick cas
4-3
not provide a good fit t
1.00
・密度10 cm , 温度10000–20000 Kはjet起源と矛盾しない
For regions A1 andB model for n = 10
・A1はESO-Hα 574自身を含んでいるが放射の大半はdecrements, as well as,
0.10
T = 20 000 K. A dens
outflowからきているof 10 000–20 000 K isagrees with results pu
0.01
2011, 2008). This ana
0
2
4
6
8
10
12
14
Par-Lup 3-4
emission from A1 comnup
the appearance of the
A&A 565, A80 (2014)
・Case B (108–10 cm-3, ~10000 K)
Fig. 9. Balmer decrements for ESO-H↵ 574, regions A1 (top) and A astrometric analysis of
The gas from Par(bottom), fittedwith Case B recombination (black and red lines) and
ESO−Ha 574 A1
Par−Lup3−4
LTE optically thin models (blue line). For the Case B scenario, the best B models, a high den
・Magnetospheric accretion flow起源
fit is 100.0o↵ered by n = 10 cm , T = 10 000 K. The LTE optically thin ture of ⇠10 000 K. Su
100.00
2000K

12

2500K
3000K

Fn / Fβ

5000K

20000K

−3

12

−3

8

−3

8

−3

4

−3

4

−3

e

LTE

4

opt.thine1012 cm−3, 1000 K
1012 cm−3, 10000 K
108 cm−3, 1000 K
108 cm−3, 10000 K
104 cm−3, 1000 K
104 cm−3, 10000 K

2500K
3000K
5000K

Fn / Fβ

10.00

20000K

1.00
LTE
opt.thin

0.10

0.01
0

2

4

6

8

10

12

14

nup
ESO−Ha 574 A

100.00

2000K
2500K
3000K

1012 cm−3, 1000 K
1012 cm−3, 10000 K

2000K
e

4

3

e

case with T e = 20000 K provides an equally good10
fit. Theagree
cm results, 1000 K
2500K
with the fact thatfor region A1, emission from thezone
cm , 10000
K is
10 accretion
3000K
dominated by the outflow emission.
10 cm , 1000 K
5000K
10 cm , 10000 K
10.0

Fn / Fβ

2000K

20000K

12

−3

12

−3

10

−3

10

−3

8

−3

10 cm , 1000 K

K A
10 cm , 10000
both sources and the Balmer decrements for ESO-H↵
574 A1,and Par-Lup 3-4 were computed (with respect to H ). The line
1.0
ratios are plotted as a function of their upper quantum number(nup ) and results are presented in Figs. 9 and 10. While the specLTE
tral range of X-Shooter
covers all of the Balmer lines, any lines
with nup0.1> 13opt.thin
were found to be too noisy for inclusion in theanalysis. Additionally, at the intermediate spectral resolution of
X-Shooter the H7 and H8 lines are blended with other lines andare therefore
It6 is noted
only the12brightest
0 also 2not included.
4
8 that 10
14
Balmer lines were detected in knot A. To constrain the physn
upical conditions in the emitting gas the decrements were firstly
compared
to standard
Case
predictions(red andcurves)Fig. 10. Balmer
decrements
forBPar-Lup
3-4. Fluxes
are blackextracted
fromcalculated
for a range
temperatureand density.
Case Bthe source spectrum.
Theofdecrements
are fittedwith CaseThe
B recombicurves
were andderivedusingand
theLTE
calculationsof Hummer& Storeynation (black
red lines)
optically thinmodels (bluelines).
3
The bestand
fit istheo↵eredtheprovided
Case B scenarioand &ne Hummer= 1010 cm(1995).
, Te =(1987)
data byfiles
by Storey
8−3

an origin in magnetos
For example, Muzero
spheric accretion in T
magneto-spherically ac
for 10 6 M yr 1 Ṁatemperatures correspon
rates. The origin of thgas is in agreement wi
H↵ line (Bacciotti et al
The temperature e
the results of Bary et aof the gas emitting theof TTSs, was best fittethat while this value ismagnetospheric accretallow for the absorptio
the hot corona and/or aof heating and ionisat
can still be consistentgas, as the ionising ph

電離輝線の診断
Post-shocked ガス内で物理状態の勾配？
観測の輝線比と5準位モデル原子の数値計算から予測される値を比較
ESO-Hα 574
・1.4x104 K, 水素電離率 > 0.65, BとEでより高励起・電子密度は内から外に向かって4500 cm-3から160 cm-3へと減少
Par-Lup 3-4・電離度は低い xe=2–6×10-3, 温度と電子密度は高い(2.5–3×104 K, 4000 cm-3)
ESO-Hα 574のほうが高電離の輝線が豊富なのはこのため・励起度は高いが電子密度が低いためにBとEでは高励起輝線は検出されなかった・BとEはshock速度がA1とAおよびPar-Lup 3-4よりも速い
E. T. Whelan et al.: X-Shooter observations of the ESO-H↵ 574 and Par-Lup 3-4 Jets
ESO-Hα 574••

Par-Lup 3-4

JetにおけるOutflow rateのよりよい診断
5段階モデル原子を用いる

E

A1

A

B

R

B

Fig. 11. Various line ratios tracing the excitation, ne , T e and xe of the gas in ESO-H↵ 574 A1, A, B and E and the Par-Lup 3-4 blue and red jets areplotted. As some lines were not detected in B and E and in the Par-Lup 3-4 jets upper limits of 3 were assumed.

