Star Formation Newsletter
No. 294 6-10
高橋実道（工学院大学）

1. Multiplicity and disks within the high-mass core NGC7538IRS1
Resolving cm line and continuum emission at ∼ 0.06′′ × 0.05′′ resolution
H. Beuther1 , H. Linz1 , Th. Henning1 , S. Feng2 , and R. Teague1

左：continuum, 白線は92年の結果
continuumに二つのピーク
Version
of
May
17,
2017
大質量星形成: 連星多い、円盤を通した降着
一つのジェットか二つのHII領域か？
ABSTRACT
大質量星形成領域は遠く観測的研究困難
Context. High-mass stars have a high degree of multiplicity and most likely form via disk accretion processes. The detailed
physics
23年移動無しof the binary and disk formation are still poorly constrained.
NGC7538の高分解能観測（VLA）
Aims. We seek to resolve the central substructures of the prototypical high-mass star-forming region NGC7538IRS1 at the highest

Max Planck Institute for Astronomy, Königstuhl 17, 69117 Heidelberg, Germany, e-mail: beuther@mpia.de
NGC7538:大質量星形成領域
Max Planck Institute for Extraterrestrial Physics, Giessenbachstrasse 1, 85748 Garching
12

possible spatial resolution line and continuum emission to investigate the protostellar environment and kinematics.
Methods. Using the Karl G. Jansky Very Large Array (VLA) in its most extended configuration at ∼24 GHz has allowed us to study
H. thermal
Beuther etCHal.:3Multiplicity
and disks
the high-mass
NGC7538IRS1
the NH3 andOH emission
and within
absorption
as wellcore
as thecm continuum emission at an unprecedented spatial resolutionof 0.06′′ × 0.05′′ , corresponding to a linear resolution of ∼150 AU at a distance of 2.7 kpc.
Results. A comparison of these new cm continuum data with previous VLA observations from 23 yrs ago reveals no recognizableproper motions. If the emission were caused by a protostellar jet, proper motion signatures should have been easily identified. In
combination with the high spectral indices S ∝ να (α between 1 and 2), this allows us to conclude that the continuum emission is fromtwo hypercompact Hii regions separated in projection by about 430 AU. The NH3 spectral line data reveal a common rotating envelope
indicating a bound high-mass binary system. In addition to this, the thermal CH3 OH data show two separate velocity gradients acrossthe two hypercompact Hii regions. This indicates two disk-like structures within the same rotating circumbinary envelope. Disk andenvelope structures are inclined by ∼33 ◦ , which can be explained by initially varying angular momentum distributions within the
natal, turbulent cloud.
Conclusions. Studying high-mass star formation at sub-0.1′′ resolution allows us to isolate multiple sources as well as to separate
circumbinary from disk-like rotating structures. These data show also the limitations in molecular line studies in investigating the disk
kinematics when the central source is already ionizing a hypercompact Hii region. Recombination line studies will be required forsources such as NGC7538IRS1 to investigate the gas kinematics even closer to the protostars.

cm continuum, NH3, CH3OHのライン

中央：Spectral index
ピーク付近で~２両方の結果が二つのHII領域であることを支持
連星系

Key words. Stars: formation – Stars: massive – Stars: individual: NGC7538IRS1 – Stars: rotation – Instrumentation: interferometers

1. Introduction
The formation processes leading to the most massive stars arestill puzzling in many ways. While there is a clear consensus thathigh-mass stars shape the interstellar medium, whole clusters,
and even entire galaxies, questions related to the physical processes at early evolutionary stages, in particular associated withthe fragmentation processes of young, dense, and rotating cores,& Sonnhalter 2002; Krumholz et al. 2009; Kuiper et al. 2010,2011; Kuiper & Yorke 2013; Peters et al. 2010; Commerçon
et al. 2011). Are massive disks similar to their low-mass counterparts, hence dominated by the protostar and Keplerian rotation or are they self-gravitating non-Keplerian entities (e.g.,Sánchez-Monge et al. 2013; Cesaroni et al. 2007)? The answer
may be a small inner Keplerian accretion disk that is fed froma larger-scale non-Keplerian structure (toroid or pseudo-disk).

