星形成ゼミ
2017/11/10
SFN #297 #6-#10
長谷川哲夫（国立天文台）

1

iv:1708.02961v1 [astro-ph.SR] 9 Aug 2017

1

Department of Space, Earth and Environment, Chalmers University of Technology, Onsala Space Observatory,
Observatorievägen 90, 43992 Onsala, Sweden; e-mail: daria.dallolio@chalmers.se2
INAF–Osservatorio Astronomico di Cagliari, Via della Scienza 5, 09047 Selargius, Italy
297
Methanol masers
revealthe17,
magnetic
fieldGermanyof the high-mass protostar IRAS
3
Max-Planck-Institutefor Astronomy,
Königstuhl
69117 Heidelberg,
4
JodrellBank Centre for Astrophysics, Department of Physics and Astronomy, University of Manchester, M139PL Manchester, UK#618089-1732
D. Dall’Olio, W. H. T. Vlemmings, G. Surcis, H. Beuther, B. Lankhaar1, M. V. Persson, A. M. S. Richards, and E. Varenius
ReceivedarXiv:1708.02961v1
2 June 2017; accepted 9 August 2017
ABSTRACT
Context. The importance of the magnetic field in high-mass-star formation is not yet fully clear and there are still many open questions

concerning its role in the accretion processes and generation of jets and outflows. In the past few years, masers have been successfully
used to probe the magnetic field morphology and strength at scales of a few au around massive protostars, by measuring linear
polarisation angles and Zeeman splitting. The massive protostar IRAS 18089-1732 is a well studied high-mass-star forming region,
showing a hot core chemistry and a disc-outflow system. Previous SMA observations of polarised dust revealed an ordered magnetic
field oriented around the disc of IRAS 18089-1732.
Aims. We want to determine the magnetic field in the dense region probed by 6.7 GHz methanol maser observations and compare
it with observations in dust continuum polarisation, to investigate how the magnetic field in the compact maser region relates to thelarge-scale field around massive protostars.
Methods. We reduced MERLIN observations at 6.7 GHz of IRAS 18089-1732 and we analysed the polarised emission by methanol
masers.
Results. Our MERLIN observations show that the magnetic field in the 6.7 GHz methanol maser region is consistent with themagnetic field constrained by the SMA dust polarisation observations. A tentative detection of circularly polarised line emission is
also presented.
Conclusions. We found that the magnetic field in the maser region has the same orientation as in the disk. Thus the large-scale fieldcomponent, even at the au scale of the masers, dominates over any small-scale field fluctuations. We obtained, from the circularpolarisation tentative detection, a field strength along the line of sight of 5.5 mG which appeared to be consistent with the previous
estimates.
Key words. magnetic field – stars: formation – stars: massive – masers – polarization

大質量星形成領域 IRAS 18089-1732 (L = 12.89d, B = +0.50d)
観測• 距離 2.34 kpc, L = 1.3E4 L_sun, M_gas ~ 1E3 M_sun
6.7 GHz
CH3OH
maser
1. Introduction
netic braking which removes
mostof theangular momentum• Disk + ouMlow（南北方向）
MERLIN
from the circumstellar gas. This is the so-called magnetic brakSMAでの観測（Beuther+
2008年3月,
4月, 7月The roleof magnetic fields 2010）
during the formation of high-mass ing catastrophe in disc formation(Mestel& Spitzer 1956; Allen•
Diskのdust
conTnuum偏波→diskに沿った磁場（Chandrasekhar-Fermi法でB_pos
~
11
mG）
Beam
0.18”
x
0.03”
stars is not yet fully understood. As in the case of low-mass star et al. 2003; Mellon & Li 2008; Li et al.
2011, 2013). Only by• OuMlowのCO3-2でGoldreich-Kylaﬁs効果（P~ 8%） field apformation,
simulations
have shown that the magnetic
inserting non-ideal MHD e↵ects such as Ohmic dissipation or
CH3OH
6.7 GHzmaserのモニター（Goedhart+2009）→29.5日周期の強度変動
pears toprevent
fragmentation
around massive
protostars (Pe-ters et al. 2011, Myers et al. 2013) to influence accretion and todrive feedback phenomena such as collimated outflows and jets
(Seifried et al. 2011, 2012). Both the core accretion model (e.g.

ambipolar di↵usion, or the combined action of both, is it possible to overcome the apparent conflict between observations andsimulations (Machida et al. 2014, Zhao et al. 2016). Moreover,Article number, page 6 of 12

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A&A proofs: manuscript no. IRAS15

Fig. 1. 6.7 GHz methanol maser spectra of the three epochs, obtained by summing all the pixels in the image for each channel.

Maser groups• Blue group: YSO• Red group: ouMlow

Fig. 2. Spectra of the brightest maser feature F.01 in the three epochs.

Fig. 3. Masers identified in March (left), April (centre) and July (right) as listed in Table 3–5. The bottom panel shows a zoom of the region marked by the dashed grey boxes in the top panels. Eachmaser is represented by a triangle. The di↵erent sizes of the triangles represent the intensity, while the colours indicate the velocity of the maser feature, according to the scale reported in the colour
bar. Line segments mark the direction of the polarisation angle for the maser features that show linear polarisation. The average direction of the resulting magnetic field B obtained for two groups
D. bottom
Dall’Olioet corners
al.: Magnetic
in IRAS18089
of masers as defined in Sect. 5.2 is indicated in therightof eachfield
panel.The
Julyobservations were in dual circular polarisation only.

maser features. Each maser is represented by a triangle. The
di↵erent sizes of the triangles represent the intensity, while thecolours indicate the velocity of the maser feature, following thescale reported in the colour bar. For March and April only, theline segments mark the direction of the polarisation angle forthe maser features that show linear polarisation. Under the assumption that the angle between the magnetic field and the lineof sight is ✓ > ✓crit ⇠ 55 , where ✓crit is the Van Vleck angle(e.g. Surcis et al. 2011), we considered the linear polarisation
perpendicular to the magnetic field. The vectors are also scaled
logarithmically according to Pl as reported in Tables 3 and 4. We
also plot, in the bottom right corners, the average direction of theresulting magnetic field B obtained for two groups of masers as
defined in Sect. 5.2. For the F.06 feature we tentatively detect acircular polarisation signature and the magnetic field along theline of sight (see Sect. 5.4).

5. Discussion
5.1. Maser distribution and kinematics

Goedhart et al. (2009) presented a spectrum of the 6.7 GHz
methanol maser emission of IRAS 18089. The spectrum is divided in two distinct blocks around two main peaks: one locatedat a velocity of 33.7 km s 1 and another at 39.2 km s 1 . In ourobservations, we found a similar spectrum, with all our maser
features grouped around those two values (see Tables 3–5): asshown in the Tables, the velocity shifts are within the channel
width, apart from F.06 which presents a shift of ⇠0.2 km s 1 ,probably due to a blend of two components (see Sect. 5.4). Since
the same velocities occur in all three of our epochs, we divided
the masers in two groups: a blue group spanning a velocity range
from 30.0 to 36.4 km s 1 (containing F.04, F.05, F.06, F.07, F.08,
F.09, and F.11), and a red group from 37.7 to 39.2 km s 1 (containing F.01, F.02, F.03, and F.10). The blue group presents avelocity that is similar to the vlsr of the compact core (⇠33.8

Article number, page 4 of 12

Zeeman - TentaTve detecTon (再現せず）
A&A proofs: manuscript no. IRAS15

Blue group

March F.06

April F.06

B_los = 5.5 mG
Fig. 4. Masers in the blue group (blue triangles and blue segments) superimposed on the integrated I image of the dust continuum emission
observed by Beuther et al. (2010) at 880 µm with SMA (red contours; the contours are drawn in 10 steps). The magenta line segments show
the magnetic field orientation obtained by linearly polarised dust emission (Beuther et al. 2010). The blue segments represent the magnetic fieldorientation obtained by our linearly polarised methanol maser emission (Sect. 5.2); therefore the magnetic field follows the same direction indicatedby the dust emission. The red and blue ellipses show the beams of SMA (1.6500 ⇥ 1.0500 , position angle 51 ) and MERLIN, respectively. Left panel:
March; right panel: April.

Fig. 6. Fit to maser feature F.06 assuming the presence of two hyperfine components; left: March observation, right: April observation. Upper
panels: total intensity I spectrum plotted as the black histogram. The solid red curve in the upper panel is the best fit using two Gaussian components, representing two maser hyperfine components separated by 0.2 km s 1 (see Sect. 5.4). Component 1 is the blue long-dashed line andcomponent 2 is the green short-dashed line. Middle panels: circular polarisation V spectrum (black histogram). The solid magenta line is the sumof the derivatives of the two components of the fit from the upper panel. The blue long-dashed line is the derivative of component 1 and the green
short-dashed line is the derivative of component 2. Lower panels: we plot the residuals between the circular polarisation (middle panel, black
histogram) and the sum of the derivatives (middle panel, solid magenta line).

of the maser line can also occur in hyperfine interactions or in
3D maser propagation e↵ect. Therefore, as suggested by Surcis
et al. (2015), a detection of circular polarisation should only beconsidered real if it presents a V peak flux density at least fivetimes higher than the RMS.

in red, component 1 is the blue long-dashed line and component
2 is the green short-dashed line. We summarise the best fit parameters in Table 9.
It is possible to see that the I spectrum in April presents areversed profile with respect to March. A similar behaviour is1

Guido De Marchi1 , Nino Panagia2,3,4, and Giacomo Beccari5

European Space Research and Technology Centre, Keplerlaan 1, 2200 AG Noordwijk, Netherlands; gdemarchi@esa.int

2
297 Photometric
Determinationof SantheMartin
Massofpanagia@stsci.edu
Pre-main-sequence
Space Telescope
Science Institute, 3700
Drive,Accretion
Baltimore, MD Rates
21218, USA;3
INAF-NA, Osservatorio Astronomico di Capodimonte, Salita Moiariello 16, I-80131 Napoli, Italy#7 Stars. V. 4RecentStar Formationin theRd.,
30Virgin
DorGorda
NebulaSupernova Limited,
OYV #131, Northsound
VG1150, Virgin Islands, UK

Guido 5De
Marchi,Southern
Nino Panagia,
and Giacomo
Beccari
EuropeanObservatory,Karl-Schwarzschild-Str.
2, D-85748 Garching, Germany; gbeccari@eso.org
ApJ 846, 110(2017)2017 May 1; revised 2017 July 27; accepted 2017 August 9; published 2017 September 7
Received

