exploring the most massive stars beyond the local group the … · 2020-07-13 · if we take into...

University of Amsterdam

Faculty of Science

Astronomical Institute ‘Anton Pannekoek’

Master’s ThesisAstronomy & Astrophysics

Exploring the most massive stars beyond the Local Group

The first intermediate-resolutionspectrum of a massive star candidate

in NGC 55

key words: extra-galactic spectroscopy, massive stars, H ii regions

Author:

Olga [email protected]

5 November 2010

Supervisors:

Dr. Hugues [email protected]

Prof. Dr. Alex de [email protected]

Second corrector:

Prof. Dr. Lex [email protected]

Master coordinator:

Prof. Dr. Ralph [email protected]

Abstract

As part of the X-Shooter science verification, we obtained the first intermediate-resolution spectrumof a massive star candidate in NGC 55, an LMC like galaxy at a distance of 2.0 Mpc. Based onthe morphological analysis of the UBV (300− 550 nm) and VIS (550− 1000 nm) spectrum and onthe modelling of the nebular and atmospheric lines, we investigated the nature and evolutionarystatus of the central object and its influence on the surrounding interstellar medium. We concludethat the source, NGC 55 C1 31, is most likely a small cluster which contains at least one hot (Teff =40000 K), probably evolved, star with a high mass loss rate (2.0×10−5 M� yr−1) and a helium-richcomposition (X = 0.237, Y = 0.757). Around λ = 500 nm, this star provides 10% of the luminosityof this cluster. The remaining light is produced by objects with a luminosity-averaged effectivetemperature of 30000 K, solar abundances (X = 0.710, Y = 0.284), and lower mass loss fluxes8.9 × 10−9 M� yr−1 R−2

� . The surrounding region that is, most probably mainly, ionised by thiscluster has an electron density ne ≤ 20 cm−3 and an electron temperature Te = 11200 K. To bringthe conclusions on the cluster composition (the ionising spectral energy distribution) in agreementwith the properties of the nebula, we need to assume a metallicity of Z = 0.2 Z�, lower than earliermeasurements of this region.Despite the low signal-to-noise ratio of the resultant spectrum and the problems with the nebularcontamination, we have shown that it is possible to do spectroscopy on massive star candidates inNGC 55 with X-Shooter. This project provided important insights for follow-up observations.

Popular summary

Zware sterren zijn ontzettend zeldzaam. Dat wil zeggen, voor iedere ster die twintig keer zo zwaaris als de zon zijn er honderdduizend ‘lichte’ sterren van ongeveer een zonsmassa. De zwaarstesterren zijn zo zeldzaam dat er maar een paar van bekend zijn. Zware sterren hebben een veelkorter en heftiger leven dan bijvoorbeeld onze zon. Hoewel ze zwaarder zijn en dus meer brandstofhebben, maken ze deze onevenredig veel sneller op. Gedurende hun leven zijn ze vaak veel heteren lichtkrachtiger. Ook verliezen ze veel massa door middel van hun sterke sterwind, die aan heteind van hun leven nog veel sterker wordt. Ondanks hun kleine aantal hebben deze sterren toch eenbepalend effect op hun omgeving. Verreweg de meeste ioniserende UV straling van melkwegstelselskomt van zware sterren. Met hun sterwinden en de supernovae aan het eind van hun leven verrijkenze het interstellaire medium met elementen zwaarder dan helium: ‘metalen’.

De reden dat we naar sterren buiten ons Melkwegstelsel of zelfs buiten de Lokale Groep kijken,is dat het ons iets zal vertellen over zware sterren in een omgeving met meer of juist minder vandeze metalen. Deze ‘metalliciteit’ blijkt een erg grote invloed te hebben op bijvoorbeeld sterwindenen de vorming van (zware) sterren. Bovendien was in het vroege heelal het aandeel zware elementenlager. Het bestuderen van sterren in een omgeving arm aan metalen zal ons dus wellicht meer lerenover de vorming en evolutie van de eerste sterren.

In dit project analyseren we het spectrum van een kandidaat zware ster in een melkwegstelselbuiten onze Lokale Groep, NGC 55, dat een relatief lage metalliciteit zou moeten hebben. Debuitengewoon gevoelige spectrograaf X-Shooter, sinds kort operationeel op de Very Large Telescopein Chili, stelt ons in staat het spectrum op te nemen met een groter golflengtebereik en hogerespectrale resolutie dan eerder voor dit verre object gedaan is.

Het spectrum van een ster kan veel onthullen over haar eigenschappen. Maar het spectrum vandit object was niet in overeenstemming te brengen met een soort ster. En de bron was ook welerg helder voor slechts een object. Het is daarom zeer waarschijnlijk dat we een klein cluster vansterren hebben waargenomen. Analyse van het spectrum met behulp van modellen wees uit dat erin ieder geval een geevolueerde, hete ster met een hoog massaverlies bij zit. Bestudering van deeigenschappen van de omringende geıoniseerde ijle materie die deze cluster omringt doet vermoedendat de metalliciteit misschien nog wel lager is dan uit eerdere metingen bleek, slechts 1/5 van desolaire waarde.

Contents

I Introduction 9

1 Massive stars and their environments 101.1 Evolution of massive stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 111.2 Stellar winds . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 121.3 H ii regions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 14

2 Host Galaxy NGC55 and source 162.1 Distance determinations of NGC 55 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 172.2 Metallicity determinations of NGC 55 . . . . . . . . . . . . . . . . . . . . . . . . . . 182.3 Source NGC 55 C1 31 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

3 This project 23

II Observations and Data Reduction 25

4 Instrument X-Shooter 264.1 Specifications . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 264.2 Technical description . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28

5 Observations 30

6 Data reduction 336.1 Data reduction cascade . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

6.1.1 Determining the spectral format . . . . . . . . . . . . . . . . . . . . . . . . . 346.1.2 Measurement and subtraction of detector bias and dark current levels . . . . 346.1.3 Bad pixels and detector linearity . . . . . . . . . . . . . . . . . . . . . . . . . 356.1.4 Guess solutions for order centre location and wavelength calibration . . . . . 366.1.5 Tracing of echelle orders . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 376.1.6 Flat field correction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 376.1.7 Two dimensional mapping and wavelength calibration . . . . . . . . . . . . . 386.1.8 Computing efficiency and response, flux calibration . . . . . . . . . . . . . . . 396.1.9 Telluric correction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 396.1.10 Processing science frames . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40

6.2 1D extraction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40

III Observational Results 45

7 2D spectra 467.1 2D-1D correspondence . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 467.2 Other objects . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46

CONTENTS 7

8 1D spectra 518.1 Object spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

8.1.1 Comparison with other observations . . . . . . . . . . . . . . . . . . . . . . . 558.2 Nebular spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

IV Scientific Analysis 62

9 Central star NGC55 C1 31 639.1 Luminosity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 639.2 Modeling line profiles: Teff, M and vrot sin(i) . . . . . . . . . . . . . . . . . . . . . . 64