２.

arXiv:1705.10167v1 [astro-ph.EP] 29 May 2017

Tidal dissipation in rotating low-mass stars
and implications for the orbitalE.evolutionofE↵ect
close-inBolmont et al.:of stellarmassive
metallicity
planets

We also investigated the influence of the metallicity for the
II. Effect of stellar metallicity
0.4 M star. The trend we discussed is much less visible than in
6 a 1 M
theY.casestar. For example, for planets beginning at
E. Bolmont1, 2 , F. Gallet3 , S. Mathis1 , C. Charbonnel3, 4 , L. Amard3, 5 , andAlibertof
0.03 au around initially fast rotating stars, we find that at an age
1
Laboratoire AIM Paris-Saclay, CEA/DRF - CNRS - Univ. Paris Diderot - IRFU/SAp, Centre de Saclay,
F-91191theyGif-sur-Yvette
of
5
Gyr,
are at 0.035
au, 0.036 au 小and 0.037
au around theHot
Jupiterの軌道分布に中心星のmetallicity依存性
metallicity星はコンパクトで潮汐の効果
小Cedex, France2
Z = 0.004, Z = 0.0134 and Z = 0.0255 stars respectively.
NaXys, Department of Mathematics, University of Namur, 8 Rempart de la Vierge, 5000 Namur, Belgium3星と惑星軌道の進化を同時に解く
潮汐による惑星軌道の変化で説明する星の進化が早い(星の自転周期)<(惑星軌道周期)

Department of Astronomy, University of Geneva, Chemin des Maillettes 51, 1290 Versoix, Switzerland
IRAP, UMR 5277, CNRS and Université de Toulouse, 14, av. E.Belin, F-31400 Toulouse, France
5
LUPM UMR 5299 CNRS/UM, Université de Montpellier, CC 72, F-34095 Montpellier Cedex 05, France
6
4.2. Effect of the initial stellar spin
Physikalisches Institut & Center for Space and Habitability, Universitaet Bern, 3012 Bern, Switzerland
4