Abstract
We report on the properties of the low-mass stars that recently formed in the central ~2.7
¢ ´ 2.7¢ of 30 Dor,
including the R136 cluster. Using the photometric catalog of De Marchi et al., based on observations with the
Hubble Space Telescope, and the most recent extinction law for this ﬁeld, we identify 1035 bona ﬁde pre-mainsequence (PMS) stars showing Ha excess emission at the 4s level with an Ha equivalent width of 20 Å or more.
We ﬁnd a wide spread in age spanning the range ~0.1–50 Myr . We also ﬁnd that the older PMS objects are placedin front of the R136 cluster and are separated from it by a conspicuous amount of absorbing material, indicatingthat star formation has proceeded from the periphery into the interior of the region. We derive physical parameters
for all PMS stars, including masses m, ages t, and mass accretion rates Ṁacc . To identify reliable correlations
between these parameters, which are intertwined, we use a multivariate linear regression ﬁt of the typelog Ṁacc = a ´ log t + b ´ log m + c . The values of a and b for 30 Dor are compatible with those found in
NGC 346 and NGC 602. We extend the ﬁt to a uniform sample of 1307 PMS stars with 0.5 < m /M:<1.5 andt < 16 Myr in six star-forming regions in the Large and Small Magellanic Clouds and Milky Way withmetallicities in the range of 0.1–1.0 Z:. We ﬁnd a = -0.59 0.02 and b = 0.78 0.08. The residuals are
systematically different between the six regions and reveal a strong correlation with metallicity Z, of the typec = (-3.69 0.02) - (0.30 0.04) ´ log Z Z:. A possible interpretation of this trend is that when themetallicity is higher so is the radiation pressure, and this limits the accretion process, in both its rate and duration.
Key words: galaxies: star clusters: general – galaxies: star clusters: individual (30 Dor) – galaxies: stellar content –Magellanic Clouds – stars: formation – stars: pre-main sequence

HST mulTband photometric catalog (De Marchi+ 2011)
1. Introduction
ages in excess of ∼10 Myr and are preferentially located to the• Concentrate on F555W (V), F814W (I), and F656N
east of(H-alpha)R136. This bands
indicates that formation of low-mass stars in

The Tarantula Nebula (30 Dor) in the Large Magellanic
Cloud (LMC) is the nearest extragalactic region of extended
star formation. It is energized by the OB association
NGC 2070, which hosts at its center the Radcliffe 136 (R136)
cluster, the closest example of a massive extragalactic starburst

this ﬁeld has proceeded over a period of time longer than thatassociated with the birth of the massive stars of R136 some
~2–3 Myr ago (e.g., De Marchi et al. 1993; Hunter et al. 1995;
de Koter et al. 1998; Massey & Hunter 1998; Walborn et al.

The Astrophysical Journal, 846:110 (16pp), 2017 September 10

The Astrophysical Journal, 846:110 (16pp), 2017 September 10

Figure 2. Identiﬁcation of objects with excess Ha emission. The dashed line is
the median photospheric V - Ha color for stars with small (<0.05 mag)combined photometric uncertainties in all three bands. For comparison, thedotted line shows the colors in these ﬁlters for the model atmospheres ofBessell et al. (1998). Thick ﬁlled circles denote stars with Ha excess emission
higher than the 3s level. The AV=1.8 reddening vector is shown, as an
example, for the 30 Dor extinction law.
The Astrophysical Journal, 846:110 (16pp), 2017 September 10

reddening caused by the circumstellar disk. Therefore, theeffect is negligible.
4. Searching for PMS Stars
As mentioned above, differential reddening makes it
impossible to identify bona ﬁde PMS stars based on theirbroadband photometry alone. Instead, we search for stars with astrong Ha excess emission. This feature is characteristic of
low-mass star formation and is attributed to the gravitationalenergy released by infalling matter that ionizes and excites thesurrounding gas (e.g., Königl 1991; Shu et al. 1994). A reliable
method to detect these objects using multicolor photometry isaddressed in a series of papers by De Marchi et al. (2010,2011a, 2011b, 2011c, 2013a), to which we refer the reader.
Brieﬂy, the median V−I and V - Ha colors of stars withsmall (<0.05 mag) photometric uncertainties in the V , I , andHa bands deﬁne an empirical reference template against which
Figure4. Radial objects
distributionwithof UMSstars (solidline) and Evenof low-mass
Ha blueemission.
beforewe identify
excess
PMS candidates younger than 5 Myr (short-dashed green line) and older thanV -agesHa,
correction
forreddening,thecolor–color
diagram
5 Myr (long-dashed red line). The dotted line corresponds to stars withexcess,V−I provides
identiﬁcation
ofnormalizedstars withintermediate
betweena5 robust
and 10 Myr.
The curves are
by H
theanumberofobjects
each bands
sample, theand reddening
the solid linevector
is further
rescaled
by aparallelfactor of tosince
in intheseruns
almost0.65 for ease of comparison. Note that crowding is the cause of the apparent
themedianphotosphericcolorsofnormalstars(DeMarchidrop observed at small radii in the proﬁles of PMS stars (the effect is negligible
et theal. much2010).
forbrighter and sparser UMS objects). Error bars correspond to the
Following
thisprocedure,
we identiﬁed stars with V - Ha
Poissonuncertainty
on thenumber counts.
color exceeding the V - Ha color of normal stars with noexcessfor
andcomparison
with the sameby at ofleast
0.1between
mag or byshows
the V−I
distribution
stars
thethree5
a color(i.e.,
times
uncertainty
on shown
the V and10 theMyrphotometricisochrones (thelatter is not
in H
Figure
3 but

De Marchi, Panagia, & Beccari

Figure 3. Same as Figure 1, but including all candidate PMS stars with Ha
excess emission and the isochrone for PMS age of 5 Myr (thin dashed line) forAV=1.22. Stars with Ha excess emission located above and below the 5 Myr
isochrone are indicated with green thick ﬁlled circles and red crosses,
respectively, and are approximately younger and older than 5 Myr.
De Marchi, Panagia, & Beccari

are typically older MS stars and do not have appreciable Ha
the thin
shows
model
theexcess.
adoptedFor
MS,comparison,i.e., that would
makedotted
them line
appearbluerthethanet Thisal. (1998)
for stars with1s
theatmospheres
solid blue lineofinBessell
Figure 3.
would correspondto theeffective
deviation.
This isinthe
for 3500
AV=1.22.
all reddening
K - TThus,
K , surface
temperatures
thecaserange
eff - 40000
possible.
willvaluesin the0.22< aAmetallicity
log grange= 4.5[M H ]We=0.5, as
, and
index
gravityV < 1.22 area representative
and will
adopt
AV = 0.72
appropriate
for the0.50
LMCas(Dufour1984). Thevalue
theoreticalcolors,or aspeciﬁcallyeffects WFC
that assumingAV=0.22
calculatedaddress
for thethespeciﬁc
3 ﬁlters, alreadyincludefor these candidateon theAV=1.22, madeup oftypical reddening
value ofPMS
E (Vstars- Iwould) = 0.5have
derivedphysicalparameters.
owing to intervening MW absorption along
E (V I ) = 0.1
An
conclusion
is thattheowing draw
to LMCabsorption
the importantline of sightand E (Vthat- Iwe) =can0.4alreadyolder
PMS ofcandidates
are 30placedfront
theseeR136
cluster
in front
and within
Dor.inAs
we ofwill
later,
this is aandrepresentative
are separated
from reddening
it by a conspicuous
averagevalue for this amountﬁeld. ofabsorbingmaterial.
This
stems fromthe fact
that (1) asthethicka excess
emission
are shownObjects
with His wemaximumtheseDeobjects
symbols reddening
in Figure 3.toward
Following
Marchi(AetV=1.22)
al. (2011c),considerably
lower
the reddening
towardin theyounger
will consider
theirthanpositionsin the CMD
relationto theirand how(2) their
spatial their
PMS
starsdistribution(17th percentile,AV=1.63)
spatialin orderto determine
to correct
distribution is uniform over the ﬁeld. The separation in AV andphotometry for reddening using neighboring reference UMS
spatial distribution conﬁrms that younger and older low-massstars. All these objects with Ha excess emission are candidateFigure 5. H-R diagram of the bona ﬁde PMS stars (ﬁlled circles, with an
PMS candidates are associated with distinct star formation
additional cross if older than 5 Myr). Objects hotter than 10,000 K (openPMS
as suchprobed
are relatively
We usedepisodes stars
acrossandthe volume
by these young.
observations.
In acircles) are shown but not not considered in this work. The typical o1s
theoretical
PMS isochrone
for a roughpreliminaryinto
separation
uncertainties on Teff and L are shown by the error bars. The three panels
particular,
it appears
that star formation
has proceeded
theof younger and older PMS stars. A more accurate agecorrespond to different reddening values for the stars classiﬁed as older thanregion.
in the CMD of Figure 3 (crosses). The adopted values are, respectively,
determination will be possible after correction for reddening. 5A Myr
=1.22 (panel (a)), 0.72 (panel (b)), and 0.22 (panel (c)). The dashed lines
The thin short-dashed line is the theoretical isochrone for aareV the PMS isochrones from the models of Tognelli et al. (2011), for ages of 1,
5. Physical Properties of Bona Fide PMS Stars
PMS age of 5 Myr and metallicity Z=0.007 from the models2, 4, 8, 16, 32, and 64 Myr, from right to left (note that where multiple
Inorder to etselectbona ﬁdePMS stars,
after
reddeningof Tognellial. (2011),
translated
into theWFC
3 bands forisochrones reach the MS and overlap, lines appear solid instead of dashed).
Also shown are the representative evolutionary tracks for masses ofcorrection
in all bands
we repeatedprocedure
discussed
inthe assumed
distance
modulustheanda combined
reddening
0.7, 1, 1.5, 2, 3, and 4 M: (solid lines).
a excess as
emission
but weSection4 to identify
objects
Using the5 Myrwith
PMSHisochronea reference,AV=1.22.
imposing
stricterasconstraints.
particular,
we onlyrespectively,
considered thehave marked
thick ﬁlledIncircles
and crosses,
4s level
V -line,Ha andexcessat theapproximately
stars
with aabove
reddening-corrected
objects
and below thathence

5.1. Bolom

The value othe dereddenedof interest hedeterminations
Braun et al. 19
discussed inbolometric lum
magnitude andmodulus of 18
L are shown b(these uncerta
correction). N
larger for bri
the V−I colvariations.
The three pa
each panel westars classiﬁed
discussed in S
0.72 (panel (objects are ma
in each panel
models of To
metallicity. Isoand 64 Myr,isochrones isand L (see eevolutionary t(solid curves)
The paucity
panels (b) andsuggest that thstars in thosefrom the mas
cannot a priorcorresponding

By means oFigure 5, webona ﬁde PM
discussed by R
De Marchi et ato the one p
provides the phave a givenof the measure
the properties
the initial mas
Brieﬂy, thecompared withmost likely ag
tracks used areinterpolated tof 0.025.
We ﬁrst deﬁ

ber 10The Astrophysical Journal, 846:110 (16pp), 2017 September 10

The Astrophysical Journal, 846:110 (16pp), 2017 September 10

De Marchi, Panagia, & Beccari

De Marchi, Panagia, & Beccari

Figure 7. Histograms of the age distribution of PMS stars. A constantlogarithmic step is used (factor of 2). The solid line provides the number ofstars inside each age bin, while the dashed line shows the apparent starformation rate in units of Myr−1.

