THE STAR FORMATIONNEWSLETTER
No. 278 (#18-22)
塚越崇(茨城大)11 Mar. 2016

1. On the universality of interstellar filaments: theory meets
simulations and observations
2. The SPHERE view of the planet-forming disk around HD100546
3. On the 2015 outburst of the EXor variable star V1118 Ori
4. A new insight into the V1184 Tau variability
5. Hot ammonia around young O-type stars.
III. High-mass star formation and hot core activity in W51 Main

On the universality of interstellar filaments: theory meets simulations and observations
Christoph Federrath
Universality of interstellar filaments

柱密度• hydrodynamical simulationで分子雲フィラメント形成を計算• 観測の~0.1pcを説明できるか?
• gravity, turbulence, magnetic field, jet/

• 全計算でフィラメントを生成• DisPerSEでフィラメント同定• 観測でも使用されている• 面密度はΣ∝r-1でよく再現

Downloaded from http://mnras.oxfordjournals.org/ at Ibaraki University on Fe

outflow feedbackのoff/on

L=2pc, M=388M⦿

379

•

•

•

proportional to the square of the Mach number, M2 = σv2 /cs2 , theanalysis box scale L
turbulent density fluctuations will quickly vanish
near
and
below
the
各フィラメントでの幅比較
a magnetic field), w
sonic scale. This means that the incompressible (no density fluctuscale is around λsonications) Kolmogorov theory then provides a good approximation ofagreement with their
重力のみでは細め。乱流を入the turbulence below the sonic scale, with a power-law exponent ofthough the prediction
arecontrolledbysupersonic
α
∼
1/3,
while
scaleslargerthan
λ
sonic
れると0.1pcに。
in Smith et al. (2014turbulence with an exponent α ∼ 1/2. We clearly see that the sonic(4) strictly speaking
scale
is
a
fundamental
characteristic
scale
on
which
the
turbulence• 重力はdense partの形成のみに影響
Smith et al. (2014) inchanges behaviour from being highly compressible and supersonicin their simulations,with α ∼ 1/2 for ℓ > λsonic , to subsonic, nearly incompressible withover a limited range
α ∼ 1/3 for ℓ < λsonic . This transition may in fact be the same as
speed
et al.Fie2
Figure 3. Filament width as a function of column density in three of our simulations: Gravity only
(squares),(Glover
Turbulence + Magnetic
sonic
scale
λによる幅の説明
what
has beentermed ‘thetransitiontocoherence’inobservations
Turbulence + Magnetic Fields + Gravity (circles). The horizontal dotted lines show the average filament width. Gravity-only models proofsonic ofscales
ratherUniversality
interstellarfiwhile turbulence models produce widths consistent with observations. Turbulence alone produces
relatively low column
density filamentof dense cores (Goodman
et
al.
1998;
Pinedaetal.
2010).
leads to denser filaments, similar to the difference between the quiescent Polaris and the star-forming
Aquila cloudfilament
shown in fig. 7 inwid
Arzo
possible=> 今回の計算ではλ=0.05Based on this, the sonic scale is explicitly defined as [seesee in their simulatio
equation (22) in Federrath & Klessen 2012],
0.15pcとなりconsistent
!
$2
"
#
cs
1/2
1 + β −1,(4)λsonic = L
σvwhere L, σ v , cs and β are the cloud scale, the velocity dispersion on
the cloud scale, the sound speed, and the ratio of thermal to magnetic
フィラメント内で亜音速、外で
pressure, plasma β = pthermal /pmagnetic . Note that equation (4) only
takes magnetic pressure into account, while magnetic tension isは超音速となるignored, which would require a (so far uncertain) correction for

2 β → ∞ is the limit in

3 Whether equation (4)recently studied in Butl
Tremblay et al. (2015),
cosmic gas such as thesonic scale and thus in

Figure 5. Filament velocity dispersion profiles along the line of sight (LOS), σ vz (solid line), and perpendicular to the LOS, σ vy (dottedline)
in ourthat includes
gravity,
turbulence,magnetic
fields (SFE)and jet/outflow
feedback(model
6 ingravity,
Table 1).
The velocity
disp
Figure4. simulation
Filament width
as a functionof thestar formation
efficiency
in our simulation
that
includesturbulence,magnetito feedback
the sound(model
speed,6such
that
these
curves
show
the
Mach
numberinsideandoutsidethefilamentsasafunctionofradiuscentredonin Table 1), shown as diamonds. The hatched region shows the observed range of filament widths and column densit