9.2.1 Constraints on effective temperature . . . . . . . . . . . . . . . . . . . . . . . 659.2.2 Constraints on rotational velocity . . . . . . . . . . . . . . . . . . . . . . . . . 669.2.3 Constraints on mass loss . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 709.2.4 Accordance with observed magnitude . . . . . . . . . . . . . . . . . . . . . . 709.2.5 He iiλ4686 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 729.2.6 Multiple sources producing blended lines. . . . . . . . . . . . . . . . . . . . . 81

9.3 Conclusion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85

10 Circumstellar nebula 8610.1 Phenomenological description of nebular spectrum . . . . . . . . . . . . . . . . . . . 8610.2 Properties of the surrounding region . . . . . . . . . . . . . . . . . . . . . . . . . . . 87

10.2.1 Electron density . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8710.2.2 Electron temperature . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88

11 A consistent picture 9211.1 Procedure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9211.2 Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9311.3 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

V Discussion and Conclusions 100

12 Discussion 10112.1 A comparison with former studies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101

12.1.1 Central source . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10112.1.2 Surrounding nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

12.2 Future investigation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

13 Conclusions 104

Bibliography 104

A FASTWIND model grid 109

B CLOUDY abundances 112

Part I

Introduction

1Massive stars and their environments

The most determining factor in stellar evolution is the initial mass of a star. Stars belonging todifferent classes of mass evolve on different time-scales, are subject to different internal nuclear pro-cesses, and leave different final products. Massive stars (M ≥ 9 M�) are characterised by relativelyshort and turbulent lives, that end with the formation of a neutron star or a black hole duringextremely energetic explosions which we observe as supernovae or gamma-ray-bursts.

Massive stars are rare. Initial mass functions (e.g. Salpeter 1955) predict that far more low massstars than high mass stars are formed. If we take into account the shorter lifetimes of massive stars,we expect to see about a hundred thousand solar-type stars for every 20 M� star in the Galaxy.For every 100 M� star, there should be over a million solar type stars (Massey, 2003). Despitethe fact that they are not numerous, massive stars are key players in the universe. Their influenceextends well beyond their immediate surroundings.

� By far the most ultra-violet radiation from galaxies comes from massive stars. Their highluminosities make them visible even at very large distances. The first massive stars are heldresponsible for the re-ionisation of the early universe (Haehnelt et al., 2001). Because of thehigh numbers of high-energy photons they emit, massive stars are able to ionise large gasclouds in galaxies. Via the heating of dust, massive stars also contribute strongly to the farinfra-red luminosities of galaxies.

� Massive stars provide most of the input of mechanical energy, momentum and mass in theinterstellar medium (Abbott, 1982); during their lives via their strong stellar winds and evenmore when they end their life as a supernova. Stellar winds and supernovae may createpressure waves in the interstellar gas, that can initiate the formation of new generations ofstars. Also, massive stars can bring star formation to a halt by blowing away away the materialin the star forming environments in which they reside.

� The bulk of elements heavier than helium in the universe originate from fusion in stars andsupernova explosions. The stellar outflows and supernova ejecta of massive stars contributestrongly to the chemical enrichment of the interstellar medium.

� Massive stars play a profound role as the source population for exotic compact objects likeneutron stars and black holes, and the high energy phenomena and processes associated withthose objects.

The reason that we study massive stars outside our Galaxy, and even in galaxies beyond theLocal Group is that the metallicity of the gas out of which these stars are born differs by a factorof 20 amongst the galaxies that are currently forming stars. Metallicity appears to be an importantfactor in the formation and evolution of stars. For instance, the temperature of star-forming clouds

1.1. EVOLUTION OF MASSIVE STARS 11

depends on metallicity. This affects the Jeans length, the critical radius for a molecular cloudto collapse and start to form stars. Therefore the initial mass function depends upon metallicity.Furthermore, stellar winds of massive stars are driven by radiation pressure in highly ionised metallines (see also section 1.2), and are thus very sensitive to metal abundances. Mass loss processesare dominant in the evolution of massive stars. Depending on the metallicity of the host galaxy, wethus expect to see large differences in the evolved massive star population (Massey, 2003).

1.1 Evolution of massive stars

Spectral classification of stars is based on the morphology of the spectrum. The temperature anddensity of a stars atmosphere mainly determines its ionisation and what atomic transitions are mostprominent. The spectral type reflects the surface temperature. From hot to cold the spectral classesof the Morgan and Keenan (MK) system are O, B, A, F, G, K, M, L and T. For historical reasons,hot O, B and A-type stars are referred to as ‘early types’. Within one spectral class, a subclasscan be indicated with a number: 0 for the hottest (‘early’) type, to 9 for the (‘late’) coolest. Thespectral type is defined by the relative strength of spectral lines.

The other dimension in this system is the luminosity class. This is based on the pressure broad-ening of lines, and reflects the size (luminosity) of stars of similar temperature. Luminosity classI is the brightest and labels the supergiants. In addition, there are II (bright giants), III (giants),IV (sub-giants) and V (main-seqeunce stars or dwarfs).

During the lifetime of a star, its properties, and thus the spectral type and luminosity classchange. With effective temperatures between ∼ 25000 and ∼ 50000 K, massive stars are born asO-type or early B-type main sequence stars. Their luminosities are between twenty thousand andone million solar luminosities. While on the main sequence, the hydrogen in the core is fused viathe CNO cycle. If the star does not rotate very fast, the radius and luminosity increases graduallywhile the surface cools. Except for the most massive stars (Minit & 60M�), the stars have evolvedinto B-type supergiants when the core hydrogen is exhausted.

The post main-sequence evolution depends strongly on mass and mass loss. If the initial massof the star does not exceed ∼ 20M�, the star becomes a red supergiant (RSG). As an effect ofhydrogen shell burning, the star expands and its surface temperature decreases. Red supergiants arethe largest stars known: their radii are typically 200 to 800 solar radii. Red supergiants can becomeblue supergiants when other, slower, fusion reactions take over, that cause the star to contract.RSGs have slow, dense winds, while those of BSGs are fast but sparse. The star is likely to end itslife in a supernova.

Stars with 20 . Minit . 40M� also become red supergiants when the core hydrogen burningstops. Because of the high mass loss they suffer from in this stage, they will at some point haveexpelled their entire envelope. We observe deeper, hotter layers that are strongly enriched withH-burning products: a so-called Wolf-Rayet star. Wolf-Rayet stars are characterised by a hotcontinuum, with broad emission lines form highly ionised elements. These lines are formed inthe strong, optically thick stellar wind resulting from the high mass loss. Three main types aredistinguished: WN, WC and WO stars, referring to the strongest lines in their spectra: ionisednitrogen, carbon and oxygen. A star with an initial mass & 25 M� probably suffers from enoughmass loss to evolve through the WC phase, where it reveals more advanced He-burning products atits surface.

Stars with Minit & 40 M� move to the Wolf-Rayet phases without becoming a cool objectfirst: the outer layers already become unstable before they expand too much. In very massive stars(Minit & 85 M�) this is expected to lead to severe mass loss events that characterise the Luminous

12 CHAPTER 1. MASSIVE STARS AND THEIR ENVIRONMENTS

Table 1.1: Modern version (Massey, 2003) of the “Conti scenario” (1976) for the evolution of massive stars.