The
valuesof. nThis
T edonegiventhroughin Tablea1comparisonarenow6583
usedorestimate
Mout
ofto
the
observed[SII]
linesvelocity
for example.
These
lines
e andisin[NII]
ESO-H↵
574
are6731highershocksthan
A1,warrant
Aetandthe andstella1
accretion
onto
the
star
(Fischer
al.
2008)
in
(Edwards
etal.
2006).
Confirmationo 60 km s 1 are indicated,
while
speeds
reach
⇠90
km
s
in
B
mate
Ṁ
.
This
is
done
througha
comparisonoftheobserved
extra
discussion.
out
and theoreticalline luminosity producedin 3-4a givenforbidden
Par-Lupjets.
Thisexplains
why
ESO-H↵
574
hasa
spectrum(Edwardstraces
et al. 2006).
Hethe
I 1.08nd E of ESO-H↵ 574.
disk Confirmation
winds came thatfromstandline.
theoreticalproduced
a givenforbidden
To this line
aimluminosity
the luminosity
LSII in
of
[SII]
line
was
much
richerin6731lines
highionisedspecies,suchas[OIII]
and
HeIthe
1.083
µm:
the
Hefrom
I 1.083µm
line
traceshighexcitation
retraces
disk
winds
came
from
the
study
of
the
the3-4detectionofHartigan
PC temprofiles
andfroliOline.
To this aim
luminosity
LSII ofthe [SIII],
[SII]
6731
line energy
wasin
4.7.
consideredin thetheseparateknots
makingup
the
ESO-H↵
574
outnotdetectedthe
Par-Lupoutflow.
et
al.
gions
(excitation
of
He
I
at
1.083
µm
is
20
eV)
and
E. T. Whelan et al.: X-Shooter observations of the ESO-H↵ 574 and3-4 Jets
thePar-Lup
detectionofPCexample
profiles andfrom spectro-as
inin
thetheseparateknots
makingup
the
ESO-H↵
574
outies.
For
Azevedoal. (20.6.3. Estimates ofconsideredṀflowout /Ṁandacc の見積もり
peratures
inFollowing
excessof ⇠15000For
K.lineThisline hasbeenwelletstudied
alsoplot thevariousratiosasa functionof(2007)
shocketuse
Par-Lup 3-4 blue and(1994)
red jets.
Nisiniies.
exampleAzevedo
al.
specs
The
flow
and 574in the Par-Lup 3-4 blue and red
jets.
Following
Nisinitodetectanextended
disk
wind
in
ESO-HαPar-Lup
3-4
in
TTSs
and
has
been
found
to
have
contributions
from
both
the
velocity
(see
their
Figs.
10
to
14).
We
used
these
modelstoeset
al.
(2005);
Podio
et
al.
(2006)
weuse
therelationshiptodetectanextended
disk
wind
in
the
jet-less
C
表の電子密度と電子温度から質量放出率を見積もる
The values of ne and
T
given
in
Table
1
are
now
used
to
estilineseet al. (2005); Podio et al. (2006) we use the relationshipObservations
of
the
He
I
1.083
µm
accretion
onto
the
star
(Fischer
et
al.
2008)
and
inner
disk
winds
and foundofthat
for Ithe
Par-Lup
3-4in jets
Observations
the He1.083µm linehamate Ṁout . This is done through a comparison of the observed timate the shock velocities
some(Edwards
etal.
2006).
Confirmation
that
He
I
1.083
µm
emission
1
Ṁ
=
µ
m
(n
V)
v
/l
(3)
tigatedfar
lessfrequently.
Takamiepoutt
574 inner
shock
speeds Takami
of 30 kmnd theoretical lineṀluminosity
int at given
forbidden jets and ESO-H↵(3)(nHH V)Hvt /ltigatedknots,
far lessfrequently.
et sal. (2002)
out = µ mH produced[NII]
traces disk winds
came from the study
of ofthethelineHeprofilesi.e.1 µm emission
1
grams
Iemission
1.083
⌘was
ne. To this aim the luminosity LSII of⇣ the ⇣[SII] 6731
line
grams
of
the
He
I
1.083
µm
to
60
km
s
are
indicated,
while
speeds
reach
⇠90
km
s
ininBthe DG1⌘i1extraTi
X X
thedetectionofPC
profilesandfrom
spectro-astrometric
studX
X
with
n
V
=
L
h
⌫
A
f
.
Here
µ
=
1.24
is
the
mean
gue thatshock
heating
be thedonsidered in the separate
upout- µand
SII⌫thewith nknots
LSII hAi ESO-H↵
fi X i Hi X574
. HHere= 1.24isthe meanthat shockheating
must
be the must
dominantheat
EFor
ofexample
ESO-H↵
574. gue
H VH=making
ies.
Azevedoetal.
(2007)
use
spectro-astrometry
HeI
1aow and in the Par-Lup
3-4weight,
blue
andmredFollowing
Nisiniinorderthe jetordertheto explain
prese
in windthe jetinin
toin
explainpresencetheof suchatomicweight,mjets.
theprotonmass,
the volumee↵ectively
atomicmass,
V thetoV
volumee↵ectively
H protonH thedetectanextendeddisk
thejet-less
CTTS
TW
Hya.
t al. (2005); Podiofilled
etfilled
al. (2006)use
the gas,relationshipetalsoal.
(2012)
alsodeet(2012)
detect He
Igions1.08
by by
the weemitting
vt andvtltand
the ltangential
velocity
andtheemittinggas,
tangential
velocity
andEllerbroekt theObservationsof theHe
I line.
1.083µm line.
line inEllerbroek
jetsal.have
been
inves質量放出率と円盤による掩蔽を考慮した質量降着率から4.6.3.
Estimates
of
Ṁ
/
Ṁ
peratuinleveltheaccHHet
and
1043
jets.
Thus
ESO-H↵
in1042the(2002)
HHHH
1042
and
HH
1043
jets.
Th
of theknot,knot,
Ai and
rate
and
upper
level
lengthof theAi fandradiative
rate
andupperout
i the fradiativeṀout = µ mH (nH V)length
vt /ltigatedfar
lessfrequently.
Takamial.
present
PVdiai the(3)
t
i
A
E
A1
B
i
X
in TT
ofgivena Xfew in
jets
He
Inow
1.083
µm
has
beenµmidentifi
・Par-Lup