The SPHERE view of the planet-forming disk around HD100546
A. Garufi et al.A&A proofs: manuscript no. aa(d)(b) Qφ FastPol

Uφ SlowPol• HD100546• B9V, d=97pc• 2 companion candidates
0.5″• SPHERE/VLT

(c)

0.5″
(e)

Qφ SlowPolb

Northernarm• ZIMPOL•

highest resolution at
opt. (0.02”)

AO artifact
SpiralDark wedge• IRDIS

0.5″

Bright
wedge

FA
R

c

NE

AR

Qφ SlowPol (unsharp masking)

0.1″

(b)
A&A proofs:
manuscript no. aa

Major axis (SE)
Major axis (NW)
Near side
Far sideSurface brightness (arbitrary units)

200

100

70
50

SPHERE R' band
NACO H band
NACO Ks band

1

Normalized surface brightness

(a)
300

IRDIS image

large scale (H)

0.7
0.5

たくさんのarm

0.3

Northernspirals

0.2←HD100546c
Discontinuity
Spiral arm

30
0.10

0.15

Angularresolution

0.1
0.20

0.25

0.10

0.15

0.20

0.25

0.30

Radius ('')

Radius ('')
(c)

Wing
structure(d)1

Surface brightness (arbitrary units)

300

250

200

Medium1″

150

Far side

100

Major axis

Backward scattered photons

050

Near side

Major axis

Forward scattered photons

100

150

200

250

Westernspirals?

SPHERE R' band
NACO H band
NACO Ks band
SPHERE R' band (UΦ )0.9Normalized surface brightness

1″

Garufi et al. 2016: The SPHERE view on HD100546

0.8Southernspirals

0.7
0.6
0.5
0.4
0.3

⇽ HD100546b

0.2

Composite

Large
1″

0.1

1″

1
0
-1

300

350

0

Azimuthal angle (degree)

50

100

150

200

250

300

350

Azimuthal angle (degree)

Fig. 2. ZIMPOL polarized light brightness profiles. (a): Radial profile in the inner 0.300 . The location of the abrupt decrement along the near sideaxis and the spiral arm bump along the far side axis are highlighted by the horizontal dashed lines. (b): Radial profile along the SE axis comparedto the same from the NACO dataset. Each profile is normalized to the respective peak intensity. The predicted radial location of HD100546c isindicated by the vertical brown area. The angular resolution of the three datasets is shown by the horizontal lines. (c): Projected azimuthal profileof the inner rim (at r = 0.1500 for i = 42 ). Angles from North to East. The emission from the four di↵erent quadrants is shown with di↵erent
colors. (d): Projected azimuthal profile at the predicted radial location of HD100546b (see red ellipse in Fig. 1c) compared with the same from the
NACO dataset. Each profile is normalized to the respective peak intensity. The azimuthal location of b is indicated by the vertical brown region.
The vertical dashed lines define the four quadrants as in (c). In the lower space, the same normalized brightness distribution from the SPHERE U
image is shown (note the di↵erent y-axis scale therein). All errors are given at 3 level as described in Sect. 3.1.

• ZIMPOL• 北西-南東方向にinner cavity

Enhancedbrightnessarm(K band only)• 赤外(NACOデータ)と位置のずれがある

An unsharp masked version of the SPHERE Q image wasobtained by subtracting the smoothed version (by ⇠ 10⇥FWHM)
from the original image. Since the noise in the science image ishigh, we firstly subtracted the variance over a FWHM from theoriginal image. This operation had the e↵ect of decreasing thesignal of the inner regions so as to avoid over-subtraction fromthe smoothed image, and resulted in a much sharper image ofthe inner 0.500 . Finally, the resulting image has been smoothedby one FWHM. The morphology of the final image strongly
depends on the choice of the mentioned parameters. However,some features are visible for a broad range of parameters. As

to an imperfect data reduction as the equalization of polarization states (minimizing instrumental polarization, see Sect. 2.1)
cannot eliminate such a signal and the crosstalk e↵ect between
Stokes parameters (see Avenhaus et al. 2014b) is a minor effect in SPHERE. Furthermore, the spatial distribution of this signal is consistent through instruments and wavebands (see the U
NACO images by Avenhaus et al. 2014a), with a positive signal
being di↵usely detected to the East and to the West, and with anegative signal to the NW and SW. Both the bright and the darkwedge highlighted in Fig. 1c seem to have a counterpart in the
SPHERE U image. However, a ⇠ 20 o↵set exists in the az-• maskを強めるとspiral armが見える• IRDISKnot(mainly H band)arm

small scale (composite H&K)
RDIS imagery of HD100546. Images in the top row are in the H2H3 band, in the bottom row in the K1K2 band. The first three• 大小のスケールでarmが見える• 色々なデータと比較•