Initial mass range Evolution

10 M� < Minit < 20 M� OB→RSG→BSG→SN20 M� < Minit < 25 M� O→RSG→WN→SN25 M� < Minit < 40 M� O→RSG→WN→WC→SN40 M� < Minit < 85 M� O→WN→WC→SN85 M� < Minit O→LBV→WN→WC→SN

Blue Variable (LBV) phase. LBVs have extreme bolometric luminosities, and are found near theedge of the observed upper luminosity limit in the H-R diagram. They undergo episodic periods ofhigh mass loss, during which they brighten visually, while remaining roughly constant in bolometricluminosity (Massey, 2003). These massive stars become Wolf-Rayet stars after the LBV phase. TheWolf-Rayet phase is the last phase before the star collapses in a supernova.

The sketched scenario is the current paradigm that is valid for non-rotating or modestly rotatingobjects. Different scenarios unfold when taking into account effects of rotation (see e.g. Maeder &Meynet 2000). Rotation induces mixing processes in the stellar interior that increase the luminosityand the amount of stellar material available for fusion. This can drastically change the evolutionarypaths (see e.g. Yoon & Langer 2005). For stars rotating at 30 to 60 percent of their break-up speedthe rotationally induced mixing is so efficient that they remain chemically homogeneous during theirevolution. Models predict that these stars move blue-wards from the zero-age main sequence andupwards to higher luminosities, instead of cooling and expanding first. The physics of rotation itselfis complicated, but we should also be aware of the interplay between rotation and mass loss, theeffect that governs most of the evolutionary processes of massive stars. Fast rotation may lead toenhanced mass loss, while the loss of angular momentum by the stellar wind can slow down therotation (Langer, 1998).

1.2 Stellar winds

Stars do not only emit electromagnetic radiation but also particles. The continuous spherical out-flow of material from a star is called a stellar wind. Almost all kinds of stars have stellar winds, butthere are differences in driving mechanisms and strengths between different types of stars. Espe-cially in the evolution of massive stars, mass loss mechanisms play a major role. Before it ends itslife in a supernova explosion, a massive star may have ejected half of its initial mass by stellar winds.

Two important observable parameters of a stellar wind are the mass loss rate M and the ter-minal wind velocity v∞. The mass loss rate, expressed in solar masses per year, defines how muchmaterial is lost by the star per unit time. The terminal wind velocity is the velocity that theout-flowing material reaches at a large distance from the star. For understanding stellar evolution,measuring values of M and v∞ is important. Different mechanisms of mass loss predict differentbehaviour for these parameters. Measuring them helps to determine what processes are responsiblefor mass loss. Additionally, M and v∞ are required to quantify the energy feedback by stellar

1.2. STELLAR WINDS 13

winds to the interstellar medium, an important effect mentioned in section 1. The amount of ki-netic energy per unit time that is deposited into the interstellar medium by a stellar wind is 0.5Mv2

∞.

The equation of mass continuity

M = 4πr2ρ(r)v(r)�� 1.1

relates the mass loss rate M to the density ρ(r) and the velocity v(r) of the material at any point inthe wind at a distance r to the centre of the star. This equation implies that the wind is sphericallysymmetric and that no material is created or destroyed in the wind. The velocity profile of the windcan be approximated by a β-law, often represented in the following form:

v(r) ' v∞

(1− r0

r

)β �� 1.2

with

r0 = R∗

{1−

(v0

v∞

)1/β} �� 1.3

in which R∗ is the stellar radius and v0 is the initial velocity of the out-flowing material. Theparameter β is the acceleration of the out-flowing material, i.e. the slope of the velocity profile.Velocity profiles with different values of β are shown in figure 1.1. In the stellar wind of a hot star,the wind accelerates fast: values of β = 0.8−1.0 are typical. Cooler stars have winds that acceleratemore slowly, corresponding to higher values of β (Lamers & Cassinelli, 1999). In this project weadopt a default value of β = 1.0 (see section 9.2).

In a stellar atmosphere, photons from the star exert a radiation force on particles. When theradiation pressure and the gas pressure overcome gravity, gas streams out. The gas in the atmo-sphere is ionised and consists therefore of free electrons and ions of hydrogen, helium and metals.The radiation field can transfer its energy and momentum to the stellar wind by photon interactionswith free electrons or ions. These interactions can be divided into continuum and line interactionprocesses. The continuum processes are collisions of photons with free electron or with ions. Theexcess energy of the photon appears as kinetic energy of particle with which it collided. The numberof photons and the density of the medium determine the continuum radiation force. In dense regionsdeep in the wind, the continuum radiation force may play an important role. However, in the layersfurther out the line interactions will be dominant.

When a photon has exactly the energy of a transition in the ion with which it collides, a lineinteraction can take place. The absorption of a photon excites the ion. When it de-excites, a newphoton in emitted in an arbitrary direction, implying that the nett force on the ions is directedoutwards. The number of photons that can be absorbed is high for complex ions that have a highnumber of possible transitions. Fe iii, for example, is only a factor 1/25000 as abundant as hydro-gen in our galaxy. But since it has millions of spectral lines, Fe iii is a very important wind-drivingion. Ionised hydrogen only has a few hundred possible line transitions. The dependence of M onZ, the metallicity, is predicted to be a power law, such that M ∝ Z0.5−0.7 (see e.g. Vink et al., 2001).

The terminal wind velocity of line-driven winds depends on the escape velocity at the surfaceand the effective temperature of the star. The effective escape velocity is given by

vesc =√

2(1− Γe)GM∗/R∗�� 1.4

14 CHAPTER 1. MASSIVE STARS AND THEIR ENVIRONMENTS

in which

Γe =σeL∗

4πcGM∗

�� 1.5

corrects the Newtonian escape velocity for radiation pressure due to electron scattering. L∗ and M∗are respectively the luminosity and the mass of the star; c is the speed of light, G is the gravitationalconstant and σe is the electron scattering coefficient per unit mass. It has a value σe ' 0.30 cm2 g−1for winds of hot stars. The ratio vesc/v∞ is about 2.6 for stars with Teff > 21000 K (Lamers &Cassinelli, 1999).

1.3 H ii regions

Diffuse nebulae, also known as H ii regions, are a class of gaseous nebulae. These are regions ofinterstellar gas that are ionised and illuminated by O or early B-type stars. Because these hot andluminous stars formed relatively recently, the interstellar material that is not used to form thesestars is still present. Very young O and B stars are often surrounded by so-called ultra compactH ii regions. These regions are so dense that the massive star in formation can only be observedin the near-infra-red. After ∼ 104 yr, the region has expanded enough for the central star to be-come visible. Classical H ii regions are expected to represent a more mature phase in this process(Garcıa-Segura & Franco, 2004).

Gaseous nebulae are observed as bright extended objects, that emit strong emission lines super-imposed on a weak continuum. The emission line spectrum is dominated by forbidden lines fromions of common elements such as [O ii], [O iii], [N ii] and [S iii]. Furthermore, we observe recombina-

0

0.2

0.4

0.6

0.8

1

1 1.5 2 2.5 3 3.5 4 4.5 5

v(r

)/v

∞

r/R*

β-velocity laws for different values of β

β = 0.5β = 0.8β = 1.0β = 2.0β = 3.0

Figure 1.1: Stellar wind velocity as a function of distance from the star (equation 1.2) for different values of β, therate of acceleration of the outflow. For the initial velocity we assume v0 = 0.01v∞, which is close to the sound speed.