modelと一致
ofTablefew
where
HearI 1.083
ha
Rtransition
relative
tolaunchingthetransition
and
grams
of thefinallyHe
inawheretheDG
Taujet. They⇣population⌘3-4はjetThevaluesofnIeB1.083
andTµm
1 jetsareusedtoestipopulation
to considered
the considered
andfinallye emission
1 relative
X
i
X
X
X
would
thatexpect
futureof
X-Shooter
observations
accrewith nH V = LSII andh・ESO-Hα⌫ AiX fi are.ionisation
Here µ = fraction
1.24 is themean
gue thatshock heating
must
be theexpect
dominant
heatingmechanism
would
that
X-Shooter
XX Hthemate
Ṁabundance
isofdonethrougha comparison
thefutureobserved
andtherelativeout . This
andare574は大きい誤差があるがこのモデルを棄却はできない
the ionisation
fractionandtherelative
abundanceof
H H
increasethenumberofjets for
which
this
line
is
(Edw
in thetheoretical
jet thein orderto as
explainthe presenceof sucha ahighly
excitedtomic weight, mHthe
theconsidered
proton mass,
V the The
volumee↵ectively
increasethenumber
of
jets
for
whic
species.
tangential
velocities
of
knots
and
line
luminosity
producedingiven
forbiddentheThetangential
velocities
ofettheknots
asPV alsodiagramofthe
ESO-H↵
574
He I 1.083µm
tracesline.
Ellerbroek
al. (2012)detectHe
I 1.083µm6731emissionlled by the emitting
gas,consideredvline
lt tracing
the species.
tangential
velocity
andxusedt andPVdiagramofthe
ESO-H↵
574
H
estimatedinratios
Bacciotti
etexcitation,al. (2011)
are
here
and
the
knot
line.
To
this
aim
the
luminosity
L
of
the
[SII]
line
was
Fig. 11.
Various
the
n
,
T
and
of
the
gas
in
ESO-H↵
574
A1,
A,
B
and
E
and
the
Par-Lup
3-4
blue
and
red
jets
are
SII
e
e
e
estimatedin
Bacciottietal.
(2011)
are
used
here
and
the
knot
regionispresented
in
Fig.
1.
Both
knots
A1
and
in
the
HH
1042
and
HH
1043
jets.
Thus
ESO-H↵
574
is
only
one
ength of the knot,
A
and
f
the
radiative
rate
and
upper
level
the
i
i
plotted.
As somearelinesgiven
were notin
detected
B andThese
E and invalues
the Par-Lup
jetsinupper
limitsassumed.
lengths
Sect.in4.1.
are3-4also
given
in of 3 wereregionup
is thepresented
in574
Fig.
1. Both dk
consideredtheseparateknots
makingESO-H↵
outi
X
lengths
are
given
in
Sect.
4.1.
These
values
are
also
given
in
while
B
and
E
are
not
seen
(see
discussion
aboveof
a fewfromjets where
He I 1.083 µm has been identified to date. We
ies.
F
opulation relativeTable
to the1.consideredandabundance
finally X isflow
The S = S+transition, elemental
taken
Asplund
while
B
and
E
are
not
seen
(see
disc
and
in
the
Par-Lup
3-4
blue
and
red
jets.
Following
Nisini+wouldisexpectthatfutureX-Shooter
observationsof emission
jets wouldline profileshow
the HeI 1.083 µm
an
Table=variables
Srelative, elemental
abundance
taken
from
Asplund
1. Valuesof1.
theThedi↵erent
used
in Eq. (3).
to dend HX are theTable
ionisation
fraction
andStheabundanceetal.
(2005),
thelevel
population
isofcalculated
numericallyetincreaseal. (2005);
Podio
et
al.
(2006)
weuse
therelationship
show
the
He
I
1.083
µm
emissionlthe numberof jets
for whichthis line
is observed.
The
regionextracted
fromthe 2D X-Shooter
spectru
al. tangential(2005),
and
the
level
calculated
numericallyObse
he considered species.
The
velocities
of(2006).
thepopulation
knots
as thatisthe
as etdescribed
in Podio
et
al.
Note
new
estimateregionextracted
fromthehas2D
X-Sh
PV
diagramofthe ESO-H↵
57410 tracing
He
I 1.083µmṀisemission
line
00
3
4
1 the
the
line
jet
narrowanda velocit
Objectvt (kms 1 )used) etandnal.
T e (10K)firstly
xe to
Ṁ
M yr
)
aset describedinPodio(2006).
Note
that
the
new
estimatee (cmout (10out / Ṁacc
stimated in Bacciottial.accurate(2011)
are
here
the
knot
Ṁ
=
µ
m
(n
V)
v
/l
(3)
tigateismorethaninlt (Bacciottiet )al.
(2011),
due
the
out
H H
t t
regionispresented
in
Fig.
1.
Both
knots
A1
and
A
are
detectedtheline
tracingthejetisnarrow
and
ESO
HA
574
(as
seen
for
the
other
jet
lines).
An
IPC
profileengths are given in
Sect.
4.1.accurate
These
values
are
also
given
inetoriginate
is
more
than
in
Bacciottial.
(2011),
firstly
due
to
the
gram
new
valuesoftheobserved
fluxes,
that
from
an
im⇣
⌘
Knot A1
130 ± 302.84500 ± 225
1.4 ±while
0.2 B and>0.65
5.0
± 1.5
1 seeni discussion
E00 are not seen
(see
above).
In Fig.
5 we
(as
for
the
other
jet
lines). An
veal
important
information
about
accretion
X
X
Table 1. The S = S+proved, elemental
abundance
is taken
00
with
n
V
=
L
h
⌫
A
f
.
Here
µ
=
1.24
is
thethemean
newvaluesobserved
fluxes,that
originatefromanimgue isththefluxfrom
calibration
and
the
new
extinction
Knot
A treatment
130 of
± of
30the
2.4
2100
±Asplund
105
3.7
±
1.1
0.5
(+1.0)(
0.2)
H
SII
i
i
X
H
show
the
He
I
1.083
µm
emissionline
profileandtheemission