散乱光では内側のrimまで拾えている • ZIMPOLのarmはb-sourceに向かって伸びている•

Garufi et al. 2016: The SPHERE
view
on HD100546
広がった放射は波長によらず見える
scalearmsの位置は2010-2015で変化なし• large

(a)

(b)

(c)

(d)

0.2″

0.2″

0.2″

H band: polarized light vs intensity

R′ polarized light vs M′ intensity

R′ polarized light vs K vs M′ intensity

0.2″
K band intensity: 2015 vs 2010

Fig. 4. Multi-wavelength and -epoch images of HD100546. (a): NACO polarized light H band (colors) and IRDIS H band (contours). (b): ZIMPOL
polarized light R0 band (colors) and NACO M0 band (contours). (c): ZIMPOL polarized light R0 band after unsharp masking (colors), IRDIS K
band (yellow contours), and NACO M0 band (cyan contours). The inset image is the NACO polarized light KS band image from 2006 after unsharp
masking. (d): IRDIS K band from May 2015 (colors) and NICI K band from March 2010 (contours).

•

observed in the visible (colors) are also evident from the IR 4.1), on the global disk morphology (Sect. 4.2), on the disk fea円盤構造等を過去の観測と比較して議論
NACO data (contours). Particularly good is the match between tures (Sect. 4.3), and on the environment around planet b (Sect.
the bright/dark edge to East. Also worth mentioning is the tenta- 4.4).
tive detectionof the spiral at 0.200 (Fig. 1e) from the M0 dataset.
惑星候補天体(b-source)位置でK-bandでのみ広がった放射がある
In Fig. 4c, we show the spatial connection between the northern arm spotted from the ZIMPOL image after unsharp masking 4.1. Cavity and inner rim
H,theK bandが散乱するτ(高さ)でちょうど円盤構造が変わっている(colors)di↵use emission from the M0 band (cyan con- Our results on the disk cavity (from Sect. 3.1.1) support the pre• andtours) surrounding planet b (Quanz et al. 2015). The same emis- vious findings in DPI (Quanz et al. 2011; Avenhaus et al. 2014a),sion from heating
the IRDISsourceが円盤に埋もれていて視線方向に減光されている
K band is shown (yellow contours). The that are that the disk around HD100546 is truncated at roughly•
arm seems to head to the location of b and a blurred counterpart 15 AU and that this cavity is consistent with being intrinsically•0On the 2015 outburst of the EXor variable star V1118 Ori
T. Giannini et al.

• V1181Ori• EXor型増光天体• 過去に4回増光(いずれも数年ほど持続)–1• データ: EXORCISM project• JHK photometry; 1.1m AZT-24望遠鏡Fig. 1.— V1118 Ori JHK lightcurve (blue, green and red, respectively).
10−6• 増光前後の比較• 質量降着率の増大

dMacc/dt (M ∀/yr)• optical spectra; 8.4m LBT, 3.6m TNG
10

10

−7

−8

• 質量損失率の増大([OI],[SII]など)
HeI447HγHδ

HeI403

CaIIK

H8

H9

H10

H11

H12

H13

H14

−9

H15

• 過去のバーストよりも弱い

10• もしくは増光/減光途中かもFig. 3.— Mass accretion rate determination coA new insight into the V1184 Tau variability
T. Giannini et al.

T. Giannini et al.: A new insight into the variability of V1184 Tau• V1184Tau (d=1.5kpc)• G5 star, 7-39L⦿
• UXor型の変光と言われはじめた• データ: EXORCISM project

Fig. 1. V1184 Tau optical and near-IR light curves vs. Modified Julian Date (MJD). Red points are new infrared owhile black points are literature data. Vertical black lines indicate the periods when mid-infrared data have been re• + published photometry data
corresponding telescope is also labeled. In the second panel we show, as an example, the fit through the V light curve. Th• EXorよりUXor的な増光の証拠

and minimum values of the fit are T.
assumedto be V magnitude of the low- and high-state of the source, and used to c
Giannini et al.: A new insight into the variability of V1184 Taucorresponding SED (see Section 3.2).