Figure 6. Histograms of the mass distribution for the 1035 bona ﬁde PMS Figure 7. Histograms of the age distribution of PMS stars. A constantstars. Different line types correspond to different AV choices for the older PMS logarithmic step is used (factor of 2). The solid line provides the number ofobjects: AV=1.22 (solid line), AV=0.72 (long-dashed line), and AV=0.22 starstimes
insidewider
each thanage bin,while thedashed lineonshows
the apparent
star
the typicaluncertainty
the mass.
Different(short-dashed line). For completeness, the dotted line shows the mass formation rate in units of Myr−1.
The
AstrophysicalJournal,846:110(16pp),
2017
September
10
choices
for
the
older
PMS
line
types
correspondtodifferent
AV
Ha excess foremission
hotterbona
than ﬁde
10,000K thatof ofthethestars
withdistribution
Figure 6.distribution
Histogramsmass
the 1035
PMS
objects as mentioned above: AV=1.22 (solid line), AV=0.72
we do line
not consider
in this work.
for thePanagia,
older PMS
stars. Different
types correspondto different AV choices
De Marchi,& Beccari

The Astrophysical Journal, 846:110 (16pp), 2017 September 10

Figure 9. Mass acage. Diamonds ref
the stellar mass (sePMS stars older thbest ﬁt to the 30 D
De Marchi, Panagia,& BeccarTauri stars
from thof the symbols ind
solid line object
shows th
∼16temperatures
Myr. If wetheconsider
lowreddeningthe ∼300
evolution of a visc

Figure 8. Histograms of the mass accretion rate values of the bona ﬁde PMS
stars. As before, different line types correspond to different AV choices for theolder PMS objects: AV=1.22 (solid line), AV=0.72 (long-dashed line), andAV=0.22 (short-dashed line).

youngersincethan=0.22 (short-dashedThe(long-dashedand AVuncertainty
radius. However,attimes
wider thanline),
the typicalon the mass.line).
Different
histograms show a larger number of more massive stars whenvector The
crossestheevolutionary
tracks,in any
with
these
characteristics
our unaccounted
30 Dor sample, we ﬁnd a =Journal, 846:110 (16pp), 2017 September 10line types correspond to different AV choices for the older PMS extinction Astrophysical
would
lead
to
an
underestimateofthestellarmass,
the value of AV is higher, and correspondingly a smaller-0.55 0.03 and b = 1.25 0.19. De
Marchi consistentet al. (2013awith(solid ofline),
objects
as mentioned
V=1.22
does not allow
thus resulting in a systematically higher value of the R* M*
lower.
numberof massiveabove:stars Awhen
the value
AV AisV=0.72
a=0.6
already
pointedoutthatapproximate
valuesofandHartmannetal=0.22(short-dashedline).
The(long-dashedline),
andA
points.
SinceanVratio. On the other hand, this effect is relatively small, and a(difference between the times of exit and entrance) andHowever, the three mass distributions are rather similar, andThe 346
observ
known
and
aand
ve
b=1 resultin ﬁtswithshow
a largerdifferences
number of(~4moremassivestarsinwhen
relevant infor
ourboth NGC
typical uncertaintyof 0.1 mag
in E(V -smallI ), asresiduals
characteristic age (average of exit and entrance times). We histogramss ) areonly seen
thestatisticallysigniﬁcantagreement
withthis region
to anand
uncertaintyof ∼6% onalso
the forR* M
we identify
the processcells thatthroughout
its the themass
higher,and correspondingly a smaller case, corresponds
value
of 0.7–1.0
AV is M
NGC 602,
this is conﬁrmed
30* Dor.ignore
repeatallthis
forthatall star
the crosses
tracks resulting
from
range
:.
the
correlation
namely,NGC
3evolution.
For eachcell, weabove.
compute the crossing timeis range
lower. ratio. In summary, the combined statistical uncertainty on Ṁacc
numbermassive show
stars awhen
the broadvalue peakof AinV theinterpolation
mentioned
The ofhistograms
relatively(p=0.082).
NGC 602 (aIf=amounts
to
∼23%.
When
also
the
systematic
uncertainties
due(differenceThebetween
the fourtimesmatrices
of exitand entrance)
andHowever,massthat,
distributions
aremeasurements
rather similar,
andresults arecontaining
all the relevant
Wethreeunderline
while theseshow
1–2 M:.theZ=0.015,
7. Mass Accretion Rate and Environment
However, itasis tc
measurements
information
associatedofwith
evolutionary
tracks crossing
characteristicage (average
exittheandentrance times).
We~4s )stars,
the relative
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forming regions in the MCs and in our Galaxyc(dashedcan provid
observed
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6. Mass
Rate andParameters
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the rate,
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Figure 11. Effective mass accretion rate c shown
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depends
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7
show
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choices for thebirth rate,thecomes
weightislegend).
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we
have
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the
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Figure
10
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602 itself. From
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Although
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to
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5.
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Table
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circlessomeseveral
often howeverafunctionofoft. derived
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a
1
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:
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(1998;
12
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H
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in
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unaccounted(∼6.55; Russell & Dopita 1990). Thus, adopting Z: = 0.019
do not reveal the
we measured
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et al.at (2011c),
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e.g.,
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al.
2010,
for
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thenearest-neighbors
full range
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individual
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ait objectsLodders
Moreover,
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aslocal0.5–1.1
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consistentwithmodelsofviscousdiskevolutionby2010),
theSk
183
metallicity
value
on
the
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themethod,typical
on
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,
and
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formation
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past
2
Myr
is
∼150
M
the
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and ages. Objects appear more massive when A grows, and theby the best ﬁt to the data (dashed line). The other symbols, as
given in each p
emission over hours −1
or days (e.g., Fernandez et :al. 1995; Smithobjects: AV=1.22 (solid line), AV=0.72 (long-dashed line), and AV=0.22(short-dashed line). For completeness, the dotted line shows the mass
Hacells
excessemission
hotterthan 10,000
K that its
distribution
the stars all
withtheweofidentify
thatthat starcrossesthroughout
we do notevolution.
consider in thisFor work.
each cell, we compute the crossing time

Yasui+ 2009,10,16の主張とは逆センスV

Although the N enrichment in Sk 183 could in part result fromlower radiation p

Embedded Filaments in IRAS 05463+2652: Early Stage of Fragmentationand Star Formation Activities
1

2

3

3

L.
K. Dewangan, R. Devaraj, T. Baug, and
D. K.ofOjha
297 Embedded
Filaments
in IRAS05463+2652:
Early
Stage
Fragmentation and1
Physical Research Laboratory, Navrangpura, Ahmedabad—380 009, India; lokeshd@prl.res.in
2#8 StarFormation
Activities
Instituto
Nacional de Astrofísica,Óptica y Electrónica, Luis Enrique Erro # 1, Tonantzintla, Puebla, C.P. 72840, México
3

Departmentof Astronomy
and Astrophysics,of Fundamental Research, Homi Bhabha Road, Mumbai 400005, India
L. K. Dewangan, R. Devaraj, T. Baug ,Tata
andInstitute
D. K. Ojha
Received
2017
June
20;
revised2017
August31;
accepted 2017 September 1; published 2017 October 10
ApJ 848, 51 (2017)

Abstract
We present a multiwavelength data analysis of IRAS 05463+2652 (hereafter I05463+2652) to study starformation mechanisms. A shell-like structure around I05463+2652 is evident in the Herschel column density
map, which is not associated with any ionized emission. Based on the Herschel submillimeter images, several
parsec-scale ﬁlaments (including two elongated ﬁlaments, “s-ﬂ” and “nw-ﬂ” having lengths of ∼6.4 and∼8.8 pc, respectively) are investigated in the I05463+2652 site. The Herschel temperature map depicts all thesefeatures in a temperature range of ∼11–13 K. 39 clumps are identiﬁed and have masses between ~70–945 M:.
The majority of clumps (having Mclump 2 300 M:) are distributed toward the shell-like structure. 175 youngstellar objects (YSOs) are selected using the photometric 1–5 μm data and a majority of these YSOs are
distributed toward the four areas of high column density (25 ´ 10 21 cm−2; AV∼5.3 mag) in the shell-likestructure, where massive clumps and a spatial association with ﬁlament(s) are also observed. The knowledge of
observed masses per unit length of elongated ﬁlaments and critical mass length reveals that they are supercritical.
The ﬁlament “nw-ﬂ” is fragmented into ﬁve clumps (having Mclump ~ 100–545 M:) and contains noticeable
YSOs, while the other ﬁlament “s-ﬂ” is fragmented into two clumps (having Mclump ~ 170–215 M:) without
YSOs. Together, these observational results favor the role of ﬁlaments in the star formation process in I05480+2545. This study also reveals the ﬁlament “s-ﬂ,” containing two starless clumps, at an early stage offragmentation.
Key words: dust, extinction – ISM: clouds – ISM: individual objects (IRAS 05463+2652) – stars: formation –stars: pre-main sequence – stars: protostars

データ：+2652 site.Herschel
They found an extended molecular cloud associated
1. Introduction• 182.0-00.3,
Vlsr = -10.6 km/s, distance = 2.1 kpc
with I05463+2652
and referred
to as “182.0−00.3” cloud (ID
UKIRT UKIDSSGP survey
The investigationofﬁlaments
instar-forming
regionshas−1• dV = 1.7 km/s, R ~ 10 pc, Mcloud ~ 9000 Msun #70; Vlsr∼−10.6 km s ; line width=1.7 km s−1; radius
NRAO VLA Sky Survey
recently received much attention,c with the availabilityof the(Rc)∼
Herschel continuum data (e.g., André et al. 2010, 2016;
WISE,
Glimpse360
10 pc; Mass
(Mcloud) ~ 9000 M:) (see Table 1 and Figure 9(j)13CO 1-0 観測（Kawamura+ 1998）

Kainulainen et al. 2016, 2017; Li et al. 2016, and references
therein). There are several works of theoretical and observationalin Kawamura et al. 1998). Herschel far-infrared (FIR) andsubmillimeter (sub-mm) images are available toward the I05463

The Astrophysical Journal, 848:51 (10pp), 2017 October 10

Dewangan et al.