00
00
For
example,Fischeretal.
(2008)
discusshowveal
important
informationabout
tht al. (2005), and thelevel
populationiscalculated
numerically
Knot
B
220
±
50
2.2
230
±
11.5
8.2
±
2.5
estimates.
most
importantlyṀoutisextracted
nownew
calcuin theproved Secondly,
treatment andof theflux
calibration
andtheatomicweight,theprotonmass,
Vspectrum.
thevolumee↵ectively
fromthesorption
2D3.5
X-Shooter
The
part
of
Hextinction
00region
00 m
KnotetE al. (2006).
325 Note± 50 that
2.3the new
160 ±estimate
8.0
± 1.1
0.3the(+0.6)(
0.1)
features
in
He
I
1.083
µm
line
ca
s described in Podio
For
example,Fischeretal.
(2008)
1
lated
using
theSecondly,
jet parameters
deriveddirectlyfromthespectra,line.
Eestimates.
and
mostimportantly
Ṁ
is
now
calcufilledbytheemitting
gas,
v
and
l
the
tangential
velocity
and
the
line
tracingthejetisnarrowhasavelocity
of
⇠0
km
s
out
Par-Lup3-4tt
mationabout
theaccretiongeometry
of
YSOs.
Sp
s more accurate thanin
Bacciottietal.
(2011),
firstlyduetothesorptionfeatures
in
the
He
I
1.083
3
rather
fromapproximate
onof
the
value
Blue
170the
±an
30jet
1.9
4300 ±formula200
3.1based±(as
0.5seen
(2.0
±from
0.2)
⇥other10ofAspectra,
0.3±An
0.1radiative
0.05rate(+0.10)(in theforthelines).
IPC
profilecan
re- level accretio
latedthanusingparameters
deriveddirectlythelengththeknot,
fi theand 0.02)
upper
ijetandew values of the Red
observed
fluxes,that
originatefromanimprobed
thegeometry
ofwhichmagnetospheric3mationabout
theaccretiongeometryi
170 ± 30at an1.8
4600 ±unknown
300
2.4 ±veal
0.6 important(6.3 ± 0.2)
⇥
10
0.29
±
0.09
0.05
(+0.10)(0.02)
therathercriticalthan
density
assumedtemperature,
as
was
X
about
accretionis alsoand
detected.
ofa fI
fromandanthe
approximate
formulabasedrelative
oninformation
the tovalueof thethe
populationtheconsideredtransition
finally
roved treatment of the flux calibration
new extinction
modelling
red-shifted
absorption
at
the
He
X
probedthegeometryofmagnetosph
done
for outflows
stars of very
mass.Fischer etasal.was
For low
example,(2008) discuss how red-shifted abwould
X temperature,
thepreviously
criticaldensityatoutanisfrom
assumedstimates. Secondly, andmostimportantlyṀ
now calcu-unknown
via
scatteringofthestellarandinforveilingabsorptio
continua
and
are
the
ionisation
fractionandtherelative
abundanceofsorptionfeatures
in
the
He
I
1.083
µm
line
can
revealmodellingthe
red-shifted
These
updatedmeasurements
of
Ṁ
and
thus
Ṁ
/
Ṁ
H
out higher density.
outThe acc
3-4 haspreviously
aderived
lower valueof outflowsṀfromsuggests
ratio [NII]6583/[OI] 6300,
proportional
to hydroincrea
acc , whichated using thePar-Lup
jet parameters
directly
the spectra,
done
for
from
stars
of
very
low
mass.
the
red-shifted
absorption
featureis sensitive
totheconsidered
species.
The
tangential
velocities
of
the
knots
as
mationabout
theaccretiongeometryof
YSOs.
Specifically
they
are
given
in
Table
1.
Ṁ
/
Ṁ
is
derived
for
each
lobe
of
the
via
scatteringofthestellarveil
temperatures
than forformula
the sources
in ontheaccstudy
of ofBary et al. gen ionisation fraction xe , indicatesxe > 0.65 in the
flow,
and
in- andPVoutdather than froman approximate
based
the
value
These
updatedmeasurements
ofestimatedṀout theandthusṀofaway/ Ṁal.
ofthemagnetosphere
of
the
star
and
the fillingoutmagnetosphericacc(2011)
ingeometry
Bacciotti
et
are
used
here
and
the
knot
probedaccretion
in
CTTSs
by
(2008).
creasesmovingfromthesource.
This
isinagreement
with
ESO-H↵
574
and
Par-Lup3-4
outflows
bydividing
by
the
avthe
red-shifted
absorption
featurehe critical density at are
an assumed
unknowntemperature,
as wasregio
accretion
shock.
Finallysomeblue-shiftedemisi
given ofin
Table
1.
Ṁ
/
Ṁ
is
derived
for
each
lobe
of
the
out
acc
the
value
found
on-sourcebyGianninietal.
(2013).
The
highexlengths
are
given
in
Sect.
4.1.
These
values
are
also
given
in
modellingthe
red-shifted
absorption
at
the
He
I
1.083
µm
line
erage
value
Ṁ
.
For
the
ESO-H↵
574
blue
outflow,the
1 magnetosphere of the staroftostheone previously for outflows from starsacc
of very low mass.
whileaf
+and
to
⇠200
km
is
also
observed.
See
Sect.
4.2.2
citation
of
B
E
would
seem
contradictthenon-detection
ESO-H↵
574
and
Par-Lup3-4
outflows
bydividing
by
the
avvia
scatteringofthestellarandveiling
continua
and
found
that
Table
1.
The
S
=
S
,
elemental
abundance
is
taken
from
Asplundaverage Ṁout of A1,
From Table 1 it can
˙ A and B ˙is used.
˙Multi-epoch Sub-arcsecond [Fe II] Spectroimaging of the
DG Tau Outflows with NIFS. I. First data epoch
M. C. White, P. J. McGregor, G. V. Bicknell, R. Salmeron and T. L. Beck