• 強度変化が緩やか• 強い赤外超過• 現在増光のピーク

photometric data, we plot the locus of the main-sequence and the plot indicates ∆AV ∼ 2 mag, while the V light c
two extinction vectors (Rieke & Lebofsky, 1985 and Cardelli et fading of more than 5 mag (Figure 1).
al., 1989), which were applied at the location of the colors of aG5-type star.
In the right panel we show the near-infrared J
Both plots indicate that the source is bluer when brighter, a two-color diagram. Two main conclusions can be dtypical behavior of all the accretion- and extinction-driven fluc- data points do not align with either of two depictetuations (Lorenzetti et al. 2012).
curves but rather exhibit a clear shift to the right th
To go into more detail: the optical two-color plot shows that with the decrease in the source brightness; 2) if ondata with V ≤ 17 mag (blue dots) and with 17 mag ≤ V ≤ 18 tinction, their spread should imply a ∆AV larger tmag (red filled dots) tend to align along the extinction vector incompatible with the value derived in the opticalof Cardelli et al. (1989), and cluster at AV ∼ 3 mag and AV ∼ Instead, the points indicate the presence of a stron5 mag, respectively. Conversely, data points corresponding to that we quantify by plotting the locus of the combin
Fig.
Tau redopticallight
Modified thermal
Julian Date
(MJD). Red
pointsare new
infrared ocomponents,
one atthe stellar
photosphericV> 1.
18V1184
mag (opendots)anddo near-IR
not follow
anycurvesof thevs.
extinction
while
black
pointsare
literature
data.
Vertical
black
lines
indicatetheperiods
when
mid-infrared
data
have
been
re
rang
vectors. This eﬀect, already noted in the V vs. V − I and V vs. of 5700 K (that of a G5 star) and the second
corresponding
telescopeisalsolabeled.
Inthesecond
panel
we
show,
as
an
example,thefitthroughtheVlightcurve.
Th
V − R color-magnitude diagrams by Barsunova et al. (2006) and 1000 and 2000 K, and for diﬀerent ratios between
and minimum
values of
fit areto be V magnitude
lowsource,emission
and used istolac
areas.
In andthe high-statehigh state,ofthetheinfrared
Semkov
et al. (2008),
is thelikelydueassumed
to the scatteringof the stel- of thecorresponding
SED(see
Section3.2).
lar light by small dust particles, which tends to prevail against nated by the stellar emission, while in the faint s
the direct light from the star when this latter becomes fainter. counted for by a thermal source between 1000 K
This eﬀect makes the optical color-color plot eﬀectively unable with an eﬀective emitting area > 104 larger than th• 距離は1.5kpcでなく200pc• SED fittingより

C. Goddi et al.: Hot NH3 around Young O
Hot ammonia around young O-type stars. III. High-mass star formation and hotcore activity in
W51 Main
and (
and (
Ciriaco Goddi et al.

W51 Main

NW• W51: 大質量星形成領域

WE• 2x107L⦿, 7000M⦿
• Main(UC,HC HII region) (, IRS1, IRS2)• JVLA• NH3(J,K=6~13), CH3OH,CH3CN• ~<0.3”• W51e2(substructure), W51e8の
構造に着目

ever,
respe
sense
nent,
symm
are m
they c
comp
of th
the h
lite li
Whil
the s
W (i.
with
of rot
an ou
B
inver
most
The c
tion o
maps
evide
ening
also d3.1.2
color: 25GHz cont.