Submitted to A&A

惑星は外側に移動

E. Bolmont et al.

In the previous section, we considered initially fast-rotating
1 of
M stellar
star. metallicity
We consider here models for the same initial mass
ABSTRACT
E. Bolmont et al.: E↵ect
7ofyr
<10
but
with
slowerinitialrotation(periods
3 and 8 days).
ObservationsofhotJupiter
type
exoplanetssuggestthat
theirorbitalperioddistribution
depends
onthemetallicity
of
their
host
star.
惑星軌道の進化
Figure6 shows
We investigate here whether the impact of the stellar metallicity on the evolution of the tidal dissipation inside
the convective
envelope the di↵erent evolutions for di↵erent initialWe
also
investigated
the influenceof themetallicity
for theof rotating stars and its resulting e↵ect on the planetary migration might be a possible explanation for this observed
statisticaltrend.
初期の星の自転周期：1日<
惑星軌道周期
spins.
For
close-inplanets
(a
=
0.03
au),
changingtheinitialinit
We use a frequency-averaged tidal dissipation formalism coupled to an orbital evolution code and to rotating stellar evolution models
0.4 Mspin
star.
The
trend
we
discussed
is muchFor
lessvisiblethan inin order to estimate the e↵ect of a change of stellar metallicity on the evolution of close-in planets. stellar
We consider
here twodi↵erent
leads
to
very
di↵erent
evolutions.
afastrotating
Z大ほど軌道大
惑星は外へ移動stellar masses: 0.4 and 1.0 M evolving from the early pre-main sequence phase up to the red giant branch.
the
case
of
a
1
M
star.
For
example,for planets
beginning
at
star,
the
planetsurvivesandmigratesoutwards.
WhileforslowerWe show that the metallicity of a star has a strong e↵ect on the stellar parameters which in turn strongly influence the tidal dissipation
0.03
aroundrotating
stars,thewecorotation
find that at
an agein the convective region. While on the pre-main sequence the dissipation of a metal poor Sun-like staris higherau
than
the dissipation
rotating
stars,
theinitiallyplanet isfast
initially
insideradiusmetal-poorの星は進化が早いため先に外へ移動
of a metal rich Sun-like star, on the main sequence it is the opposite. However, for the 0.4 M star, the dependenceof thedissipation
of
5
Gyr,
they
are
at
0.035
au,
0.036
au
and
0.037
au aroundand migrates inwards until it falls onto the star, or rather
until itthewith metallicity is much less visible.
Z =tidally
0.004,
Z migrate= 0.0134
andreaching
Z =! 0.0255
stars respectively.
Using an orbital evolution model, we show that changing the metallicity leads to di↵erent orbital evolutions(e.g., planets
gets
disrupted
when
the
Roche
limit.
farther out from an initially fast rotating metal rich star). By using this model, we qualitatively reproduced the observational trends
showsof the population of hot Jupiters with the metallicity of their host stars. However, more steps are needed to Figure
improve our6bmodel
to try the planet’s evolution for a stellar initial pe7 ~10
9theto quantitatively fit our results to the observations. Namely, we need to improve the treatment of the rotation evolution in the orbital
riod
of
3
days.
Initiallytheplanetsareinside10
yrcorotation raevolution model and ultimately we need to consistently couple of the orbital model to the stellar evolution
model.
4.2.
Effect of the initial stellar spin
dius but above the line corresponding to Porb = 1/2 P? , meaning
Key words. Planet-star interactions – Stars: evolution – Stars: rotation – Stars: abundances – Stars: solar-typeの方が軌道大that
migrate inwards
the dynamical
tide. fast-rotating
All planIn they
the previous
section,due
weto潮汐の影響が強いmetal-richconsidered
initiallyets
migraterapidlyhere
untilreaching
1/2 Pmass
1 M
star.inwards
We consider
modelsfor thethe Psameorb =initial?
1. Introduction
only the number of planets depends on the stellar metallicity but! Mathis (2016), once the planet
line.
As
discussedinBolmont&but
withslowerinitialrotation(periods of 3 and 8 days).
↑also the mass and radius of the planets, namelythat there areThanks to space missions such as CoRoT (CoRoT Team 2016), fewer small size/low-mass planets
reaches
thisline,
thethanequilibriumtide
takes over.
However,asaroundmetal-poorstars
9Figure6showsthedi↵erent
evolutionsfordi↵erent
initial>10yr
Kepler (Borucki et al. 2010), and K2 (Howell et al. 2014), around metal-rich stars. Furthermore, studies such as the works
thecorresponding
dissipation
is muchthe inward
spins.
For
close-in(afactorau),lower,changing
the initial
the amount of detected and confirmed exoplanets reaches an of Adibekyan et al. (2013) haveinit = 0.03
shownZ=Zsun
that
planets
betweenplanets
10
exquisite high number of objects belonging to an expended range M and 4 Mjup orbiting aroundmigration
slows
Meanwhilethestar continues
shrinkingmetal-poorstars
havedown.
ontoaverstellar
spin
leads
very
di↵erent
evolutions.
Forafast rotating星の進化で自転周期> 惑星軌道周期of star-planet system’s configurations. The high number of de- age a wider semi-major axis (longer orbital period) than thosespinning
up,survives
leading andto thedecreaseof the PWhileP?
orb = 1/2
star,
theplanetmigratesoutwards.
forslowertected exoplanets allows us to do some statistics over the char- orbiting metal-rich stars. Moreandspecifically, the last two aspects
acteristics of the host star. For instance, observational studies appear to be true for the hot Jupiterlimit.
The stars,
planetthe
entersonceisagain
the regionofexcitationof radiusthe
惑星は落ち始める
population.
rotating
planetinitially
insidethecorotationsuch as Gonzalez (1997), Santos et al. (2003) or Adibekyan et al.
inertial waves and the dissipation factor increases to the dynam-

３. In situ accretion of gaseous envelopes on to planetary cores
embedded in evolving protoplanetary discs
Gavin A. L. Coleman1⋆ , John C. B. Papaloizou2, Richard P. Nelson1,3
1 Astronomy Unit, Queen Mary University of London, Mile End Road, London, E1 4NS, U.K.
2 DAMTP, University of Cambridge, Wilberforce Road, Cambridge, CB3 0WA, U.K.

arXiv:1705.08147v1 [astro-ph.EP] 23 May 2017

Kavli Institute for Theoretical Physics, University of California, Santa Barbara, CA 93106, USA.
ガス惑星形成：その場形成か軌道移動の結果か？
3