The Astrophysical Journal, 848:51 (10pp), 2017 October 10

toward the four are
structure (see black
clumps are also fou
clusters of YSOs a“nw-ﬂ.” However
associated with any
Together, the st
toward the clumps
elongated ﬁlament

Figure 4. (a) Herschel column density (N (H2)) map of I05463+2652. Several ﬁlaments are highlighted by curves, which are similar to those shown in Figure 1(b).
Black circles show at least four areas of high column density (5×1021 cm−2 (AV∼5.3 mag)). The boxes also encompass two other elongated ﬁlamentary
structures (“nw-ﬂ” and “s-ﬂ”; also see Figure 3(b)). (b) A two color-composite map (Herschel 250 μm (red) and Herschel column density map (green) images) of
I05463+2652. Here, the column density map is processed through an “Edge-DoG” algorithm. (c) It is the same as that shown in Figure 4(b). A shell-like structure anda cavity are labeled and are also highlighted by ellipses (also see Figure 2(b)). Other marked curves and boxes are similar to those shown in Figure 4(a). (d) The
distribution of identiﬁed Herschel clumps in our probed ﬁeld around I05463+2652. The identiﬁed clumps are highlighted by asterisks and the boundary of each clump
is also shown in the ﬁgure along with its corresponding clump ID (see Table 1). A shell-like structure and a cavity are also labeled. In all the panels, the position of
I05463+2652 is highlighted by a star. In each panel, the scale bar corresponding to 10 pc (at a distance of 2.1 kpc) is shown in the bottom left corner.

sources (e.g., Gutermuth et al. 2009). Using the 2MASS andGLIMPSE360 photometric data at 1–5 μm, we have examinedthe dereddened color–color space ([K−[3.6]]0 and[[3.6]−[4.5]]0). The dereddened colors were estimated using
the color excess ratios given in Flaherty et al. (2007).
Following the dereddened color conditions listed in Gutermuth
et al. (2009), we ﬁnd 134 (18 Class I and 116 Class II) YSOs in

help of previously known conditions (i.e., [3.6]0 < 15 mag forClass I and [3.6]0 < 14.5 mag for Class II; e.g., Gutermuth
et al. 2009). The dereddened 3.6 μm magnitudes werecomputed using the observed color and the reddening laws
(from Flaherty et al. 2007). Figure 5(a) shows the color–color
diagram ([K−[3.6]]0 versus [[3.6]−[4.5]]0). The red circles andblue triangles refer to ClassI and ClassII YSOs, respectively.

Figure 6. (a) Herschel column density map is overlaid with the identiﬁed YSOs in
our selected ﬁeld around I05463+2652. The marked symbols are similar to thoseshown in Figure 5(c). Several ﬁlaments are highlighted by curves, which aresimilar to those shown in Figures 1(b) and 4(a). Black circles indicate the areas ofhigh column density (5×1021 cm−2 (AV∼5.3 mag)). The background map issimilar to the one shown in Figure 3(b). (b) A two color-composite map(Herschel 250 μm (red) and Herschel column density map (green) images) is
superimposed with the surface density contours (in black) of YSOs. The
background map is similar to the one shown in Figure 4(b). The maps show thepresence of the clusters of YSOs toward the areas of high column density in theshell-like structure and the ﬁlament “nw-ﬂ.” The contours are shown at [3, 6,
9]×σ (where, 1σ=0.4 YSOs pc−2), from the outer to the inner side. In bothpanels, the boxes encompass two elongated ﬁlamentary structures (“nw-ﬂ” and “sﬂ”; also see Figure 3(b)). In each panel, the position of I05463+2652 ishighlighted by a star and the scale bar corresponding to 10 pc (at a distance of
2.1 kpc) is shown in the bottom left corner.

1σ=0.4 YSOs pc−2), from the outer to the inner regions. The

Several authors
found at the juncti
star-forming comp
SDC335.579-0.292
etc. (e.g., Myerset al. 2013; Baug
This particular reformation of star-fo
I05463+2652 siteembedded shell-li(including “nw-ﬂ”
virial mass of th
published NANTE
we ﬁnd that the mo
is unstable against
embedded shell-lik
and many of thesehigh column densit
where star form
Section 3.3). Furthwith the areas of hi
Section 3.2). Note
data to trace the emsite does not conta
surface density >clusters of YSOs
indicating that thestellar cluster is oc
conditions of cluste
this site using hobservations canvelocity ﬁeld of thlinked with the I05
In recent years,
identifying embedd
tion, where starKainulainen et al.
using the NIR, 870
and suggested tha
candidate for a ﬁlashows very few sig
identiﬁcation of su(e.g., Kainulainen
have selected two“nw-ﬂ”). The Hers
with a temperaturethermally supercrit

P. Elliott , A. Scholz , R. Jayawardhana , J. Eislö↵el , and E. M. Hébrard

1708.03711v2 [astro-ph.SR] 22 Aug 2017

1

Department of Physics and Astronomy, York University, Toronto, ON M3J 1P3, Canadae-mail:297
Apmelliott0@gmail.com
deep staring campaign in the σ Orionis cluster Variability in substellar
2
SUPA, School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, UK
3#9members
Thüringer
Landessternwarte Tautenburg, Sternwarte 5, D-07778 Tautenburg, Germany
P. Elliott, A. Scholz, R. Jayawardhana, J. Eislöﬀel, and E. M. Hébrard
Received
18/04/2017; accepted 09/08/2017
arXiv:1708.03711v2
ABSTRACT
Context. The young star cluster near

Orionis is one of the primary environments to study the properties of young brown dwarfs
down to masses comparable to those of giant planets.
Aims. Deep optical imaging is used to study time-domain properties of young brown dwarfs over typical rotational timescales and to
search for new substellar and planetary-mass cluster members.
Methods. We used the Visible Multi Object Spectrograph (VIMOS) at the Very Large Telescope (VLT) to monitor a 240 ⇥ 160 field in
the I-band. We stared at the same area over a total integration time of 21 hours, spanning three observing nights. Using the individual
images from this run we investigated the photometric time series of nine substellar cluster members with masses from 10 to 60 MJup .
The deep stacked image shows cluster members down to ⇡5 MJup . We searched for new planetary-mass objects by combining our
deep I-band photometry with public J-band magnitudes and by examining the nearby environment of known very low mass members
for possible companions.
Results. We find two brown dwarfs, with significantly variable, aperiodic light curves, both with masses around 50 MJup , one of which
was previously unknown to be variable. The physical mechanism responsible for the observed variability is likely to be di↵erent for
the two objects. The variability of the first object, a single-lined spectroscopic binary, is most likely linked to its accretion disc; thesecond may be caused by variable extinction by large grains. We find five new candidate members from the colour-magnitude diagram
and three from a search for companions within 2000 au. We rule all eight sources out as potential members based on non-stellar shape
and/or infrared colours. The I-band photometry is made available as a public dataset.
Conclusions. We present two variable brown dwarfs. One is consistent with ongoing accretion, the other exhibits apparent transientvariability without the presence of an accretion disc. Our analysis confirms the existing census of substellar cluster members down to⇡7 MJup . The zero result from our companion search agrees with the low occurrence rate of wide companions to brown dwarfs foundin other works.

1. Introduction
The discovery of brown dwarfs in 1995 (Nakajima et al. 1995;
Rebolo et al. 1995), in conjunction with the discovery of the first
exoplanet around a solar-type star in the same year (Mayor &Queloz 1995), has triggered a significant revision in our ideasof star and planet formation. In particular, instead of a bimodal
view where stars form from cores and planets in discs, our cur-regions (see review by Luhman 2012). For each star, about 0.20.5 brown dwarfs are formed, in all regions studied so far, with
only a minor, if any, dependence on environmental conditions(Scholz et al. 2013; Muzic et al. 2017). Brown dwarfs can hostmassive discs (Testi et al. 2016) and show signs of accretion,
just like young stars. The widely-accepted view we have today
is that most of the more massive brown dwarfs are an extensionof the stellar mass function and form in a way similar to low-t al.: A deep staring campaign in theOrionis cluster

P. Elliott et al.: A deep staring campaign in theMOS) at the veryof deep images toollowing the prel. (2004), Scholz
dy & Hillenbrandcontrast to theseted masses below
cking our time seep optical images.
or new candidate
metry and by con-Table 1. Basic properties of young members and photometric candidates instudied in this work.

still somewhat unon is significantly
what younger thanthat is, between 1
et al. (2008), andge estimates of 2-

RA
hh:mm:ss.s

DEC
dd:mm:ss

Mayrit 258337
Mayrit 396273
Mayrit 379292

05:38:38.1
05:38:18.3
05:38:21.4-02:32:03-02:35:39-02:33:36[MJO2008] J053852.6-023215[BNM2013] 90.02 782[BNM2013] 90.02 1834
[BZR99] S Ori 51[BZR99] S Ori 50[BZR99] S Ori 58

05:38:52.6
05:39:12.9
05:39:00.3
05:39:03.2
05:39:10.8
05:39:03.6-02:32:15-02:24:54-02:37:06-02:30:20-02:37:15-02:25:36

RV, g, d
RV, g
RV, Li,
g, d
RV, g
H↵
H↵, d
g
...
H↵, d

[BZR99] S Ori 60[BZR99] S Ori 62[BZR99] S Ori 65

05:39:37.5
05:39:42.1
05:38:26.1-02:30:42-02:30:32-02:23:05

H↵, d
H↵
d

Orionis from Peña Ramı́rez et al. (2012) that are
Jb
(mag)

M
(MJup )

Var? cComments

Y(4.5, 8.0, 12.0)
N
Y(12.0)

15.07
15.29
15.31

56
47
47

Y
Y
...

SB1, Known var.

N
N
Y(4.5, 8.0)
NN
Y(4.5, 8.0)

16.18
16.68
17.19
17.16
17.47
18.42

29
24
19
19
17
11

...
...
...
...
...
...

19.02
19.14
20.30

8
8
5

...
...
...

IR Excess

Variability and deep image analysisr

Photometric cand.

Deep image analysis onlyY(8.0)
N
Y(4.5, 8.0)

P. Elliott et al.: A deep staring campaign in theNotes. (a) RV: Radial velocity consistent with systemic cluster velocity, g: low-gravity atmosphere, Li: Lithium absorption, H↵: Strong, broad H↵
VIMOS/VLT
I-band(c) observations.
Below
are the frommainthefindings
emission, d: Presence of a disc.(b) J-band magnitude from Peña
Ramı́rez et al. (2012).
Indicates if variability
was identified
analysis
P. Elliott et al.: A deep staring campaign in the Orionis clusterin this work.
and conclusions from our analysis.

Fig. 1. The observed pointings of our VIMOS/VLT observations
overlaid on the 2MASS J-band image. The navy blue and red
rectangles are the four quadrants for Fields A and B, respectively. The blue and green sources are young and photometriccandidate sources, respectively, presented in Peña Ramı́rez et al.
(2012).