1.3. H II REGIONS 15

tion lines of hydrogen (Hα, Hβ etc.) and singly ionised helium (i.e. He iλ5876). The continuum ofa gaseous nebula consists of atomic and reflection components. The atomic spectrum is dominatedby free-bound emission in the Balmer and Paschen continua. The reflection component originatesfrom starlight that is scattered by dust. Dust particles are heated to a temperature of order 100 Kby radiation from the central stars, and cool by emission of infra-red radiation. The amount of dustcan vary strongly from nebula to nebula, and so does the strength of this reflection continuum.

The energy source of H ii regions is photoionisation of hydrogen, the most abundant elementin these regions, by the UV-radiation of the central stars. When a galactic cluster resides in theregion, the high-energy radiation of the few hottest members dominate the ionisation. Photonswith energies > 13.6 eV, the ionisation potential of hydrogen, are absorbed in this process. Theexcess energy of each absorbed photon is transferred to the liberated electron, that obtains a higherkinetic energy. Collisions between electrons, and between electrons and ions, distribute this energy.A Maxwellian velocity distribution with a temperature 5000 < T < 20000 K is maintained in typicalnebulae.

Collisions between free electrons and ions excite the lowest energy levels of the ions. Radiativede-excitation of these levels is forbidden by the electric dipole selection rule, therefore the transitionprobabilities are low. But the density in these regions is so low (ne ≤ 104 cm−3) that collisionalde-excitations are even less probable. As a result, almost every excitation leads to the emissionof a photon, producing an emission spectrum which is dominated by forbidden lines. Forbiddenlines provide useful diagnostics for measuring element abundances, densities and temperatures inthe nebula (see also section 2.2 and chapter 10).

Balmer and Paschen lines are emitted in the recombination of hydrogen. In this process, re-capture of an electron by an hydrogen ion results in an excited atom. In the cascade of radiativede-excitation, that may eventually lead to the ground level, line photons are emitted .

The equilibrium between recapture of free electrons by ions and photoionisation determines thedegree of ionisation at each point in the nebula. The higher the temperature of the central star,the higher the degree of ionisation. A nebula which shows forbidden transitions of highly ionisedelements (up to [Ne v] and [Fe vii]) is very likely to have a hot central star. This does not necessarilyindicate a high nebular temperature.

The nebular temperature is determined by the thermal equilibrium. While photoionisation,minus the losses from recombination, provides the heating, the most important source of radiativecooling is the collisionally excited line radiation (forbidden lines). The contributions of free-freeemission and collisionally excited hydrogen line radiation to the cooling are small (Osterbrock &Ferland, 2006).

2Host Galaxy NGC 55 and source

Our target is located in NGC 55, one of the brightest galaxies of the Sculptor Group. The SculptorGroup is very loosely concentrated and stretches from the outskirts of the Local Group out at about1.5 Mpc to 6 Mpc (Van de Steene et al., 2006). NGC 55 lies in the foreground of the Sculptor Group,forming a smaller loosely bound subgroup with the bright spiral galaxy NGC 300 (type Sd) andtwo spheroidal dwarf galaxies: ESO 410-05 and ESO 294-10 (Karachentsev et al., 2003).

Determinations of the distance to NGC 55 vary between 1.3 and 2.3 Mpc (Gieren et al., 2008),see table 2.1. The different methods and results of distance determinations will be discussed in

Figure 2.1: Image of NGC 55 obtained with the Wide Field Imager on the 2.2-metre MPG/ESO telescope at ESOLa Silla Observatory. The image is based on data obtained through B, V, and H-alpha filters. North is up, East tothe left. The field of view is 30 arc minutes wide. Picture and caption from ESO.

2.1. DISTANCE DETERMINATIONS OF NGC 55 17

section 2.1.Though difficult to estimate due to its high inclination of ∼ 80◦ (Hummel et al., 1986), the

morphological type of NGC 55 is likely to be SB(s)m: a barred spiral with an irregular appearance,very much like the Large Magellanic Cloud. However, unlike the LMC, NGC 55 is not a dwarfgalaxy: it has a diameter of ∼ 20 kpc, and a mass log (M/M�) = 10.5 (de Vaucouleurs & Freeman,1972).

Estimates of the metallicity of this galaxy range from 0.25 to 0.75 Z�, and even lower in extra-planar regions (Tullmann et al., 2003). More details on the metallicity determinations can be foundin section 2.2.

In section 2.3 we will introduce our selected massive star candidate.

2.1 Distance determinations of NGC 55

Table 2.1 shows a chronological list of distance determinations of NGC 55. We will briefly discussthe applied methods and comment on the obtained results.

� Pritchet et al. (1987) performed narrow-band photometry on individual C- and M-stars inNGC 55. From the mean magnitude of fourteen of the carbon stars they inferred a distancemodulus of (m−M)0 = 25.66±0.13 (d = 1.34±0.08 Mpc). The mean magnitude is comparedwith the mean luminosity of a sample of carbon stars in the Large Magellanic Cloud. Thelarge discrepancy with the other values listed in table 2.1 is probably due to an inadequatespatial resolution of 0.59 arc sec/pixel in their CCD images. The short distance they deriveis likely the result of the blending of some of their target stars with other nearby stars in thefield (Gieren et al., 2008). One would then measure a lower magnitude for these sources andconclude that they are closer.

� Karachentsev et al. (2003) used the empirical Tully-Fisher relation to determine the distanceto NGC 55. The intrinsic luminosity of a galaxy is proportional to the stellar mass. Theamplitude of the rotation curve of a spiral galaxy is related to the gravitational mass. Takinginto account a ∼90% contribution from dark matter to the gravitational mass, one can relatethe velocity dispersion to the intrinsic brightness of spiral galaxies. Since apparent brightness,as well as spectral line width, can be observed, this relation can be used to measure distances.Karachentsev et al. (2003) present a revised relation between absolute magnitude and H i line

Table 2.1: Distance determinations to NGC55. The first column gives the determined distance in Mpc with anerror estimate. Column 2 indicated the used method; for details we refer to the text. The last column provides thereference. Table from (Gieren et al., 2008).

Distance (Mpc) Method Reference

1.34± 0.08 Carbon stars Pritchet et al. (1987)1.8± 0.2 Tully-Fisher Karachentsev et al. (2003)2.12± 0.10 TRGB Tikhonov et al. (2005)2.3± 0.35 PNLF Van de Steene et al. (2006)1.91± 0.10 Cepheids, V I Pietrzynski et al. (2006)1.94± 0.08 Cepheids, V IJK Gieren et al. (2008)

18 CHAPTER 2. HOST GALAXY NGC 55 AND SOURCE

width based on eight other Sculptor galaxies with accurately known distances. They find adistance of 1.8± 0.2 Mpc for NGC 55.