その場形成でできる惑星の位置、質量を調べる
ABSTRACT

円盤進化とその中での惑星へのガス降着を同時に解く
The core accretion hypothesis posits that planets with significant gaseous envelopes
accreted them from their protoplanetary discs after the formation of rocky/icy cores.
Observations indicate that such exoplanets exist at a broad range of orbital radii, but
it is not known whether they accreted their envelopes in situ, or originated elsewhere
and migrated to their current locations. We consider the evolution of solid cores embedded in evolving viscous discs that undergo gaseous envelope accretion in situ withorbital radii in the range 0.1-10 au. Additionally, we determine the long-term evolution of the planets that had no runaway gas accretion phase after disc dispersal. We
find: (i) Planets with 5 M⊕ cores never undergo runaway accretion. The most massive
envelope contained 2.8 M⊕ with the planet orbiting at 10 au. (ii) Accretion is more efficient onto 10 M⊕ and 15 M⊕ cores. For orbital radii ap ≥ 0.5 au, 15 M⊕ cores always
experienced runaway gas accretion. For ap ≥ 5 au, all but one of the 10 M⊕ cores experienced runaway gas accretion. No planets experienced runaway growth at ap = 0.1 au.
(iii) We find that, after disc dispersal, planets with significant gaseous envelopes cooland contract on Gyr time-scales, the contraction time being sensitive to the opacity
assumed. Our results indicate that Hot Jupiters with core masses ! 15 M⊕ at ! 0.1 au
likely accreted their gaseous envelopes at larger distances and migrated inwards. Consistently with the known exoplanet population, Super-Earths and mini-Neptunes at
small radii during the disc lifetime, accrete only modest gaseous envelopes.

In situ accretion of gaseous envelopes

3

10

2

10

1

10

0

10

Key words: planets and satellites: atmospheres, formation, gaseous planets – planet–
disc interactions – protoplanetary discs

−1

10

−2

10
1

−1

0

10
10
INTRODUCTION

1

10

Two distinct paradigms have been proposed to explain
the origin of gas giant planets: core nucleated instability (Perri & Cameron 1974; Mizuno et al. 1978; Stevenson
1982) and gravitational instability (e.g. Helled et al. 2014,
and references therein). In principle, given the presence ofa solid core of suﬃcient mass, core nucleated instability canlead to the formation of gas giant planets at almost any
orbital location within a protoplanetary disc during its life-

In this paper, we examine the accretion of gas on toalready formed cores located at a wide variety of stellocentric distances, ranging between 0.1 au and 10 au within
evolving protoplanetary disc models, to determine wheregas giant planets can form via in situ gas accretion, andthe likely envelope masses of planets that do not achieve
runaway growth. Most previous studies of gas accretionon to planetary cores have used either static models ofplanetary envelopes, where there is a thermal equilibrium for which planetesimal accretion on to the core pro-

5

8

G. A. L. Coleman et al

µδ (

0

µδ (

4. Multiple kinematical populations in Vela OB2 from Gaia DR1 data
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F. Damiani1 , L. Prisinzano1 , R. D. Je↵ries2 , G. G. Sacco3 , S. Randich3 , and G. Micela1

INAF Osservatorio
Astronomico di Palermo G.S. Vaiana, Piazza del Parlamento 1, 90134 Palermo, Italy
Vela OB2
associatione-mail: damiani@astropa.inaf.it
Astrophysics Group, Keele University, Keele, Sta↵ordshire ST5 5BG, UK
2
γ Vel：
Radial
Velocity
で二つのグループ（A,B）
INAF Osservatorio
Astrofisico
di Arcetri, Largo
E. Fermi 5, 50125 Firenze, Italy
1

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2
3

0

1

2

3

4

5

Received 24 March 2017 / Accepted 10 May 2017

NGC25476

π (mas)

−40

−30

−20

−10

0

10

20

µα (mas/yr)

30

40

Fig. 1. Proper motion along Declination µ vs. parallax ⇡ for the ABSTRACT
en- Fig. 3. As in Fig. 2, but using data from the external regions in Vela OB2tire Vela OB2 TGAS sample. Vertical lines correspond to literature dis- (772 stars).
2固有運動で見ると？
Recenttancesfor Context.
Vel (red)
andresults
NGC using
2547radial-velocity(green), and measurements
to our samplefrom
se- the Gaia-ESO Survey have led to the discovery of multiple kinematic
populations
acrosserrors
the Velaassociation. We present
herearea proper-motion study of the same region.
lection (blue).
The average
forOB2the parallax-selected
sampleOur aim is to test
whether or not the radial-velocity populations
have alocalised
counterpartoverdensities
in proper-motion
and if so,against
how the a much
RadialAims.
Velocityのグループとの関係は？
shown.
of twoofspace,
datapoints