Figur
contour: NH3(6,6) (blue:emm., red:abs.) sionline-wout, w(disp• 吸収はHC HIIコア

W51e2+ionized flow

• 高励起に行くほど速度

A&A proo

勾配は東西になる• 低い励起ではionized flowのコンタミがある• 東西方向がリアルなコアの速度勾配 (回転)W51-W

• W51e8についても同様の解析=>EW方向、NS方向にわずかに速度勾配

A&A proofs: manuscript no. w51e2_nh3_arxiv

Fig. 5. Total intensity images of 5 inversion transitions of NH3 as well as the CH3 OH line towards W51e2. Light colours indicate emission
and dark colours absorption. Contours of the NH3 total intensity are also overplotted in the image: emission is displayed with black contours,
representing factors 1, 2, 3, 4,.... of 20 mJy beam 1 km s 1 ; absorption is displayed with white contours, representing factors 1, 5, 9, 13,....
of –50 mJy beam 1 km s 1 ), for each transition, respectively. The images were integrated over the velocity range covering the main hyperfine
component for each transition (from 48 km s 1 to 70 km s 1 ). The velocity resolution was smoothed to 0.4 km s 1 for all transitions. The synthesized
beams (0.00 19–0.00 26) are shown in the lower left corner of each panel (see Table 1). The images were constructed with 0.00 04 pixels for all transitions.
The cyan crosses mark the positions of the 0.8mm continuum peaks identified by Shi et al. (2010a) and corresponding to sources W51e2-E, W51e2W, and W51e2-NW (from left to right), claimed to be individual high-mass YSOs.

3.2. Spectral profiles

We extracted spectral profiles towards individual sources bymapping each spectral channel and summing the flux density ineach channel map for each transition. Figures 11, 12, and 13
show such spectral profiles extracted towards the W51e2-W HC
HII region, the sources seen in NH3 emission in the W51e2 core
and the W51e8 core, respectively. As displayed in these plots,
the NH3 profiles are not simply composed of a single spectral
component. In fact, owing to the interaction with the quadrupole
moment of the nitrogen nucleus, each NH3 inversion line is actually split into five components, a “main component” and foursymmetrically spaced “satellites”, which make up the NH3 hyperfine structure (HFS). The frequency separations and relative

the middle and right panels show in more detail the profiles of theNH3 main central hyperfine components and the CH3 OH line, respectively. Individual spectral profiles do not show the presenceof multiple velocity components, so they can be reasonably wellfitted by single Gaussian profiles.
In order to derive the line parameters, we used fivecomponent Gaussian models to fit the HFS in each of those inversion transitions, where we fixed their velocity separations according to the calculated values, assuming the same line-widths
for all four hyperfine components, to reduce the number of freeparameters. In order to derive the kinematics in optically thinlines, we also fit the hyperfines only with a 4-component gaussian model again with fixed o↵sets and linewidths. The fitted line
parameters are reported in Table 2 for the main components and• スペクトルフィットによりτ, Trot, Ncolを導出• W51e2 complexについて• W51e2-E

hot molecular gasが付随、molecular outflowを駆動、ジェットによるH2O maserが付随、低速のOH
expanding shell=> high mass protostar• W51e2-W
Fig. 14. Overlay of the CO J=3-2 0 moment image from the SMA
W51e2のactive star forming regionだと思われていたが実はEのほうがactiveだった。低th(blue and red contours; Shi et al. 2010b) onto the total intensity (0thmoment) map of the (6,6) inversion transition of NH3 (gray scale andwhite+black contours). The NH3 emission is displayed with white contours, representing 30% to 90% with steps of 20% of the line peakfor the (6,6) line, 107 mJy beam 1 km s 1 ). The NH3 absorption isdisplayed with black contours, representing
factors 1, 5, 9,.... of –⦿
50 mJy beam 1 , for all transitions. The CO 3-2 line contours are from
10% to 90% with steps of 10% of the line peak (67 Jy beam 1 km s 1 forthe red contours and 82 Jy beam 1 km s 1 for the blue contours). Blueand red contours correspond to blue- and redshifted gas, with integrated
velocity ranging from –124 to –12 km s 1 and from +10 to +116 km
s 1 , respectively. The cyan contour locates the peak of the total intensity map of the most highly-excited NH3 transitions (13,13): we assume
that this peak pinpoints the high-mass YSO driving the CO outflow.

速のOH expanding shellがあり、UC HII regionによるものと思われる。回転は見えるがケプラーかどうかは不明。中心YSO質量20M⦿に対しNH3ガス質量5M=> 降着フェーズが終わったO型星• W51e2-NW

ホットガスが付随、コンパクトアウトフローが付随=> (high mass) YSO forming region

0.00 06, 320 AU). This material could be potentially accreting onto
the central YSO, either in the form of an accretion disk or an infalling envelope, provided that it has not been stopped by theintense radiation pressure from the central O-type star and thethermal pressure from the ionized gas in the HII region. Indeed,
the observed velocity gradient provides evidence for rotation in
the molecular core surrounding the HC HII (§ 3.1.1).

• W51e8 => massive hot core (70M⦿、T~200K)