� Tikhonov et al. (2005) analysed images from the HST/WFPC2/ACS archive to study thespatial distribution of AGB and RGB stars in three nearby galaxies: M 81, NGC 300 andNGC 55. They used the I-magnitude of the tip of the red giant branch in colour-magnitudediagrams to determine the distance and found 2.12 ± 0.10 Mpc for NGC 55. The I-bandmagnitude of the tip of the red giant branch is relatively insensitive to the age and chemicalcomposition of the population. This means it can be used as a distance indicator when thezero point is calibrated with an object of which the distance can be derived by other methods.Distances found with this method typically have an accuracy of ≤ 10%

� Van de Steene et al. (2006) analysed 21 planetary nebulae in NGC 55 to determine its distance.Planetary nebulae can be identified by blinking on- and off-band images in [O iii], Hα + [N ii]and the continuum. There are a few criteria to distinguish them from other emission-linesources such as H ii regions. Van de Steene et al. (2006) constructed the [O iii]λ5007 PlanetaryNebula Luminosity Function. Their most likely distance, corrected for the reference oxygenabundance of 12 + log(O/H) = 8.36, is 2.30± 0.35 Mpc.

� Pietrzynski et al. (2006) performed a study on Cepheids in NGC 55, which was extended byGieren et al. (2008). As a result of the precise relationship between pulsation period andintrinsic luminosity, Cepheids can be used as standard candles. Pietrzynski et al. (2006) andGieren et al. (2008) find 1.91± 0.10 and 1.94± 0.08 Mpc respectively.

There are indications that the distance obtained by Pritchet et al. (1987) is an underestimation.Excluding this value, the determined distances in table 2.1 agree with a value of 2.0 Mpc, whichshall be the distance we will use in further calculations.

2.2 Metallicity determinations of NGC 55

In this section we will summarise metallicity measurements for NGC 55 that can be found in liter-ature. All these studies derive oxygen abundances. After hydrogen and helium, oxygen is the mostabundant element in a typical H ii region (see also table B.1). Oxygen produces strong (forbidden)lines and the abundance is therefore relatively easy to measure even for a distant region. Becausethe individual nebular lines are optically thin, abundance determinations for gaseous nebula are lesscomplicated than for stellar atmospheres.

Many of the following oxygen abundance estimates are based on calibrations of the line fluxratio of

[O ii]λ3726, λ3729 + [O iii]λ4959, λ5007Hβ

.�� 2.1

This ratio is referred to in literature as R23 and is believed to be the most useful oxygen diagnosticwhen the temperature sensitive lines such as [O iii]λ4363 are unavailable (Pagel et al., 1979) becauseit correlates more tightly with oxygen abundance than other bright-line ratios do (Zaritsky et al.,1994). To derive abundances with this method, the electron temperature Te and electron densityne should also be measured. The [O iii] or [N ii] lines can be used for measuring Te and a pair of[S ii] lines is often used to derive ne. More details on this techniques are given in section 10.2.

� Webster & Smith (1983) performed spectrophotometrical observations of about 80 H ii regionsin different Sculptor Group galaxies; one of them being NGC 55. For this galaxy they obtained

2.2. METALLICITY DETERMINATIONS OF NGC55 19

spectra of 7 regions. Region 2 in their paper is very close to the location of our source. For thisregion they derive an oxygen abundance of 12 + log O/H = 8.23 based on the strong oxygenlines (R23). The other method is based on only the weak [O iii]λ4363 line; this resulted in avalue of 8.39. The temperature of this region is derived to be 10180 K. Furthermore, theyconclude that there is no radial gradient in the oxygen abundance across NGC 55. The nitrogenabundance appears to be generally lower than any other H ii region with a comparable oxygenabundance. Webster & Smith do not give a formal error estimate on their results, but statethat, for this region, the errors in the measured line strengths are typically 20− 40%.

� Stasinska et al. (1986) present spectrophotometric observations of 15 H ii regions in 10 irregulargalaxies. In NGC 55, they investigated three different regions for which they derived theoxygen abundance via the R23 ratio. They found 12 + log O/H = 8.15, 8.36 and 8.53, forthe three regions. The latter one region is located near the centre of the galaxy, thoughthe exact location is not indicated. For this centre region they derive ne < 100 cm−3 andTe = 9200 ± 800 K, while the other regions are hotter and denser. The uncertainties in theline intensities are estimated ro range from about 10% for the strong lines to 30% for theweakest lines. A large error results from the temperature determination, which is based onthe weak [O iii]λ4363 line intensity. Unfortunately, Stasinska et al. do not give a formal errorbar on their abundances.

� Zaritsky et al. (1994) investigated the relationships between the characteristic oxygen abun-dance, its spatial abundance gradient and macroscopic properties of galaxies by examiningindivial H ii regions in those galaxies. For NGC 55, one of the 39 galaxies in the sample, theyused the data from Webster & Smith (1983) (the average of 6 of their regions). With the R23

line ratio method they derive 12 + log O/H = 8.35 ± 0.07. Their value is different becausethey use a different calibration than do Webster & Smith. Their error estimate results fromthe differences between their three used calibrations. They did not propagate the error fromthe measured line strengths.

� Tullmann et al. (2003) investigated star formation in the gaseous halo of NGC 55. Thereforethey obtained deep spectra of two extraplanar regions (EHR1; EHR2), the diffuse ionisedgas (DIG) and an H ii region in the disk. For the calculation of temperatures, densities andionic abundances they use the nebular abundance tool (NAT) based on De Robertis et al.(1987). The R23 method is applied as well. They find oxygen abundances (with the R23 resultin parentheses) 12 + log O/H = 8.05 ± 0.10 (8.08) for the central H ii region, 7.77 (7.61) forEHR1, 7.81 (7.68) for EHR2 and 7.91 (7.88) for the DIG. We are mainly interested in the diskH ii region, but it is very interesting that the extra-planar regions have these low abundances.They claim that the flux errors are 14% at most, and for the H ii region less or equal to 9%.

In table 2.2 we summarise the values that have been derived for regions that are closest to the regionof our source. In this table we compare these with the solar oxygen abundance 12 + log O/H = 8.66from Asplund et al. (2004). Assumed that the other elements scale with oxygen as their relativeabundances in the sun, we can then give an estimate of Z.

The derived oxygen abundances vary substantially between the different determinations. Al-though some determinations are given without error estimates, the abundance derived with theR23 method depends exponentially on the temperature. This is often derived from [O iii]λ4363, aweak line introducing a large error. Tullmann et al. (2003) explain the variations in oxygen abun-dance by different slit positions. Their slit cuts through a region which apparently has the highest


Table 2.2: Oxygen abundance measurements of NGC 55. The first column gives the measured oxygen abundancein numbers of particles relative to hydrogen, with the error estimate if this was given in the paper. The secondcolumn gives a metallicity estimate based on this number, using 12 + log O/H = 8.66 (Asplund et al., 2004) for thesun (Z� = 0.0126). The third column gives a key word on the method; we refer to the details in the text. The lastcolumn provides the reference.