two sets of kinematical data complement each other.
broader distribution of points (field stars). We therefore deMethods. This work is based on parallaxes and proper motions from the TGAS catalogue, as part of Gaia DR1.
fineacrossfrom⇠5thisdiagram(andfrompreviousparallax
Results. Two distinct proper-motion populations are found dispersed
degrees(or ⇠30
pc at
theirthelikelydistances).
Their limits)
two
populations,termedhereofCcommon( 9 <objects.
µ↵ <However,
3 mas/yr anddetailed correspondence to the radial-velocity populations could not
be testedbecause of
the paucity
7.5populations< µ < consists
10.5 mas/yr),
D ( of13the<youngµ↵ cluster< 7 mas/yr
compelling indications are found that one of the new proper-motion
mostly of andmembers
NGC 2547, and the other is related to the 2 Vel cluster. Constraints
the ageof thebothuseof whichappearbeandon2.5< µ< two
5.5populations,mas/yr). We
C and
D toto avoid
cononly 10 35 Myr old, and their possible mutual interactions within thelast 1.5
Myrtheare twodiscussed.
fusionwith
populations (A and B) defined by Je↵ries←Gaia で得た固有運動の分布
二つのグループ (C, D)

10

1. Introduction−100

The Vela OB2 association spans a wide sky region inside theGum nebula (Pettersson 2008) and has been subject to an increasing number of studies in recent years. Its visually brightest
member, the O8 III+Wolf-Rayet binary star 2 Vel, was foundto be surrounded by a cluster of lower-mass stars through X-ray
observations (Pozzo et al. 2000). The 2 Vel cluster was thenstudied, among others, by Je↵ries et al. (2009), and more recently by Je↵ries et al. (2014, 2017), as part of the Gaia-ESO
survey (Randich et al. 2013; Gilmore et al. 2012). This survey
providedprecisedata for 0hundreds10of candidate−40−30 radial-velocity
−20
−10
20
members of the 2 Vel cluster,
enabling Je↵ries et al. (2014)µα (mas/yr)
to find clear evidence of a double kinematical population, sep-−20

µδ (mas/yr)

20

al. (2014);
Key words. open clusters and associations: individual: Vel OB2 –etpropermotionswe discuss below the possible connection between

Radial
VelocityのグループABとのthese pairs. Population C (D) comprises 148 (101) stars. We
havetried to estimate the field-star contamination in these sam直接比較はない（共通の天体が少ない）
ples by considering TGAS stars, within the same parallax lim-2
similarinage
andkinematics
to
group
B
of
the
younger
its, falling in a rectangular region of the (µ↵ , µ ) diagramVelintercluster. Unfortunately, no additional Gaia-ESO Survey
observamediate
between those enclosing populations C 2andVel)の固有運動と近い値
D, that is
Cの固有運動はグループAB(γtions exist towards Vela OB2, most of which remains unexplored13 < µ↵ < 3 because
mas/yr ofand
5.6 < µ < 7.4 mas/yr; 53 stars
spectroscopicallyits large size.
are found.
This region has the same area in the (µ↵ , µ stud) plane
The richness of kinematical signatures found in theseasenclosingpopulations
C andD, so noareathecorrections
iesthose
motivated
us to searchfor similar
evidence
among
newly
are
required.
the expectedlevel ofcontamination
released
GaiaTherefore,
DR1 data (Gaia
Collaboration
2016a,b),the so- is⇠36%
population
C,Sinceand ⇠52%
for catalogue
populationonD.
called for
TGAS
catalogue.
the Tycho
which theToDR1
helpisunderstandthe spatial
regionGaia
based only reaches
down
to V ⇠where
12, wePopulations
do not ex- C
andwemajority
have repeatedthe same stars
selection
TGAS stars,
pectDtoexist,
find theof the individualin theofkinematicalbut
now fromfrom
an annulus
surrounding
the Saccoprevious
region,of 4
populations
Je↵ries etal. (2014) andet al.
(2015),
since
they (outer)are mostly
fainterthan this(11 <theV proper-motion< 19). However,dia(6) innerradius,which
yieldedas we(µdescribe
below, intheFig.
Gaia
showpeaks
clear aregram
3. proper-motion
Here, the two data
density↵ , µ ) shown

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F. Damiani et al.: Kinematical populations in Vel OB2 from Gaia DR1 data

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126

124

122

120

118

116

RA (deg)

Fig. 4. Spatial distributions of kinematical populations C and D (red andgreen, respectively). The small black dots indicate the datapoints outside of rectangles in Fig. 2. The big circles indicate positions of 2 Vel(north) and NGC 2547 (south). The dashed arrows indicate their inferred positions 1.1 Myr ago.