DG Tauの分子流のGemini/NIFSによる面分光観測【背景】
YSOにより駆動される分子流は原始星近傍で起きる降着・放出過程の重要な手がかり

星のスペクトル
1684
M. C. White et al.
広がった成分を調べたいのでこれを除去

Downloaded from http://mnras.oxfordjournals.org/ at National Ast

Figure 1. The stellar spectrum of DG Tau. The spectrum is extracted from a 0.8 arcsec diameter circular aperture centred on the star. Panel (a) shows the
H-band (1.49–1.80 µm) spectrum, and panel (b) shows the K-band (1.99–2.40 µm) spectrum. The flux density has been normalized to unity at 1.60 µm in
the H band and 2.2 µm in the K band. Prominent emission features are labelled, as are absorption features, which are used to determine the accuracy of the
wavelength calibration (Section 3.4). The CO bandheads visible in the K-band spectrum are also marked. The K-band data have been presented previously by
Beck et al. (2008).

adjacent to the spectral region of interest around the H2 1–0 S(1)line, averaging them and subtracting this averaged continuum image
from each wavelength plane of the data cube.

3.3 Circumstellar environment[Fe II] 1.644 µmのチャネルマップ
Fig. 2 shows channel maps of the circumstellar environment of DG
Tau, as seen in [Fe II] 1.644 µm line emission, with the stellar and
3つの成分
unresolved line emission components removed. The top-left and(i) よくコリメートした高速度Blueshiftジェット
bottom-right frames show the velocity ranges used for continuum
scaling. Here, and in all subsequent figures, the outflow axis is(ii) Blueshiftした中速度のあまりコリメートしていない縁の明るいV字構造
labelled as x, and the axis across the outflow as y. The data have been(iii) Redshiftした分子流Figure 2. Channel maps of the DG Tau outflow. Panels show images of the extended [Fe II] 1.644 µm line emission around DG Tau, binned into 40 km s−1 -wide

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on May 26, 2014(ii) An intermediate-velocity, less-collimated, edge-brightened,‘V’-shaped structure in the blueshifted outflow. Within approximately 1 arcsec of the central star, the outer edges of this structure are linear, and are aligned radially with respect to the central
star. The opening half-angle of this feature is 15◦ ± 1◦ . AgraAmboage et al. (2011) obtained an opening half-angle of 14◦ forthe same structure from SINFONI data of DG Tau obtained on 2005
Oct 15. This structure is ‘pinched’ ∼1 arcsec from the star, and thenre-expands with increasing distance from DG Tau (Fig. 2, panel[−180 : −140] km s−1 ).
(iii) A redshifted outflow, which becomes visible approximately
0.7 arcsec from the central star. The inner region of this structure is近づく(Blueshift)成分 (下図右)・高速ジェットは歪んだ軌道
・3つのノット(A, B, C)
遠ざかる(Redshift)成分 (下図左)・Blueshift成分とちがって速い成分や尾根がない・泡状の構造. The difficulty in accurately fittings in a larger uncertainty in the knot

hologyce of the DG Tau redshifted outflow,from that of the blueshifted outflow
ly discernible fast outflow, nor ridge-

1687

obscured by the circumstellar disc around DG Tau. We estimate theradial extent of the obscuration, and hence of the DG Tau circumstellar disc, to be ∼160 au, after correction for the inclination of thejet–disc system to the line of sight (38◦ ; Eislöffel & Mundt 1998).
This is in agreement with the measurement by Agra-Amboage et al.
(2011), who used this obscuration to place limits on the disc modelsof Isella, Carpenter & Sargent (2010). The redshifted outflow takes
the form of a bubble-like structure. The material with the greatest
receding line-of-sight velocity is concentrated on the outflow axis,
and at the apex of the bubble. The material along the edges of thestructure emits at progressively lower line-of-sight velocities withdecreasing distance from the central star.
3.3.1 Approaching jet trajectory

The high-velocity blueshifted jet does not travel a linear path, but
bends along its length (Fig. 2, rightmost-top panel). We define theridgeline of the jet as the location of the jet brightness centre at eachposition along the outflow axis. Single-component Gaussian fitswere performed across the jet, at every recorded position along theoutflow axis, on an image of integrated [Fe II] 1.644 µm emissionline flux formed over the velocity range −300 to −180 km s−1 . Thisvelocity range was chosen so that the ridgeline was computed forthe highest-velocity gas, corresponding to the high-velocity jet (seebelow, also, Pyo et al. 2003b). The ridgeline computed from thesefits is shown in Fig. 3(b). The uncertainties in the lateral position of
the fitted ridgeline are of the order of ±0.01 arcsec.
The jet ridgeline was fitted in spatial coordinates with a simple sinusoidal
functionorderµm
to characterize
the nature of the jet trajecFigure 4. DG Tau receding outflow.
(a) Integrated[FeinII] 1.644
line fluxof the receding outflow from DG
Tau.TheTheamplitude
line flux is computed
oversinusoid
the
tory.
of the fitted
is 0.027 ± 0.001 arcsec ≈velocity range 0–300 km s−1 . (b)3.8
Contourplot theof thewavelength
same integrated
II]au, andis [Fe1.035± 0.006 arcsec. Deprojecting
line flux. Contours are labelled in units of 10−15 erg cm−2 s−1 arcsec−2 . The
thisdistancetoaccountforthejet
inclination to the line of sight
position of the central star and the position and size of the occulting disca bothphysical
235 au. If the sinusoidal jet traused during the observations areyields
shown inpanels wavelength
by a yellow starofand
jectory is due to jet precession, the amplitude corresponds to acircle, respectively.