12 + log O/H Z/Z� Method Reference

8.39 0.54 [O iii]λ4363 line Webster & Smith (1983)8.23 0.37 R23 Webster & Smith (1983)8.53 0.74 R23 Stasinska et al. (1986)8.35± 0.07 0.49 R23 Zaritsky et al. (1994)8.05± 0.10 0.25 NAT Tullmann et al. (2003)8.08± 0.10 0.26 R23 Tullmann et al. (2003)

temperature (Te = 11500 ± 300 K) and therefore the lowest oxygen abundance. This spread is indisagreement with Webster & Smith (1983), who conclude that there is no radial oxygen abundancegradient in NGC 55.

For our standard grid of model atmospheres, that is to be used in section 9.2, we will assume ametallicity Z = 0.5 Z�, which is an average of the derived values in table 2.2. In chapter 11 we willconstruct a consistent picture of the source and its surroundings. An estimate of the metallicityof the region in which our star is located is obtained. In chapter 12 we will compare our derivedoxygen abundance with the values in table 2.2.

2.3 Source NGC 55 C1 31

In the context of the Araucaria project, Castro et al. (2008) performed VLT-spectroscopy of bluemassive stars in NGC 55. Low resolution (R = 780) optical (330 − 621 nm) spectra were obtainedwith VLT-FORS2 for approximately 200 blue intermediate to high mass stars, of which 164 couldbe classified.

NGC 55 C1 31 was selected as a candidate most massive star in this region because of its clas-sification as an early O supergiant. Objects of this spectral type are expected to be the initiallymost massive stars. The location of this star in NGC 55 is indicated in figure 2.2, which also givesan idea of the scale we are looking at. The spectrum resulting from the Araucaria project is shownin figure 2.3. Other properties of this source from Castro et al. (2008) are given in table 2.3.To investigate this source further, a more sensitive spectrograph with a higher spectral resolutionand a broader spectral range is required.

2.3. SOURCE NGC55 C1 31 21

Figure 2.2: Image of the host galaxy NGC55 with a zoomed in region (acquisition image in R-band) that indicatesthe location of our source NGC 55 C1 31. The green box is the projection of a 1′′ × 11′′ slit, according to one ofX-Shooter’s slit masks, in an arbitrary orientation. The physical region size assumes a distance of 2.0 Mpc (seesection 2.1). Colour image copyright Robert Gendler (2008).

Figure 2.3: A low (R = 780) resolution spectrum of NGC55C1 31 obtained with VLT-FORS2 in the context of theAraucaria project (Castro et al., 2008). It has been classified as early O supergiant, because of the strong wind-relatedemission at He iiλ4686.


Table 2.3: The properties of NGC 55 C1 31 in the catalogue of blue massive stars from Castro et al. (2008). Thecolumns list: (1) star identification; (2) right ascension (hh:mm:ss); (3) declination (dd:mm:ss); (4) V -magnitude; (5)V − I-colour; (6) spectral type; (7) projected galactocentric distance (”); (8) radial velocity (km s−1); (9) signal-tonoise ratio of the spectrum; (10) comments.

ID RA(J2000) Dec(J2000) V V − I ST Rg vr SNR Comments(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)

C1 31 00:15:00:01 -39:12:41.39 18.523 −0.716 Early OI 100.92 157 40 neb.

3This project

In this project we aim to reveal the nature of NGC 55 C1 31 by spectral analysis. This source wasselected as a candidate most massive star in galaxy NGC 55 because of the classification (early Osupergiant) from Castro et al. (2008). In order to obtain as much information as possible on thisdistant source, we have observed it with X-Shooter, the most sensitive spectrograph on earth. Thisnew instrument will be described in chapter 4. An overview of the observations as well as theobserving conditions is given in chapter 5. In chapter 6 we will describe the procedure of data re-duction. The resulting spectra of this source and its surrounding nebula are presented in chapters 7and 8. In chapter 9 we will derive constraints on the properties of the central star. This will bedone with help of its observed luminosity and with the modelling of photospheric line profiles. Inchapter 10 we will derive properties of the surrounding nebula based on the emission line spectrum.The conclusions on the central source and the surrounding nebula are united in a consistent picturein chapter 11. In chapter 12 we will discuss the results and compare them with former research onthis region and this source in particular. Chapter 13 contains the main conclusions. In appendix Aone will find the full grid of model parameters that is used in section 9.2. Appendix B contains atable of abundances used in chapter 11.

Observing NGC 55 C1 31 served as a pilot project for follow up X-Shooter spectroscopy of othermassive star candidates: five in NGC 55 and five in IC1613. What this project contributed to ourinsight in observing these kinds of objects and interpreting their spectra will also be discussed inchapter 12.

Part II

Observations and Data Reduction

4Instrument X-Shooter

X-Shooter is the new wide-band single target spectrograph for the Cassegrain focus of on UT2 ofESO’s Very Large Telescope in Chile. The instrument received first light in November 2008 andended its commissioning phase in October 2009. X-Shooter was built by a consortium involvinginstitutes from Denmark, Italy, The Netherlands, France and ESO (D’Odorico et al., 2006).

4.1 Specifications

In one single exposure X-Shooter covers the spectral range from the ultraviolet to the near infra-red(300 to 2500 nm). It operates at intermediate resolutions; the resolving power R = λ/∆λ variesbetween 4000 and 14000 depending on wavelength and slit width (see table 4.1 and 5.2). Thesensitivity of X-Shooter is very high: the limiting magnitudes per spectral bin (using 2 pixels bin-ning in the spectral direction) are U = 20, V = 21.5, I = 21 and K = 18.8 (1h exposure, SNR = 10).

Table 4.1: X-Shooter characteristics and observing capabilities (X-Shooter User Manual, 2010).

Wavelength Range 300− 2500 nm split in 3 arms

UV-Blue arm Range: 300− 550 nm in 12 ordersResolving power: 5100 (1′′ slit)Detector: 4k × 2k E2V CCD

Visual-red arm Range: 550− 1000 nm in 15 ordersResolving power: 8800 (0.9′′ slit)Detector: 4k × 2k MIT/LL CCD

Near-IR arm Range: 1000− 2500 nm in 16 ordersResolving power: 5600 (0.9′′ slit)Detector: 2k × 1k Hawaii 2RG

Slit length 11′′

Beam separation Two high efficiency dichroicsAtmospheric dispersion compensation In the UV-Blue and Visual-red armsIntegral field unit 1.8′′ × 4′′ reformatted into 0.6′′ × 12′′

4.1. SPECIFICATIONS 27

Figure 4.1: Schematic overview of X-Shooter (X-Shooter User Manual, 2010).

28 CHAPTER 4. INSTRUMENT X-SHOOTER

4.2 Technical description

Figure 4.1 shows a schematic view of the layout of X-Shooter. The instrument consists of four maincomponents: the backbone and three instrument arms: UVB, VIS and NIR. Below follows a listof the pre-slit optics and units of the backbone. More details can be found in the X-Shooter UserManual.

� The instrument shutter is the first element in the beam coming from the telescope. Itallows safe daytime use of the instrument for tests and calibrations.

� The calibration unit allows to select a choice of flat-fielding and wavelength calibrationlamps.

� The acquisition and guiding slide contains 3 different mirrors to put into the beam fromthe telescope or from the calibration unit and the integral field unit (IFU). The IFU set-upcan be used to obtain simultaneous spectra in a two-dimensional field.