1.3
1.2
1.1
1.0
0.9
0.8
0.7
0.6
0.5

1.0
Mayrit 258337, 05:38:38.1 -02:32:03

0.8

4. Photometry

0.6

estimated from unresolved photometry, the masses of the individualcomponents
are lower
than that.variability
The primaryinistwo
mostyoung brown– We haveidentified
significant
likely⇠40 MJup
andnewly
the secondary
MJup, assuming
a 0.7 magPrimarily,
wewere⇠25
interestedindi↵erentialindwarfs,one
identified,froma sampleof photometrynine.
di↵erence.our analysis, therefore absolute calibration of magnitudes was– Given the short time span of observations and the strong
not necessary.
Our approach
was isto ⇡0.5
use amagsamplereference
The light
curve
peak-to-peakamplitude
for of weinter-nightvariationsin theirquasi-periodic
signal,could
stars (seein eachquadrantof theeachpointing,
initially
selected by their
Mayrit 258337
Figure3). From
GLS
periodogram
anal0.0
To extract the sources’ photometry
andastrometryfromourimcalculate
aobservations,definitivefor
either1.0
lowI-band
standard
deviation,period
to foundaccountforanyvariations
in magniysisnotof thesewea seriesof object.
signifages,
we usedthe cluster
astropy-affiliatedphotutils package (Bradley
– The
first
object,Mayrit258337(a single-lined
0.8
P. Elliott et al.: A deep staring campaign
in theOrionis
icant
peaks.
The
peak bywithhighest observing
power
is atconditions.
0.77 day.spectroscopic
tude
causedthethechanging
The process
0.4
et al. 2016) in Python. Manyof the extraction algorithmsWeuseding the evolutionary models of Bara↵e et al. (2003) and Bara↵etried towas
useshows
thisperiodto fit of
the consistentdata. However,properties
there was with other0.6
asfollows.
binary),ahost
0.3within this package are the same as those used by SExtractor
et al. (2015).
clearly
still variable
a lot of higher-orderfeaturesthe variability
sig-variable mid0.4
Table 2. Photometric propertiesof all variable sources identiyoung
such unsaturated
asincorrelatedandWe first objects,
selected bright,
sourceswith a low stan0.2(Bertin & Arnouts 1996). 0.2
nature.
Therefore, at this time, it is unclear whether this peak at
fiedin this work as well as all non-variable properties fromeidarddeviationcompared
to
the
other
sourcesinthesame pointinfraredmagnitudes
andmid-infraredexcess.
Therefore,
its
⇠0.77
day
is
relatedtotherotation
of
one
of
the
objects
inthe
The orfirst
step was tobuildanmembers
input catalogue
of sources
0.1 young0.0
therphotometriccandidateof Orionis
from ing. We plotted the light curves of all of these sources and re1.0
variability
in
the
I
band
is
most
likelylinkedtoitsaccretionsystem.
3. Observations and data reductionfor
each
quadrantof
each
pointing.
We
selected
the
image
with
0.2 Ramı́rez et al. (2012).
Peña
moved any that were obviously inconsistent with the others, that0.8
the best seeing and extracted
sources using a 5 criterion. disc.
We
In the isfirstof the withrightsignificant
hand panelschangesof Figure
4, we can see
0.0
sourcesin brightness
not seen in other
The data presented in this work were obtained on 24, 25,
0.6
then used this list of positions to extract athe astrometrya andthat
there issignificant
correlation
betweenW1 andW2 curvemag–⇤phosecond
object,Mayrit396273,has
no mid-infrareda Thesources.
Wethen calculated
a themedianlightfrom the re0.2
and 26 December 2006 using the VIMOS instrument at the
ID
Arb.
m
+
3
0.4 I
m
e
tometry of sources in the time series of images. To accountfor for this
nitudesThis is showncorrelation
the consistently
lowexcessandobject.
nosignificant
variation
in indimidmaining
sources,
subtractedbythis
medianorfrom
each of the
(mag)
(mag)
(mag)
VLT, Paranal. The VIMOS instrument has four Charge-Coupled
0.4
0.2
any potential pixel drift between
individual observations, (.0.05)
we re- p values in the synthetic samples and average r value ofviduallight
curves,excluded3
outliers,
andcalculated
thesuminfraredmagnitude.
The
observedvariability
in
the
I
band
Devices (CCDs) each containing 2048 ⇥ 2440 pixels. The pixel
variablesources
0.6centred our apertures All
0.0
before1.0 extracting the photometry in⇡0.5.
eachA GLS periodogram of the two series of WISE magnitudes
oftheresiduals
for
each.
If
the
residuals
of
the
light
curve
were
scale is 000 .205, providing a field of view of 70 ⇥ 8.30 in each
0.20
may
be
causedbyvariableextinction
bylarge
grains.
(second
Mayrit258337 exposure. The typical
-0.95
0.143
0.051< 2 pixel
2.804
individual
pixel
drifts were
for right hand panel of Figure 4) shows that we found one
0.8
<0.03
the
curve
remainedin thereferencestars oflist;
quadrant.
0.15
Mayrit396273-0.33 to account0.063 for0.049
1.291
(⇡2.5
, ⇡1%
false alarmprobabil–inWe
notmag,findperiod
anylight
new
low-mass
potential
members
all our
observations.
In ordervariation marginallythe didsignificant
0.6
b
ifday
not,initthe
wasomittedandforthetheprocedure
was run again using
0.10
The observations spanned approximately seven hours each
UKIDSS
442414579467
0.09
0.170
0.130
1.308
ity)
at
0.63
W2
magnitude
secondModifiedbackground signal, we performed
a local background subtracOrionis using
ourasuncalibrated
I-bandphotometry
witha avail0.4
0.05
new range.
list
an input. The
criterionof 0.03mag wascomUKIDSS 442414579362c
-0.11
0.033
0.028
1.193 Datethenight, and alternated between the two pointings, presented in
Julian(MJD)
Additionally,a similar
periodis foundtion for each extracted source.
We calculated the average
flux
able
UKIDSS
J-bandphotometry,
consistentwiththe results
0.2
0.00
promisebetweenhavingenoughreferencestars(typicallyten or(atamuch
higherfalse
alarm
probability
of
⇡30%)
in
the
W1
this work as Fields A and B. Figure 1 shows the observed pointNon-variable fromY and Caroundcatalogues
withinaneight-pixel-width
annuluseachsource.
The
in0.050.0
of samePeñaRamı́rez
et al. the(2012).
more)andachievinga highThe observamean precision of ourin0.280
the
MJD
range.
Given
shortprecision.
span of theseings overlaid on the Two Micron All Sky Survey (2MASS) J0.0 edge
0.5
1.0this1.5annulus
2.0
2.5 -0.76
0.2 by0.4
0.6
1.0
1.2
Mayrit379292
ner
of
was
set
two0.017
times0.8the0.059
Full-Width
Halfit is hard to conclude definitely on this recovered periodic+5.4093 10
MJD
Period (day)tions,
masterreferencelight
curves low-mass
for all pointings
is between around11 andband image. Each individual exposure was 300 s long and taken
–
We
did
not
identify
anycompanions[MJO2008]
J053852.6-023215
0.36
0.024
0.138
0.177
Maximum (FWHM) of each individual observation. We signal.
multi- However, it may be that the new
0.77
dayworse,
signalthan
fromtheour7ob19
mmag, comparable,
albeitmmag physicalachievedin the I-band. The final dataset was comprised of 112 science[BNM2013] 90.02 782
0.025
0.266
forty
young
Orionissourcesintheapproximate
plied this average flux value1.20
by the areaof the 0.095
source’s apertureservations
andthe 0.63
day signal(2005).
from AllWISE
data are90.02 1834
0.174
0.221
0.787
in
Scholz& Eislö↵elmasterlightre-curves were fiimages for each of the pointings and additionally 15 dark sky
Fig.[BNM2013]
3.and
Left
panels:
VIMOS
light 1.55
curvesthevariable
subtracted
the resultant
value.forAs
ourfour
input
catalogue
wasseparation
range
au, These
consistentwithstudieslated
to thenallyrotationof the500-2000
object.
Givenwidthof
ofall
theindividual
peaksother[BZR99]S Ori 51
1.94
0.053
0.197
0.272
subtracted
from thelightthecurves
sources in
exposures from the three nights. The reduction of the data wassources
identified
in
this
work
(Table2).
The
redmarkers
aretheconstructed50from the image2.13
with the best seeing,[BZR99]0.268 the value
0.316
in theperiodograms,
valuesareresults
consistent.
ofofday)wide
multiplicity
in 2thevery
low-mass
objects.
thepointing.
Figureshows
ourfor Field A, quadrant 1.
lightcurves SofOrieach source, the grey markers0.085
are the light
curve (±0.2
done for each of the separate quadrants to account for any CCD4.1. Time series of relative0.4photometry
0.2

Mayrit 396273, 05:38:18.3 -02:35:39

Normalised Power

Arbitrary magnitude (mag.)tar Orionis haring from massiveg. Garrison 1967;. 2004; Caballeron and modest disr an ideal groundas the early evo. Deep surveys ofwn dwarfs (Béjar
th masses compat al. 2001; Bihain
deepest large-areacarried out basedd Infrared Survey
RT Infrared Deepña Ramı́rez et al.
. Throughout thisung members andmı́rez et al. (2012).
l sources.
onis cluster range
rtainties when indistance value isal. 1997). We usedtrometric Solution
ollaboration et al.
for the Orionis
arch radius of-500 pc based onmembers. Elevenmember catalogue
94 ± 0.40 mas, but
the average error)
to 352±51
40 pc. Thiss most of the prepublished estimaaefer et al. 2016)

Feat. aResolvable Simbad ID

Orionis cluster

05:39:04.2 -02:31:11

05:39:07.8 -02:30:55

4flux forvery
exposureswith much[BZR99]S Ori
58faint sources in
3.62
0.119
0.508worse seeing
0.129
source of similar magnitude, shown for comFig. 2. Standard
deviation of the light curves in Field A, quadrant of a calibration
dependent
behaviours.
could
be
0
or
below,
due
to
local
backgroundsubtraction.
In
suchparison.
Rightpanels: described
Generalised
Scargleperiodograms
Our uncalibrated I-band photometry for sources successfully
(a)
1 for
Orionis versus
medianuncalibrated
I-band magnitude.
The observing
conditions
werevariable throughout
the threeNotes.
Properties
in Lomb
Equation
1. (b)was
Backgroundfield5.2.2.
cases
the
photometry
for
these
aperturesnotincluded
instar:ourMayritfor
eachvariablesource.
The
dot-dashed,dashed,
and
solid
grey
4.2.396273
Stacking
individual images(c)cross-matchedwith UKIDSS/DR9
GCS catalogue (Warren et al.
nights of observations. The median seeing and standard devi05:39:04.2,-02:31:11 Background field star: 05:39:07.8, -02:30:55.
final
analysis.
lines
are
1
,
2
,
and
3
intervals
from1000
bootstrap
samation of the science frames from each night was 0.84±0.20,
2007)
isavailable
publiclyvia
theVizieR
service.
Mayrit396273,firstreportedinBéjaretal.
(2004),
hasno in
midples.
We
removedany
imagesthat
hadaseeing(FWHM)
attheThe
secondaspectoftheanalysispresented
this work is the
0.89±0.66, and 100 .23±0.66. This large variation in seeing lim00
infraredexcess
out
to
8
µm
(no.
446
in
the
list
by
Hernández
timeofobservation
>1
.5
(23
exposures).
This
was
to
avoid
stackingofindividual
exposurestocreatea
deep
image of theitedability tofrom
removetherespective
small-scale
fringeThestructure
was our
subtracted
each
image.
final resultdeep imAcknowledgements.
We
thank
José
A.
Caballero,etal. 2014). Orionis
The AllWISE(Wright
et al. In
2010;
Mainzer
et al.for
The calculated
valuescontaminating
are shown inthe
Table
2. In theflux
case
that
neighbouring
sourcesbackgroundcalcuregion
intheI referee,
band.
order
to
account
forconstructive
the smallingnight pointing
sky emission.
adoptedtaken
the sameapproachthat
2011)
database
contains
fluxes
at
12
and
22
µm
for
this
source,comments.
This
work
has
made
use
of
data
from
the
EuropeanAgencyagefromfor each
was Wea medianfromthe realigned
a singlesourcehad time
seriesphotometry
two presented
separate variations (typically <2 pixel) of source positions Space
lated inthe annulus
aroundeachsource. Thefrom
analysis
between
ob00
is
used
in
Alcalaetal.
(2002),
LópezMartı́
etal.
(2004),
andwithasignal-to-noise
ratio
of
14
and
4,
respectively,butthede(ESA) mission Gaia (http://www.cosmos.esa.int/gaia), processed by theindividual exposures which had a seeing value 1.5 . The stan-OBéjar, V. J. S., Zapa
Bell, C. P. M., Nay
MNRAS, 434, 8
Bertin, E. & Arnou
Bertout, C. 1989, A
Bihain, G., Rebolo,
Bozhinova, I., Scho
Bradley, L., Sipocz
Burgasser, A. J., Ki
Caballero, J. A. 200
Caballero, J. A. 201
Caballero, J. A. 201
Caballero, J. A., A
521, A45
Caballero, J. A., B
A&A, 424, 857
Caballero, J. A., M
143
Caballero, J. A., No
Casali, M., Adamso
Cody, A. M. & Hil
Cody, A. M. & Hil
Cody, A. M., Stau↵
Cutri, R. M., Skru
Catalog, 2246, 0David, T. J., Petigu
Duchêne, G. & Kra
Franciosini, E., Pal
Gaia Collaboration
Gaia Collaboration
Garrison, R. F. 196
Hambly, N. C., Col
Hernández, J., Calv
Hernández, J., Hart
Hewett, P. C., War
367, 454
Hodgkin, S. T., Irw
675
Janson, M., Bergfo
Kouwenhoven, M.