Figure 3. The DG Tau approaching outflow. (a) Integrated [Fe II] 1.644 µmline flux of the approaching outflow from DG Tau. The line flux is computed
over the velocity range −300 to 0 km s−1 . Knots A, B and C are labelled.
(b) Contour plot of the same integrated [Fe II] line flux. Contours are labelled in units of 10−15 erg cm−2 s−1 arcsec−2 . The unlabelled contour corresponds to 170 × 10−15 erg cm−2 s−1 arcsec−2 . Knots A, B and C are

Downloaded from http://mnras.oxfordjournals.org/ at National Astrono

a Gaussian profile, which is extended
h an axial ratio of ∼2.3. This knoty Agra-Amboage et al. (2011), who± 0.03 arcsec along the outflow axis
5 Oct 15. Our positions agree to 1σ .
than knots A and B, at ∼15 per cent2 and 3). As with knot B, knot C istion, but with an axial ratio of ∼1.7.
y Agra-Amboage et al. (2011), with a
05 arcsec along the outflow axis from
5. This agrees with our measurement
position is within their uncertainties.
e recent knot ejection history of DG

M. C. White et al.

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Ob

n of knot A have been observed preÅ (Solf & Böhm 1993), [O I] 6300
lley et al. 1997) and He I 10 830 Årmore, Lavalley et al. (1997) reportobserve at ∼0.15 arcsec ≈ 34 au deentral star2 exhibits very little proper
not represents a steady region in the
nced. A feature similar to knot A apof Agra-Amboage et al. (2011, fig. 3
ors did not mention it. We interpretin the jet, resulting from the recoland further on this interpretation in

1686

DG Tau outflows – I. First epoch静止ノットA
X-ray, FUVの観測も考慮するとJet recollimation shockの特徴
X-ray観測からpre-shock flowの速度は400–700 km/s
MHD円盤風を仮定
最も内側のJetの流線は中心星から半径0.01–0.15 AUの位置から出ている
Post-recollimation-shock jetは215 km/s以下まで減速されて[Fe II]のHVCとして観測
Jetは速度と直径を変化させながらまっすぐでない軌道
DG Tau outflows – I. First epoch
1691
回転しているかはわからない(←強いrecollimation shockを通過したときの効果)

Figure 9. Electron density measurements of the DG Tau approaching jet.
Black circles show the electron density derived in this work for the DG Tau
jet from the [Fe II] 1.533 µm/1.644 µm line ratio at each position along
the outflow axis, averaged over all spaxels within ±0.5 arcsec of that axis
in the cross-outflow direction. All determinations of electron density fromthe literature are made using the optical line ratio technique developedby Bacciotti & Eislöffel (1999), except for those by Agra-Amboage et al.
(2011), which use the [Fe II] line ratio technique. Where provided, electron
densities are quoted for high-velocity (HV), intermediate/medium-velocity(IV/MV) and intermediate-velocity (LV) components. Uncertainties are as
quoted in the relevant reference, except for where they have been estimatedfrom 2D maps of electron density (Bacciotti et al. 2000; Coffey, Bacciotti& Podio 2008).

Figure 10. Progression of knots in the approaching DG Tau outflow, 2005–
2009. Shown is a contour plot of [Fe II] 1.644 µm line emission fromthe approaching DG Tau outflow at 2005.87, 2006.98 and 2009.88. Images are formed by integrating over the velocity range −380 to 0 km s−1 .
Contours have levels of [25, 30, 50, 70, 100, 120, 170] × 10−15 erg cm−2
s−1 arcsec−2 . Short dotted lines represent the observation date of each epoch.

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of

2つのジェットノットB, C
ノット放出の間隔は周期的でないジェットの速度は215 km/sから315 km/sへと増加
速度の変化が動くノットの形成要因かもしれない[Fe II] 1.644 µmの第二の分子流成分を複数成分ガウシアンフィットにより分離
中速度成分は大きく開いた流れ
1688 M. C. White et al.
中心のジェットと広角の分子風の間のshocking,turbulent entrainment layerだろう

Figure 5. Two-component Gaussian fits to the spectra of four spaxels in the approaching DG Tau outflow. The spaxels are located at a distance of (a, b)
x = 0.62 arcsec and (c, d) x = 1.13 arcsec along the outflow axis. The spaxels shown in panels (a, c) are on the jet ridgeline at y = −0.03 arcsec, and the
spaxels in panels (b, d) are offset from the ridgeline at y = −0.26 arcsec. Actual data and uncertainties are indicated by black circles and error bars, fitted line
components are shown as blue dot–dashed and green dotted lines, and the total fit is shown as red dashed lines.

2
instrumental2
intrinsic

2
fittedIIhigh-velocity jet. Knots B and C are reproduced. Knot A is not
visible, as that region of the outflow is not fit due to its low signalto-noise ratio, which results from the proximity of the central star.
The ‘pinching’ of the outflow ∼1 arcsec along the outflow axis fromthe central star is also reproduced. We interpret this to be due to alack of emitting gas between the two jet knots, rather than an actual
narrowing of the jet.