� The acquisition and guiding camera allows to detect an object from the U to the z-band.It consist a CCD providing a field of view of 1.47′ × 1.47′ and a filter wheel.

� The dichroic box consists of two highly efficient dichroic beams splitters that split the lightand distribute it over the three arms.

� One piezo tip-tilt mirror for each arm is used for active backbone flexure compensation.

� An atmospheric dispersion corrector (ADC) and a focal reducer are contained in theUVB and VIS pre-slit arms. The ADC compensates for atmospheric dispersion in order tominimise slit losses and to allow any positioning angle (within a zenith distance of 60◦).

Each instrument arm is a fixed format cross-dispersed echelle spectrograph, with an individual slitcarriage and shutter. The NIR spectrograph is fully cryogenic, whereas in the VIS and UVB armsonly the detector is cooled. The NIR spectrograph is cooled with a liquid nitrogen bath cryostatand operates at 105 K. The specifications of the instrument arms are given in table 4.1.

In an echelle spectrograph, the light is diffracted in two directions perpendicular to each other.See figure 4.2 for a schematic picture of the echelle principle. Firstly, the light is dispersed in thedirection perpendicular to the slit, to transform the wavelength scale to a spatial axis. Secondly, thelight is also dispersed in the direction parallel to the slit. The light is chopped in so called orders,which cover successive parts of the spectrum. Their wavelength ranges partly overlap at the edges.These orders, visible as luminous arcs on the X-Shooter raw images (see figure 6.1), are projectedabove each other on the detectors. The number of orders varies per instrument arm (see table 4.1).Figure 4.3 shows the position of the slit and the orientation of the orders on the detector for eachinstrument arm.

4.2. TECHNICAL DESCRIPTION 29

Figure 4.2: The principle of an echelle spectrograph. The two gratings need to be positioned orthogonally in orderto disperse the light in two directions. Figure by Boris Povazay (Cardiff University).

Figure 4.3: The slit coordinate system and the correspondence between object position in the slit and spectrumlocation on the raw frames for each instrument arm. A source at the position indicated with a black star in thetop panel (positive x) produces spectra placed as illustrated in the bottom panels. The boxes indicate the full slitprojection. The curves on the lower panels indicate the orientation of the orders on the raw frames and the coloursindicate the wavelength range: the shortest wavelengths are displayed in blue, the longest in red. PA is the positioningangle of the slit with respect to the North. Figure from X-Shooter Pipeline User Manual.

5Observations

The observations of NGC 55 C1 31 obtained for this project were part of the X-Shooter Science Veri-fication (SV) Run 1 and 2. SV runs serve as a final test of a new instrument and aim to demonstrateits scientific potential. A series of Science Verification observations is an integral part of the com-missioning of all new VLT instruments. These observations are performed in service mode and thecollected data are made publicly available. These early sets of data can be difficult to reduce andanalyse. Firstly because SV observations push the telescope and instrument close to their limits.Secondly because the set-up of instrument might not yet be optimal. Thirdly, since SV runs arealso aimed for testing the pipeline and reduction tools, these are also still in a test phase and mightnot yet be fully applicable.

The X-Shooter spectra of NGC 55 C1 31 have been obtained in the nights of August 13 (SV1),September 27 and 30 (SV2) in 2009. Table 5.1 gives an overview of the observations. The range inwhich the observational seeing varied during the nights is indicated. These values are derived fromthe full width half maxima of some point sources on the acquisition images that are taken in theR band. The lunar conditions are also indicated: moon distance from the source, illumination andlunar phase. The total observing time is 2.5 hours spread over 10 frames per arm. In the last columnwe indicate the positioning angle of the slit with respect to the North (see figure 4.3). The choicefor a particular positioning angle depends on the position of the object on sky. The observations of27 and 30 September have been taken around the same time in the night. With only a few days inbetween, the positioning angles do not differ that much.

All observations have been carried out in nodding mode. This means that the slit is positionedin a way such that the source is located at position A for half of the observations, and at positionB for the others by nodding the telescope; see figure 5.1. These positions lie respectively 2.5′′

above and below the centre of the slit, which has a projected length of 11′′. This is the standardobserving mode for infra-red spectroscopy. Since the different wavelength arms are bound to thesame instrument set-up, all images are in nodding mode. In this mode one can easily ‘remove thesky’ by just subtracting two images with the source in the two different positions. With sky, wemean the homogeneous light that enters the slit, which is mostly thermal background in the nearinfra-red and mostly moonlight in the ultra-violet and visual band, when observing at bright time.In the visual and near infra-red, the sky spectrum also shows emission lines from earth’s atmosphere.These are called telluric lines.

31

Table 5.1: Overview of the observations. Column 1: date and time of the observations and shift due to barycentriccorrection. Column 2: observational seeing from R band acquisition images. Column 3: lunar conditions includingangular distance from the source to the moon, illumination of the moon and lunar phase. Column 4: exposure timeper arm on the source. Column 5: positioning angle of the slit on the sky (see figure 4.3).

Date Seeing Moon Distance Exposure Time PositioningStarting Time (UT) (R-band) Illumination (Per Arm) Angle

∆vbc Quarter

13/08/09 69◦ 2× 900 s 53.410◦

08:51:45 0.69”-0.88” 56%9.394 km s−1 3

27/09/09 66◦ 4× 900 s −52.309◦

03:43:18 0.87”-0.93” 60%−8.386 km s−1 2

30/09/09 45◦ 4× 900 s −56.446◦

03:25:41 0.96”-1.50” 85%−9.529 km s−1 2

Table 5.2: The wavelength range, projected slit size, predicted resolving power (for the used slit width) and numberof orders of each instrument arm.

Arm Range (nm) Slit Size Resolving Power Number Of Orders

UVB 300− 560 0.8′′ × 11′′ 6200 12VIS 550− 1020 0.9′′ × 11′′ 8800 15NIR 1020− 2480 0.6′′ × 11′′ 8100 16

32 CHAPTER 5. OBSERVATIONS

A BB B A

Figure 5.1: The nodding positions A and B on the slit in indicated with orange circles. Cycles of observations arerepeated in A-B-B-A sequences. The black dot indicates the acquisition position of the slit. The nodding positionshave an offset of 2.5′′ from this centre point. The distance between the two nodding positions is called the nod throw,which is 5′′ in our set-up.

6Data reduction

The data are reduced with X-Shooter pipeline version 0.9.4. In section 6.1 the general steps infor X-shooter data reduction are described. We will explain why we did or did not perform anyconsecutive step, and how this, for example, can be done with the pipeline. We will also give anindication of quality of the intermediate results.

Concerning the final results, we distinguish two different kinds: one and two dimensional spec-tra (1D and 2D). Though these concepts will be treated in more detail in chapter 7, it is useful toexplain them beforehand. A 2D spectrum is an image with vertically a projection of the slit, andhorizontally the dispersion (wavelength) direction. The intensity of the image is the flux. In fact itis a long sequence of slit-shaped images in increasing wavelength. A continuum source appears insuch a frame as a blurred horizontal line. Spectra of objects in the neighbourhood of the intendedsource can sometimes also be identified. A 1D spectrum is what is more commonly referred to as aspectrum: flux on the y axis against wavelength on the x axis. The procedure of extracting the 1Dspectra from the 2D spectra, which done outside the pipeline, is described in section 6.2.