F. Ménard4 , M.R. Meyer1, 16 , J. Milli11 , D. Mouillet4 , J. Olofsson17, 7, 20 , A. Pavlov7 , J. Pragt18 , P. Puget4 ,S.P. Quanz1 , R. Roelfsema18 , B. Salasnich3 , R. Siebenmorgen12 , E. Sissa3 , M. Suarez12 , J. Szulagyi1 , M. Turatto3 ,S. Udry19 , and F. Wildi19

iv:1709.00417v2 [astro-ph.EP] 7 Sep 2017297 The HIP 79977 debris disk in polarized lightEngler, H.M. Schmid, Ch. Thalmann, A. Boccaletti, A. Bazzon, A. Baruﬀolo, J.L. Beuzit, R. Claudi, A. Costille, S.
#10 N.
(Affiliations
can be foundafter thereferences)
Desidera, K. Dohlen, C. Dominik, M.
Feldt, T. Fusco,
C. Ginski,
D. Gisler,
J.H. Girard, R. Gratton, T. Henning, N. Hubin,
M. Janson, M. Kasper, Q. Kral, M. Langlois, Received
E. Lagadec,
F. Ménard,...; accepted... M.R. Meyer, J. Milli, D. Mouillet, J. Olofsson, A.
Pavlov, J. Pragt, P. Puget, S.P. Quanz, R. Roelfsema, B. Salasnich, R. Siebenmorgen, E. Sissa, M. Suarez, J. Szulagyi,
M. Turatto, S. Udry, and F. Wildi
ABSTRACT
arXiv:1709.00417v2
Context. Debris disks are observed around 10 to 20 % of FGK main-sequence stars as infrared excess emission. They areimportant signposts for the presence of colliding planetesimals and therefore provide important information about the evolutionof planetary systems. Direct imaging of such disks reveals their geometric structure and constrains their dust-particle properties.
Aims. We present observations of the known edge-on debris disk around HIP 79977 (HD 146897) taken with the ZIMPOL
di↵erential polarimeter of the SPHERE instrument. We measure the observed polarization signal and investigate the diagnostic
potential of such data with model simulations.
Methods. SPHERE-ZIMPOL polarimetric data of the 15 Myr-old F star HIP 79977 (Upper Sco, 123 pc) were taken in the VeryBroad Band (VBB) filter ( c = 735 nm,= 290 nm) with a spatial resolution of about 25 mas. Imaging polarimetry efficiently
suppresses the residual speckle noise from the AO system and provides a di↵erential signal with relatively small systematic
measuring uncertainties. We measure the polarization flux along and perpendicular to the disk spine of the highly inclined disk
for projected separations between 0.200 (25 AU) and 1.600 (200 AU). We perform model calculations for the polarized flux of an
optically thin debris disk which are used to determine or constrain the disk parameters of HIP 79977.
Results. We measure a polarized flux contrast ratio for the disk of (Fpol )disk /F⇤ = (5.5 ± 0.9) · 10 4 in the VBB filter. The
surface brightness of the polarized flux reaches a maximum of SBmax = 16.2 mag arcsec 2 at a separation of 0.200 0.500 along
the disk spine with a maximum surface brightness contrast of 7.64 mag arcsec 2 . The polarized flux has a minimum near thestar < 0.200 because no or only little polarization is produced by forward or backward scattering in the disk section lying in frontof or behind the star. The width of the disk perpendicular to the spine shows a systematic increase in FWHM from 0.100 (12AU) to 0.300 0.500 , when going from a separation of 0.200 to > 100 . This can be explained by a radial blow-out of small grains.
The data are modelled as a circular dust belt with a well defined disk inclination i = 85(±1.5) and a radius between r0 = 60
and 90 AU. The radial density dependence is described by (r/r0 )↵ with a steep (positive) power law index ↵ = 5 inside r0 anda more shallow (negative) index ↵ = 2.5 outside r0 . The scattering asymmetry factor lies between g = 0.2 and 0.6 (forward
scattering) adopting a scattering-angle dependence for the fractional polarization such as that for Rayleigh scattering.
Conclusions. Polarimetric imaging with SPHERE-ZIMPOL of the edge-on debris disk around HIP 79977 provides accurate
profiles for the polarized flux. Our data are qualitatively very similar to the case of AU Mic and they confirm that edge-on
debris disks have a polarization minimum at a position near the star and a maximum near the projected separation of the main
debris belt. The comparison of the polarized flux contrast ratio (Fpol )disk /F⇤ with the fractional infrared excess provides strong
constraints on the scattering albedo of the dust.
Key words. Planetary systems – Scattering – Stars: individual object: HIP 79977, HD 146897 – Techniques: high angularresolution, polarimetricThis does not include the innermost regions |x| < 0.2 . OnlyEngler et al.: HIP79977
debrispolarization
disk in polarizedsignal
light is expected at small apparent sepa small
arations for a disk or ring with an inner radius r > 0.200 ,because at small separations we observe scattering from thedisk sections located in front of and behind the star. This forward and backward scattering produces only little polarizaNtion. Thus, one can approximate the innermost disk with alinear extrapolation of the measured curve from P(x = 0.2700 )E
to P(x = 0.000 ) = 0 (red dotted line in Fig. 6(c)).
This neglects a possible contribution of polarized flux
from warm dust located very close (r < 0.200 ) to the star.
Studies on the spectral energy distribution of HIP 79977 (e.g.,
Chen et al. 2011) indicate that there is no significant (& 1%)
signal to the IR excess emission from warm dust at small separation. Therefore, we assume that there is also no significant
unresolved contribution from an inner disk to the polarization
signal.
The polarized flux in the VBB filter, covering an e↵ective
aperture area of 3.600 ⇥ 1.800 and including the interpolated
0.5
points inside interval |x| < 0.300 , amounts to 5800 counts per’’
second and per ZIMPOL arm. This value must be corrected
for the variable atmospheric transmission T atm (Fig. 1) usinga factor of fcorr = 1/T atm = 1.3 ± 0.1. This yields a corrected
count rate of 7540 ± 800 cts/s where the uncertainty is domiFig. 2. Composite image of debris disk around HIP 79977 with the VBB and I-band filters obtained with LOCI data reduction. The original
nated byfcorr
data were 3 ⇥ 3 binned to reduce the e↵ect of the noise. The positionof thestar. is marked by an asterisk in orange. The white dotted line
Engleretal.:
HIP79977
debrisdisk
in color-scalepolarized lightshows the position of the expected lines of nodes for an inclined disk
ring.
The
is givenin arbitraryThe
determinationof thestellarunits.
flux of HIP 79977 mustaccount for the saturation of the PSF core and the cloudy
p
7
weather.
Welevelfirstanddetermine
thesystematicmean value
cts/s inN
noiselarge-scaleerrorsofin1.13·10
our observaP = Q2 + U 2 . However,
P is a↵ected for low signal-tonoise data by a systematic bias e↵ect because of squaring
tions.
theofVBB
filter for frames 210-280 which were apparently notφφa↵ected
Q and U parameters.
Therefore we characterize the disk
poE
Very
close
to
the
star,
markedbyawhitecirclewithaby clouds (Fig. 1). Because the exposure is saturated
larization pattern with a locally defined azimuthal / radial Q- r ' 0.1200 (Figs. 3 and 4), the data are unreliable because of
out
to
the
radius
r
3
px
some
flux
is
lacking.
Toaccount
and U-parameter definition with respect to the central light strongly variable wings of the PSF peak. Also visible are thefor thefaintsaturatedthe00 above
PSF, and
we belowcomparethewhich
HIPare79977
source as discussed in Schmid et al. (2006). Single scatteringfeaturespart
at r &of0.12
the diskwith high-quality
ZIMPOLPSFs
of as
thepositive
standardo↵ dust particles in optically thin debris disks generatesprofilelin- negativein the Q and U
images, andappearsig- star
early polarized light with the electric field vector azimuthally
nal in the
Q' and U' images. Thesefeaturesetare
fainter
HD 183143(STD261_0013-24,
Schmidal.much2017),
whichoriented with respect to the star. Polarization in the azimuthal(factorunder< 0.1) excellent
than the diskatmosphericsignal and originate
most likely
were taken
conditions.
For thedirection is defined by the Stokes parameter Q' :from poorly corrected instrumental e↵ects because an intrinnarrowsic
band
filters
N_R to( produce= 57 nm) and N_I
c = 646nonm,
signalis expected
U' signal.
Q' = (Q cos 2' + U sin 2'),( (3)= 81 nm), these PSFs contain within a rac = 817 nm,
dius of r = 5 px a flux between ⇠ 20% and ⇠ 25% of the total
where ' is the polar angle between north and the point of in- 4. Data analysis
stellar flux measured for an aperture of 300 diameter. Based onterest measured from the north over east. The Stokes paramethis, we4.1.
assumefor ourangle
HIP 79977 data, that the round annu0.5 ’’
Disk position
ter U' :lus with
inner and outer radii rin = 5 px and rout = 416 px
U' = Q sin 2' + U cos 2'(4) 00 We measured the position angle of the disk in the Q' -image(= 3 diameter)
containsofbetween⇠ 75%ofand⇠ 80%lineof theby the determination
the orientation
the mirror
forcentral
an unsaturated
PSF profile.
Thisyieldsthrough thestar perpendicular
to the disk.
The
best fordefines the polarization pattern in the directions ±45 flux
with expected
angle wasfoundrates
by searchingan angle
respect-5to the Q' -2
direction. 1
the corrected
stellarcountbetweenwith
1.33· 10257 incrects/s and47
10position 1316
19
22
7 of 0.1 the orientation of the mirror line which produces
ment
Figure 4 shows the final Q' and U' . The Q' image clearly
1.40 · 10
cts/s per ZIMPOL arm for observations in the VBB
reveals the nearly edge-on disk structure down to a projected the smallest residuals if one side is subtracted from the otherfilter in the slow polarimetric mode.
separation of ⇠ 0.100 . Polarized light is detected across the side.
rates fromare converted
to photometricmagnitude
The results
the polarimetricand imaging data
sets
entire width of the image of ⇠ 3.600 . The peak of the surfaceThe countm(VBB)
using
theincluding
following
expression(Schmidal. 2017):
agree.
After
ZIMPOL’s
True Northo↵setetof2
brightness appears here as a narrow stripe below the expectedUQ