3.4.2 Approaching IVC
The integrated line intensity image of the IVC shown in Fig. 6(d)differs significantly from that of the HVC (Fig. 6a). The emissionis spread further from the outflow axis than the HVC. Interestingly,
the edge-brightened ‘V’-shaped structure is not reproduced. This

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on May 26, 2014

【結論】!
以上の結果は時間変化するメカニズムに依存!
さらなるマルチエポックデータが必要である!
とくに!
Tau microjet that generally show increasing absolute line velocity
the fitted line width, via the formula FWHM
= FWHM−
with distance from the central star (Bacciotti et al. 2000; Pyo et al.
. We discuss the properties of each fitted component
FWHM
2003b), although Pyo et al. (2003b) show some evidence for sinu・ノットの放出周期や固有運動!
below.
soidal velocity variations (Section 4.1.1). The fitted line width islowest along the jet ridgeline (Fig. 6a), and the region of lowest line
3.4.1 Approaching HVC・ジェットの時間進化にともなう速度の進化!
width corresponds to the region of highest line component intensity
The [Fe ] 1.644 µm emission-line intensity image of the blueshifted
at each position along the outflow axis. This indicates the presenceなどはおもしろそう!
HVC in Fig. 6(a) shows the classic morphology of a well-collimated,of a narrow jet with a relatively undisturbed core.

Multi-epoch Sub-arcsecond [Fe II] Spectroimaging of the
DG Tau Outflows with NIFS. II. On the Nature of the Bipolar
Outflow Asymmetry
M. C. White, G. V. Bicknell, P. J. McGregor and R. Salmeron

前の論文の続き遠ざかる成分のバブル状の構造に焦点を当てている
34

M. C. White et al.

35

Figure 6. Kinetic models of expanding bubbles. The top panels show velocity vectors for snapshots of hemiellipsoidal bubbles with velocity fields
expanding (a) away from the centre of the bubble, and (b) away from thelocation of the central star. Note that the vectors represent the velocity fieldof the dispersed, shocked jet material interior to the bubble walls, whichis the source of the emission; the vectors are placed on the outside of the
bubble for clarity only. The central star, denoted by the star symbol, is at
the same distance from the bubble centre, denoted by the small circle, as the
bubble centre is from the bubble apex. The dashed lines show the alignmentof the velocity vectors in each model. The bubble major axis is aligned to theoutflow axis. The polar angle is defined as the angle from the outflow axis,
measuredat thecentre. Velocities
are modulated
bytwo
the polar
Figure7. bubble
Comparisonof simulated
IFU data ofstatic,angle
radially symmetric hemiellipsoidal bubble models with an internal distribution of emitting material
such thatgas expansion
velocity
to zeromodelat the same
withtheantotal
expandingvelocity
field.reduces(a) Bubble
whereheightgas expansion velocity vectors point away from the centre of the bubble. (b) Bubble model where
above the notional circumstellar disc (grey regions) as the point from which
gas expansion velocity vectors point away from the central star. In both cases, the major axis of the bubble is coincident with the outflow axis, and the bubble
the internal velocity field expands. The star and circumstellar disc are notbase is positioned 0.70 arcsec from the central star. The emitting regions of the bubble have a fixed thickness of 15 au. Maximum gas expansion velocities areincluded in the models.
(c) Velocity profiles of the models. Profiles for the−1
setvelocity
to 230 field
km s−1. Gas expansion
at (panelthe ‘apex’internal
expandingaway fromvelocities
the bubblepeak
centrea) areof the bubble structure, and decrease to 0 km s at the same height above the circumstellar disc
as
the
point
from
which
thebubble
bubbleexpands.
Simulated
are binned into velocity slices based on a model IFU with 0.05 × 0.05 arcsec2 spatial pixels anddrawn in thin lines; profiles
for theexpandingaway
from thedata
central−1
21 km
spectralpixels.
Intensity
scalingeachvelocity slice is different. (c) Channel maps of the receding bubble in DG Tau, as observed in [Fe II]star (panelb) sare drawn
in thicklines.
The dashedlinesbetween
show radial
velocity
1.644
line
emission.
Intensity
scaling
is thesame
betweenprofiles,the µmdottedlines
show verticalvelocity
profilesandthe solidlinesindividual panels.

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical O

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Figure 5. [Fe II] 1.644 µm contour image of the receding outflow fromDG Tau, formed over the velocity range 0–340 km s−1 , for observing
epochs 2005–2009. Contours are drawn at levels of (15, 20, . . . , 40) ×
10−15 erg cm−2 s−1 arcsec−2 . The location of the central star and the position and size of the occulting disc used during the observations are shown
as a yellow star and circle, respectively.

DG Tau outflows – II. Bipolar asymmetry

A Parametric Modeling Approach to Measuring the Gas
Masses of Circumstellar Disks
Jonathan P. Williams and William M. J. Best

星周円盤のガス成分の観測からモデルを介して質量を求めようという試み
以前、麻生さんが紹介したため省略いたしますA spider-like outflow in Barnard 5 - IRS 1: The transition
from a collimated jet to a wide-angle outflow?
Luis A. Zapata, Héctor G. Arce, Erin Brassfield, Aina Palau, Nimesh Patel and Jaime E. Pineda

Barnard 5 - IRS1のSMAによる電波輝線・連続波観測
12CO (2–1)

北東–南西方向の高速双極分子流を分解
赤青それぞれ3つの速度領域で異なる構造(i) 低速では大きく開いたコーン状のローブ(ii) 中速の原始星近くの双極シェル(iii) 高速のバレット
結果として蜘蛛のような形態に星風により分子ガスが巻き込まれた結果だろうOne of the most striking characteristics of the B5-IRS1 outflow is
the different structures observed at different ranges of velocities. As
discussed above, the high-velocity bullets and intermediate-velocity
structure are driven by a collimated (jet-like) wind, while the lowvelocity lobes are consistent with being formed by a wide-angle
wind. An increasing number of outflows have been observed thatshow two components (a wide-angle structure and a collimated feature). They include L1551 (Itoh et al. 2000), HH 46/47 (Velusamy,
Langer & Marsh 2007; Arce et al. 2013), Cepheus HW2 (Torrelles
et al. 2011), and Source I (Matthews et al. 2010; Zapata et al.