( 4.6 ± 1.0, +8.7 ± 0.9) mas/yr for populations A and B, respectively). We therefore tentatively associate population C with the2
Vel populations A and B cumulatively, and population D with
NGC 2547. Figure 4 also shows that NGC 2547 is not the geometrical centroid of either population C or D, which are spread

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5. The Gaia-ESO Survey: dynamics of ionized and neutral gas in the
A&A
proofs: manuscript
no. Damiani_Lagoon
Lagoonnebula(M8) ?
F. Damiani1 , R. Bonito1, 2 , L. Prisinzano1 , T. Zwitter3 , A. Bayo4 , V. Kalari5 , F. M. Jiménez-Esteban6 , M. T. Costado7 ,P. Jofré8, 9 , S. Randich10 , E. Flaccomio1 , A. C. Lanzafame11 , C. Lardo12 , L. Morbidelli10 , and S. Zaggia13

v:1705.08194v1 [astro-ph.GA] 23 May 2017

Dec (deg)

1

INAF - Osservatorio Astronomico di Palermo G.S.Vaiana, Piazza del Parlamento 1, I-90134 Palermo, Italye-mail: damiani@astropa.inaf.it
2
Dipartimento di Fisica e Chimica, Università di Palermo, Piazza del Parlamento 1, 90134, Palermo, Italy
3
Faculty of Mathematics and Physics, University of Ljubljana, Jadranska 19, 1000, Ljubljana, Slovenia
4
Instituto de Física y Astronomiía, Universidad de Valparaiíso, Chile
5
Departamento de Astronomía, Universidad de Chile, Casilla 36-D Santiago, Chile
6
Departmento de Astrofísica, Centro de Astrobiología (INTA-CSIC), ESAC Campus, Camino Bajo del Castillo s/n, E-28692 Villanueva de la Cañada, Madrid, Spain
7
Instituto de Astrofísica de Andalucía-CSIC, Apdo. 3004, 18080 Granada, Spain
8
Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, United Kingdom
9
Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, Av. Ejército 441, Santiago, Chile
10
INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125, Florence, Italy
11
Dipartimento di Fisica e Astronomia, Sezione Astrofisica, Università di Catania, via S. Sofia 78, 95123, Catania, Italy
12
Laboratoire d’astrophysique, Ecole Polytechnique Fédérale de Lausanne (EPFL), Observatoire de Sauverny, CH-1290 Versoix,
Switzerland
13
INAF - Padova Observatory, Vicolo dell’Osservatorio 5, 35122 Padova, Italy
Received date / Accepted date
ABSTRACT
Aims. We present a spectroscopic study of the dynamics of the ionized and neutral gas throughout the Lagoon nebula (M8), using

VLT/FLAMES data from the Gaia-ESO Survey. The new data permit to explore the physical connections between the nebular gas
and the stellar population of the associated star cluster NGC6530.
Methods. We characterize through spectral fitting emission lines of H↵, [N II] and [S II] doublets, [O III], and absorption lines of
sodium D doublet, using data from the FLAMES/Gira↵e and UVES spectrographs, on more than 1000 sightlines towards the entire
face of the Lagoon nebula. Gas temperatures are derived from line-width comparisons, densities from the [S II] doublet ratio, andionization parameter from H↵/[N II] ratio. Although doubly-peaked emission profiles are rarely found, line asymmetries often imply
multiple velocity components along the same line of sight. This is especially true for the sodium absorption, and for the [O III] lines.
Results. Spatial maps for density and ionization are derived, and compared to other known properties of the nebula and of its massive
stars 9 Sgr, Herschel 36 and HD 165052 which are confirmed to provide most of the ionizing flux. The detailed velocity fields
across the nebula show several expanding shells, related to the cluster NGC6530, the O stars 9 Sgr and Herschel 36, and the massive
protostar M8East-IR. The origins of kinematical expansion and ionization of the NGC6530 shell appear to be di↵erent. We are able
to put constrains on the line-of-sight (relative or absolute) distances between some of these objects and the molecular cloud. The
data show that the large obscuring band running through the middle of the nebula is being compressed by both sides, which might
explain its enhanced density. We also find an unexplained large-scale velocity gradient across the entire nebula. At larger distances,
the transition from ionized to neutral gas is studied using the sodium lines.
RA (deg)

Key words. ISM: individual objects: (Lagoon nebula) – ISM: general – HII regions