axis to be ✓disk =major axis of ESE
an inclined circular ring (white dotted line) we obtain the position angle of the disk
WNW
114.5
0.6log(cts/s). This valueam·k
is in 1good
agreement
PA (VBB),=
with a flux minimum near the position of the star.
m(VBB)= ±2.5(VBB)
mmodewith+zpima114 reported by Thalmann et al. (2013) for the scattered lightm
imagesinH-band
and
with
PA
=
115
measuredby
LiemanBy contrast, the U' image contains no structural features
where am = 1.15 is the airmass, k1 (VBB) = 0.086 is the filfrom the disk. Assuming azimuthal polarization of light gen- Sifry et al. (2016) in the sub-mm range.
ter coefficient for the atmospheric extinction, zpimathe(VBB)=
erated in single scattering processes and no multiple scatter- m We define an x y disk coordinate system wherestarphotometricpointalong
for theVBBaxisfiltering (see Canovas et al. 2015), we do not expect to find24.61
any isisat thethe origin,+x is thezero
coordinate
the major
in andm
mmode =
1.93
is an o↵set(✓tothe zero
which
astrophysical signal in the U' image. Therefore, this image
roughly
WNW-direction), xpoint
towards
ESE (✓accounts
disk +180
disk ),
can be used for an estimation of the statistical pixel to pixel and y perpendicular to this with the positive axis towards NNE

0.5 ’’

Article number, page 10 of 20

Johnson J

1.250

0.300

8.06

0.02

5

Notes: (1) ESA (1997); (2) Høg et al. (2000); (3) Gaia Collaboration (2016);(4) this work; (5) Cutri et al. (2003).

𝒚𝒉𝒙, 𝒚, 𝒛𝒕𝒐 𝒐𝒃𝒔𝒆𝒓𝒗𝒆𝒓𝒛 = 𝑳𝑶𝑺𝒚𝒅𝜽𝒊𝒙 = 𝒙𝒅Fig. 7. Illustrative sketch of the debris disk with inclination i andcoordinate systems (x, y, z) and (xd , yd , h) used in model. The small
blue cube at scattering angle ✓ marks the position (x, y, z) of a gridelement with grain number
densityn(x,y,diskz).in polarized lightEngler et al.:
HIP79977
debrisfor the used instrument and detector mode. We obtain for HIP
79977 a magnitude m(VBB) = 8.60m ± 0.07m in good agreement with the literature values (see Table 2). The derived
photometric magnitude m(VBB) yields the color V-VBB =9.09m - 8.60m = 0.49m which is close to the color index in theJohnson-Cousins’ photometric system V IC = 0.44m ( e↵ =0.806 µm,
= 0.154 µm for IC ; Pecaut et al. 2012) for a
F2/3V star.
For the polarized flux of the whole disk we get
mpdisk (VBB) = 16.6m ± 0.3m . This yields a ratio of
total polarized flux of the disk to the stellar flux of(Fpol )disk /F⇤ = (5.5 ± 0.9) · 10 4 .
We determine for the peak surface brightness of the polarized light SBpeak (VBB) = 16.2m arcsec 2 along the inner(0.200 0.400 ) disk spine (Fig. 6(a)) and a surface brightness
contrast for the polarized flux of SBpeak (VBB) mstar (VBB) =7.64 mag arcsec 2 . For the outer disk around x ⇡ ±1.700 thesurface brightness contrast is about 10 mag arcsec 2 .

5. ModelingTo reproduce the physical appearance of the debris disk
around HIP 79977 we construct a 3D model for the scattered intensity and the polarization flux from optically thin(single scattering) dust. The disk is described by an axisymmetric dust distribution using the cylindrical coordinates

Article number, page 5 of 20

61 AU

Fig. 4. Polarimetric di↵erential imaging data of HIP 79977 with the VBB filter (590-880 nm). The original data were 3 ⇥ 3 binned to reduce
the noise. The position of the star is marked by an asterisk in red. The upper panel shows Q' (left) and U' (right) images. Lower panel:
Isophotal contours of polarized light overlying Q' image. The contours were measured from the Q' image smoothed via a Gaussian kernel
with = 1.5 px. Contour levels are given for 3 (blue line), 9 (light blue), 15 (orange) and 21 (red) counts per frame per binned pixel.
The white dotted line shows the position of the expected ring axis. The region inside the white stellarcentric circle with radius ⇠ 0.1200 is
dominated by strong speckles variations. The color-bars show the counts per binned pixel.

Fig. 6. HIP 79977 debris disk properties along the axis x. The individual points give parameters of the Mo↵at profile of the vertical crosssection as shown in Fig. 5 and described in Sect. 4.2. From the top to the bottom: (a) the profile peak SBpeak (x), (b) FWHM, (c) verticallyintegrated flux P(x), and (d) spine distance from the disk major axis y0 The vertically integrated profile flux P(x) is calculated as a meansurface brightness in a 0.100 ⇥ 1.800 bin. At separations smaller than x ⇡ 0.200 the systematic uncertainties are increased and open circles markthe low S/N points. The vertical yellow line indicates the position of the star.
Article number, page 9 of 20

and fit for x . 0.600 because the Mo↵at profile cannot fit correctly negative flux values at small angular separations which
originate from the systematic e↵ects described above.
The vertical o↵set y0 (x) of the disk spine is shown iny0 ⇡ 50 mas, and also in the outskirts (|x| & 100 ) the y0 -o↵set
is even more than 50 mas. In comparison, the o↵set y0 (x) of
the disk spine measured in the imaging data (Fig. 2) is approximately 60 ± 5 mas (⇡7.5 AU) at |x| < 0.300 . For larger
separations, the y0 -o↵set in intensity is smaller and achieves

A&A proofs: manuscript no. HIP_79977_arch

et al.: HIP79977 debris disk in polarized lightrees of freedom

use the noise isresent flux intenferred from theich look essencalized spurious
disk. The fitting
model convolul, the key proprized flux alongle.
to fit the hP|x|ihe profile fitting
e range [60, 86]e maximum. Of
by 9 parameters
que solution forl defined ranges

mined by the 2model. This apl noise is larger
are much harderk models whichg to the Q' imdel parameters.
distribution of
he 2D image fit
distribution are
eir uncertainties
calculated frommeters together
le 3 (Col. 5 andthetic image ofconvolved imFig. 9(d)) well.
me PSF-shapedd below the disk
m the disk flux.
he spine at small

pth in the disk isimate a ⌧ ⇡ 0.5
plane (⇥ = 0 ),ering, a photon
we see the diskng allows an asem where many
he data. In par-A&A proofs: manuscript no. HIP_79977_arch

5.1. The scattering phase function for polarized lighta

The phase function (PF) f (✓) in Equation(7)polarimetriccharacterizes profile. Lieman-Sifry et al. (2016) also meathe
the angle dependence of scattered radiation. In the following,+3
sured
we disregard the wavelength dependence of
the PF.a disk inclination i > 84 and PA = 115 3 which are
A very popular way to describe the scattering
phase funcconsistent
with our and previous results.
tion is the Henyey-Greenstein (HG) function (Henyey &Greenstein 1941):f (✓) =b1
4⇡(1 + g2

g2,
2gcos✓)3/2

6.2. Diagnostic
potential of polarimetry(11)

scattered
where g is the average of the cosine ofThethe scatteringan- flux has been measured for more than 20 degle which characterizes the shape of the phaseForbris function.
disks, mainlywith the Hubble Space Telescope (HST)
isotropic scattering g = 0, forward scattering
have(e.g., grains
Schneider
et al. 2014, 2016). With ground-based ob0 < g  1, while for 1  g < 0 the scattering is peaked
servations the flux measurement for the scattered light frombackwards.
However, there exists also growing evidence
thatdisksa simpledebrisis very difficult because of the speckle noise inHG-function is a poor approximation for the modeling of thetroducedbytheatmospheric turbulence. The polarized flux of
scattered intensity from debris disks. This is nicely demon- Fig. 11. Comparison between alternative models of the HIP79977 debris disk: "Model 40" with the radius of the planetesimal belt r0 = 40
mucheasierto determine,
becausea parameters as specified in Table 3. From top to the bottom: (a)and "Model70" withthe radius of r0 = 70
AU (right) it
andis
fitting
strated for the bright disk HR 4796A (Millithe
et al.disk
2017),(Fwhichpol )diskAUis(left)
model of the dust distribution in the disk, (b) model of the polarized light, (c) model of the polarized light convolved with the instrumental
shows, for small phase angles ✓ < 30 , adi↵erential
strong di↵raction
quantity
which can be distinguished from the unPSF and (d) residuals left after subtraction of the PSF-convolved model image from the Q' image. Color-scales of images (a) and (b) are
peak and, for large phase angles ✓ > 30 , a scattering inten- Fig.
Scattering
phase
functionforthe (c)polarized
light (blue)
given8.inthe
arbitraryunits.
Color-bar
ofstarimagesand (d)fluxevenin forcounts at each pixel of the image.
brightcentral((F) shows⇡ 0),sity which is roughly angle-independent. polarized
Thus, a morelightgen- fromthree di↵erent asymmetry parameters gsca =pol0.2,⇤gsca = 0.4, anderal phase function, for example, a two-component(or dou- of
gsca strong= 0.6. Redlines show the corresponding
in the presence
atmosphericspeckles. Heyney-Greenstein0000ble) HG function seems to be required for the modeling of functions for f (✓, gsca ).
For HIP 79977 the disk profile in polarized flux reveals signala from the disk for separations from 0.3 to 1.5 , comthe scattered intensity of highly inclined debris disks
pared the polarization with the intensity profile and found a0cclear maximum which traces the radial location of the disk
fractional polarization of ⇠ 10 % (1 -range [5 %, 20 %]) at