Figure 1. (a) SMA integrated intensity (moment 0) colour scale and contour map of the 12 CO(2–1) emission (grey contours), and the continuum emission at
1300 µm (white contours) from B5-IRS1. The blue colour scale represents blueshifted gas, while the red colour scale represents redshifted gas. The integrated
velocity range for the blueshifted side is from −50.9 to +6.2 km s−1 , while for the redshifted side is from +14.7 to +70.8 km s−1 . The white contours range
from 35 to 85 per cent of the peak emission, in steps of 10 per cent. The peak of the continuum emission is 0.032 Jy beam−1 . The grey contours range from 10
to 87 per cent of the peak emission, in steps of 7 per cent. The peak of the 12 CO(2–1) line emission is 150 Jy beam−1 km s−1 . The emission peak here is thearithmetic sum of the total integrated CO emission. The synthesized beam of the continuum image is shown in the lower-left corner. (b) Integrated intensity mapof the B5-IRS1 12 CO(2–1) outflow at different velocity ranges. The different blue and red colour tones of the contours represent different range of velocities
with the lighter (darker) hues representing lower (higher) outflow velocities. The low-velocity blue and red ranges are [−4.3 to +6.2 km s−1 ] and [+14.7 to+24.2 km s−1 ]. The intermediate-velocity ranges are [−39.0 to −15.9 km s−1 ] and [+35.1 to +59.7 km s−1 ], while the high-velocity ranges are [−50.9 to−40.3 km s−1 ] and [+60.2 to +70.8 km s−1 ]. The high-velocity bullets show a larger spread of velocities (see Fig. 4), however, for clarity, we choose theseslightly narrower high-velocity windows. The contours for the different blue and red contours range from about 30 to 90 per cent of the peak emission, in stepsof 5 per cent. In the high-velocity emission, the contours start from 50 per cent the peak of the line emission. The two pink hexagons mark the positions of thestrongest H2 and Hα knots reported by Yu et al. (1999).

velocities that are attributed to the presence of molecular outflows.
This is confirmed by our integrated intensity maps showing a clear
bipolar structure, see Figs 1 and 3.
In Fig. 1(a), we show the velocity-integrated intensity, or zeromoment map, of 12 CO(2–1) for both the blue and red outflow lobes.
The velocity range of integration for the blueshifted lobe is from−50.9 to +6.2 km s−1 , while for the redshifted lobe is from +14.7
to +70.8 km s−1 . With our SMA observations, we detect a subspreadof atvelocities
SMA
data
have levels◦ 1. The
Figure 5. Left-hand panel: PVD perpendicular to the outflow axis in the redshiftedstantiallylobe. Thelarger
PVD wasmadea positionbecause
angle ofthe
157.
contoursignificantly
bettersensitivity
than
previous
studies(see
Velusamy
and the peak emission are the same as Fig. 4. Right-hand panel: the same as in the left-hand panel, but for the blueshifted lobe. The contour levels and the peak
& Langerof1998).
The left-handpanelshownin Fig.in1 Fig.
shows
theThewide-angle
emission are the same as Fig. 4. The blue and red dots represent the position and velocities
the high-velocity
bullets1(b).
dashed curves
bipolaroutflow
that
has
been
reportedinprevious
high-angularresmark the different structures found in the PVDs. Both PVDs were made approximately 15 arcsec away from the continuum source in both sides of theoutflow.
olution studies of this source (Langer et al. 1996; Velusamy &Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on Ju

In Fig. 5, we present PVDs of both outflow lobes, along cuts perpendicular to the outflow axis (PA = 157.◦ 1). These PVDs present
different structures compared to those shown in Fig. 4. At low outflow velocities (!10 km s−1 ), there are two bright curved structuresthat have similar velocity ranges in both outflow lobes, and are
related to the low-velocity and very wide angle structures seen in
Fig. 1(b). In addition, there are two (faint) ring-like structures inboth sides of the flow at intermediate velocities (at outflow velocities of about 10–45 km s−1 ). Such structures, in position–velocity

3.4 Nature of the morphology of the B5-IRS1 outflow

from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on June 11, 2014

L. theA. Zapataal. that is shown in Fig. 1. The
Figure 3698
4. PVD along
axis of theetoutflow◦
PVD was made at a position angle of 67. 1. The black contours range from 6
to 90 per cent of the peak emission, in steps of 4 per cent. The peak of the line
emission is 6 Jy beam−1 . The grey contours represent the negative emission
arising mainly from velocities close to the cloud’s systemic velocity andhave the same value as the black contours, but negative. The curves markthe different structures found in our PVD. The circles mark the highest
velocity features in the outflow. Velocity is shown with respect to the LSR.

the observed intermediate-velocity features are caused by propagating bow shocks in a jet, or highly collimated wind. The high-velocity
bullets are located at the tips of the intermediate-velocity ellipticalstructure in the PVDs shown in Fig. 5 (and at the far ends of the PVD
shown in Fig. 4). Hence, it seems possible that these high-velocity
features are tracing the tip (or head) of the bow shock responsible
for the intermediate-velocity cavities.