5.2. Model fitting
Fig.
Ratio
0.500 and⇠ 4513.
% ([30%, 60of%])polarized
at 1.500 . Fromflux
their todata,theit scattered light luminosinformation is difficult to6 obtain from intensityw) · fring.
(✓, gsca )This,(12)
whetheror not theythin
see indebris
polarized
flux a maxiity for
opticallydisks
as a function of disk inclinaWe have calculated 5.28 · 10 models for a parameter grid asis not clear
00
imaging of edge-on disks, because the Stokes I disk profile
mum attion
a separationof around 0.6 asymmetry
and a flux decrease
inside.
and scatteringparameter
gsca (plotted for gsca =where the first term describes the strong di↵raction peak, the specified in Table 3 in order to find the set of model parame-Thalmann
et al. (2013) also fit the observations with model
thewhichforwardin front
ofimage.
the starbest fitscattering
the observed dust
polarized
intensity
second term represents the more isotropicis
anddominatedmuch less for-by ters
0.0, 0.2,
0.3,adopt..., 0.9).
Theisindependentof the disk geometry
calculations
but theya radiusof rratio
=
40
AU
for
the
For the fitting,
we extension
reduced the number
image
pixels byplanetesimal ring and do not investigate 0models with larger
ward scattering part (see also Min et al. 2010),
and w is thethe projected
and therefore
disk
may ofnot
be visible.
and follows from the scattering phase functions as shown in Fig. 8.
3 ⇥ 3 binning and selected a rectangular image area with ra .
scaling parameter, 0  w  1.
0
The diskfluxlength
Fdiskofand
polarized
contain
341 theand widthof 100 flux
binned(Fpixelscenteredandpol )disk
For the polarized scattered radiation from a debris
disk theThe new SPHERE - ZIMPOL observations presented into the diskabout
x and y the(majorscatteringand minor) axes
(see The
Fig.this work provide a very much improved polarimetric sensituation is slightly di↵erent. The strong complementary
forward peak seen aligned
information
dust.
in intensity, which can be ascribed to the light di↵raction by 9(d)). A round area with a radius of 16 pixels (0.1700 ) centeredsitivity whichThe
fractional
infraredexcessof HIP 79977 LIR /L⇤ =clearly
reveals a maximum
in thepolarizascattering angle dependence is strongly di↵erent because forlarge particles a, is expected to produce no significant on the star and the spurious features near the saturated regiontion profile
P(x)at 3a isprojected
separationof 0.60±0.0600et al. (2015). This yields
5.21
·
10
given
in
Jang-Condellward
andbackward
scattering
produces
no
or
only
very
littleareexcludedfromtheevaluation
of
the
fit
goodness.
Figurelight polarization. The scattering polarization is produced by(74±7 AU). This maximum location is not compatible withthering
double9 illustrates
the alsodi↵erent
in the image fitting
the smallradius ofratio
r0 = 40 AU adopted by Thalmann
the photons hitting the particle surface andpolarizedinteracting by
dif- Thisflux.
applies
to stepsthe di↵raction
peakprocedure.
(or forfuse reflection or/and refraction and transmission as described From the model dust distribution in the disk (a) the expectedet al. (2013), probably, because they only fit the intensity proward
scattering
peak)
from
large
particles
a
>
,
which
is
not
(Fno
)disk /Fthat
polfeatures⇤ could constrain the ring raabove by the second term f (✓, gsca ). But, in addition, the an- polarization flux is calculated (b), convolved with the instru-file which shows
Fig. 12.
Two-dimensional
constraints
on theradiusoftheresiduals
planetes- (e)= 0.11 ± 0.02,
dius. ⇤ =PSF(c), This
fitted tomeans
observation(d),andpolarizedtheonlyslightlypolarized.
that
the
flux
gle dependence of the linear polarization orLP(✓)
producedby ment
imal belt r0 and the asymmetry
parameter g sca for HIP 79977. The
L
/L
2
IR
⇤
are
then showused thefor68%
the andevaluation
offortheimage
fit.
Our best-fit model is in good agreement with results ofthe particle scattering needs to be taken into
account.
For
eximagecontours
95%
CLregionsthemodelsampleoriginates predominantly from scatterings with scattering anrecent observations of HIP 79977 with ALMA. Lieman-Sifry
ample, strict forward and backward scattering will produce and The
goodnessofthe
thelocations
fit was ofestimatedfor eachmodel
the dots
point out
the describedmodelsin thiswith
whereuncertainty
accounts
2
gles inrangethe45reduced
135
and the polarizedA&A
flux proofs:
produced
et al. (2016)
used
the
1240
µm continuumonly
visibilities
and de- for uncertainty in the
manuscript
no.theHIP_79977_arch
parameter
plain.
no polarization for randomly oriented particles
forthesymme-parameter:
rived basic
geometricalparametersof the disk. Theymodelled
try reasons. We adopt the Rayleigh scatteringfunction as a event can be approximated(F)/F-ratioderivedinthispaper. The double ratio ⇤per scatteringbyanaveragedparpol
disk
⇤
⇥
⇤
N
2
data
the surface density of the disk with a single power law r 1
X yi xi ( p)1
simple approximation for the angle dependence of the polar2 the induced
could
a goodthebetweenthe parameters
polarized oflu-twoticle parameterforscatteringpmresultingforthefor flux
extraction
for the intensity
extendingfrom bean inner
to an
outer
and
theyratio
derived=parameterse↵ects, 6 models
Table 3. Gridofself-subtraction
for
the
5.28
·polarization
10
andparametersfortheproxy
bestradius
fitformodel.
Also
given are thered
2ization fraction psca :+11
Ndataderived
N partheimage andRinner = 60
AU without detecting an outer cut-o↵ radius.
i=1intrinsic yiintensity slope for the diskf (✓, gdi↵ , gsca ) = w · f (✓, gdi↵ ) + (1deselected comparison
models70" and "Model 40").
the scattering
angle
of 90
(see("Model
Sect. 5.1).
along the major axis. In addition, they detected a polarimetric13
minosity
and the IR-excess luminosity of the disk

This result confirms the large ring radius r0 found by us from

where Nis a numberof data may
points bewithobtained
measurementif
data dust
New constraints
onproperties
1 cos2 ✓(Fpol )disk /F⇤ Article(Lnumber,results yi which have uncertainties yi . Each data point
psca (✓) = pm
= pm LP(✓),
page 15 of 20
pol )disk
1 + cos2 ✓quantitative polarimetric
data
of
many
debrisdisks
can ofbe linear⇤=⇡, Model 70 Model 40
corresponds
the minimization
Parameterto a binned pixel within
Range
Step
Best model
L
/L
L
IR
⇤
IR
with the scaling factor pm , which defines the
maximumfracwindow
shown
in
Fig.
9(d).
N
denotesthenumberofcollected. The dust grain size distribution
thepar and therefore
samplingmean value 68% CL
tional polarization produced at a scattering angle of ✓ = 90 . free parameters p = (p1 , p2 , ..., pN par ) used to create a model
polarimetric properties
areexpected
to dependonthe spectral if the wavelength and inclination dependence of the
Figure 8 shows some examples of obtained phase function image
withofvalueslisted in Col.
1 of90]
Table 3.
Radiusbeltxirand
[30,
10
73
16
70
40⇤ dust scat0 (AU)type oftheof central
star and di↵erent
system ages may reveal tering can be neglected. It is emphasized,
for the polarized flux LP(✓) f (✓, gsca ) for di↵erent
cases
the
that
neglecting
the
⇤
⇤
Inner
radial
index↵inprocedure[1,
10]
1
5.0
2.8
2.0
2.0
To accelerate
thepolarimetricfitting
wehave made aof
preseHG function f (✓, gsca ). For isotropic scattering(gsca = 0) theprocesses
evolutionary
in theproperties
thewavelength dependence of the polarized flux of the disk may
maximum of scattered polarized flux occurs at ✓ = 90 . For lection of disk models using the mean disk profile hP|x|i along
Outer
index
↵out10. The[-6,
0.5 not be an -2.5
1.4
-3.0
-2.5 for examscattering dust. thePolarimetricparameterswhich
can bequanFig. 9. Comparison of the best-fitanmodel
withparameter
the Q' gimage.
(a) Imasymmetry
majorradial
axis shownin Fig.
mean-1]
profilehP|x|iconacceptable simplification
for some cases,
sca > 0 the maximum is shifted
00
00⇤
tified for theare ofscatteringcross-sectionor albedo,
to smaller
scatteringa corresponding
asym-dust
sisting
15 points|x| = 0.22
to |x|= 1.80for theage visualizing the dust distribution
in thedisk.angles(b) producing
Model image
sca
Scale
heightH0from
(AU)
[0.5,
3.5]
0.5 ple, for near-IR
2.3
0.7
1.5 around 0.5
polarimetryofdisks
A-stars,which
metry within the amountof polarizedreceived
the front ofobserved
disk polarizationscattering
is obtained byaveraging p
the P(x)andparameters
the polarimetricfunctionof the polarized light non-convolved
PSF. (c)Modellightimage
of fromm andof their radiation
spectralregion.
and back sides of the disk. So, for example, the value of po- data
points fromthe negative and positive
Verticalprofile[0.5, x-axes2.5] given in Fig.
0.5 emit most0.9
0.6 in the UV-visual
1.0
1.0
the polarized light convolved withthe flux
instrumental
PSF.
imgsca. Ahigher
betterthe dust
in debris
larizedPF (gsca = 0.6)
at ✓ (d)= 20Qis' 35
times
thanunderstanding
6(c). Thus the 2Dofmodelswere collapsed
to adisks
profile would
and fitTherefore,
observationalthat depends,
2
Flare
[0.5,
2.2 ⇤ is an 1.4
2.5parameter3.5
at ✓area= 160outlined.
ted
first index
to the hP|x|i profile
calculating
the4.5]and defining
a1
age from the data. The rectangularwith an orange box
be very useful forinterpretations
regardingthenatureof their

0.5’’

61 AU

like the scattering albedo, on the ratio between dust scatter-shows the minimization windowArticle
as describedin the body text. The
Inclinationi ( ) the observed[82,dust
87] in a colli1 ing cross-section
84.6
1.7
85.0
82.0⇤
parent bodies which
producednumber, page 12 of 20
sca and absorption , and parameters of theorange circle marks the central region of the image excluded fromsional cascade. HG parameter g sca[0.0, 0.9]
0.1 polarimetric
0.43phase function
0.25 pm and
0.20
0.20
gsca asthe 2 evaluation. (e) Residual image obtained after subtraction of
Scalingfactor
A
9.04
4.03
3.10
p
the PSF-convolved model image (c) from the Q' image (d). Colorsca ( )⇤/· f (pm ( ), gsca ( ), i) .
(13)
6.2.1. Polarized flux and infrared excess
scales of images (a) and (b) are given in arbitrary units. The color-bar
for images (c), (d) and (e) shows polarized flux in counts per binned Notes: ⇤ Parameter value lies outside the 68% confidence interval.
The reflectivity or scattering albedo of the dust in the deThe inclination dependence of (Fpol )disk is illustrated inpixel.

bris disk of HIP 79977 can be characterized by a comparisonof the scattered polarized flux with the IR excess luminosity
which is a good measure for the dust absorption.

Fig. 13, which shows the polarized flux (Fpol )disk (expressed
per steradian) with respect to the scattered light luminosity
is a steeper
fall-o↵ thana Gaussian
distribution,
indicatLsca This
excludingthe di↵racted
light.
The scattered
light inter-