6.1 Data reduction cascade

In order to transform the raw output images from the instrument (see figure 6.1) to spectra thatare suitable for scientific analysis, we will need to perform the following steps:

� Determine the spectral format (6.1.1)

� Measure detector bias and dark current levels, and create master calibration files (6.1.2)

� Map bad pixels on the detectors and eventually check detector linearity (6.1.3)

� Guess solutions for the wavelength calibration and the location of the centre of the orders(6.1.4)

� Trace the order centres (6.1.5) and edges (6.1.6) to refine the guess solutions

� Measure variations in response of the detector (6.1.6)

� Determine the 2D transformation from pixel (X,Y on raw science frame) to wavelength (6.1.7)

� Compute efficiency and response of telescope, intrument and detector (6.1.8)

� Calibrate the flux (6.1.8)

� Correct for telluric absorption (6.1.9)

34 CHAPTER 6. DATA REDUCTION

� Apply all master calibration files and solutions to the science frames (6.1.10)

Each of these steps can be carried out by a routine in the pipeline, which is called a recipe. Recipesare fed with raw calibration images, reference files such as line lists and intermediate pipeline prod-ucts from previous recipes. The first sequence of recipes creates the necessary master calibrationfiles, which come together in the last step, where they are combined with the raw science framesto produce a calibrated science spectrum. In the following subsections the consecutive steps in thereduction cascade are described: what calibration image is needed, for which physical effect do wecorrect and with what recipe. The X-Shooter Pipeline User Manual provides more details on thealgorithms used at each step.

6.1.1 Determining the spectral format

The spectral format is, roughly speaking, a description of the location of the echelle orders andthe wavelength calibration. With highly curved orders, variable line tilt, dispersion and spatialscale along each order, it is relatively complex. We created a spectral format table by executingthe pipeline recipe xsh util physmod on the model configuration file (from commissioning run 4)and a reference arc line list, which are distributed along with the Pipeline. This spectral format,obtained by applying the predictions of the physical model, is an estimate that will be fine-tunedin the successive data reduction steps.

6.1.2 Measurement and subtraction of detector bias and dark current levels

A bias frame is a zero second exposure taken with the shutter closed. These frames are used tomeasure the bias level and the structured noise and to detect bad columns of the CCD detector.The bias level results from the electric field in the CCD. With the recipe xsh mbias, we stacked

Figure 6.1: Raw science frames from 13 August 2009, orientation as in figure 4.3. Left: UVB arm, middle: VIS,right: NIR. The continuum of the star is visible as a bright line along the orders. The bright spectral lines in the UVBand the short wavelength side of the VIS (arcs at the left) are emission lines from the surrounding region, recognisableby the inhomogeneous illumination of the slit. Most emission lines in the NIR and VIS raw frames are sky emissionlines. Those illuminate the slit homogeneously, resulting in smooth lines. The source’s continuum becomes less clearlyvisible at longer wavelengths. The upper arcs in the NIR raw frame (longest wavelengths) are strongly contaminatedby thermal emission from earth’s atmosphere.

6.1. DATA REDUCTION CASCADE 35

five bias frames to create a master bias calibration file for each night for the VIS and UVB arms.

Dark current is the small electric current that runs through for example a CCD, even whenit is not illuminated by photons. This current results from randomly generated electrons in thesemi-conductor material that are swept by the electric field. One can correct for this using darkframes, that are occasionally taken for the UVB and VIS arm. A dark frame is a long exposure(typically one hour or more) with the shutter closed. The recipe xsh mdark creates a master darkframe from a set of dark frames.

Master bias and master dark frames are used to monitor the detector performance via qualitycheck parameters such as bias and dark current levels, read-out noise, X and Y structure of thenoise. The other calibration and science frames can be corrected for bias and dark current by sub-traction of these master frames.

According to the pipeline and instrument manuals, dark current correction applies mainly for thereduction of the NIR data; for this the bias and dark current correction is combined. The correctionis obtained by subtracting a frame with calibration lamps switched off from ones with calibrationlamps on. In the UVB and VIS data, the dark current contribution is considered negligible: wechecked this by subtracting the master bias frame from the master dark frame for one of the nights.We did therefore not perform dark current correction for the VIS and UVB data.

6.1.3 Bad pixels and detector linearity

A defective pixel is unable to sense the light level correctly. If a reference bad pixel map is providedto xsh mbias, this recipe also produces a hot (bright), cold (dark) and an updated bad pixel map.We combined them to a master bad pixel map with xsh util bpmap coadd. Table 6.1 gives thenumbers of hot and cold pixels that were detected.

Pixels may respond in a non-linear way. Locations of non-linear pixels can be found withlinearity flat frames. Those calibration files can be processed with the recipe xsh lingain to createa bad pixel linearity map. However, during the first period of observation with X-Shooter (up untilJanuary 1st, 2010), in which our observations are done, very few linearity frames are acquired, andthey were not publicly available. Therefore we did not perform this step.

Table 6.1: The numbers of hot and cold pixels that are detected in the bias frames per night and per arm. Belowthe label of the wavelength arm we indicated size of the window that is used on the CCD detectors. The UVB CCDis 2048× 4102 pixels, but a window of only 1800× 3000 pixels is used (X-Shooter User Manual, 2010). The numbersof hot and cold pixels are negligible compared to the millions of pixels on the raw frames.

NightUVB VIS

1800×3000 2048×4096

cold hot cold hot

13/08/09 0 4 2 5727/09/09 3 6 1 4130/09/09 7 1 1 31


Figure 6.2: A format-check frame is VIS (ThAr lamp through single pinhole) over-plotted with the predicted linelocations that are identified on the calibration frame.

6.1.4 Guess solutions for order centre location and wavelength calibration

With the recipe xsh predict (poly mode), we created a guess solution of the location of the centreof the echelle orders and the wavelength calibration. It uses a format-check frame, a reference listof arc lines and the spectral format table from the first step. To create a format check frame, theCCD is illuminated with a thorium argon (ThAr) calibration lamp through a single pinhole at thecentre of the slit.

The arc lines are localised by fitting two-dimensional Gaussian functions to the bias subtractedformat-check frame, following the reference arc line list. Figure 6.2 shows a part of a VIS format-check frame over-plotted with the arc line locations that have been identified on the calibrationframe. The region file is constructed by executing quality check script test xsh resid tab on theformat-check residual line table, a product of xsh predict. The new (X,Y) positions of the linesare used to fit a guess global polynomial solution. This solution gives the pixel position (X,Y) fora given order and wavelength at the center of the slit.

The differences between the X and Y positions of the lines on the theoretical map, and the fittedX and Y positions from the clean arc line list have average and median values of zero in UVB andVIS frames for all nights. The differences between the fitted positions and the polynomial haveminimal and maximal values of respectively ∆X = −0.22 and 1.01 pixels and ∆Y = −0.05 and 0.05pixels for UVB and ∆X = −0.41 and 0.27 pixels and ∆Y = −0.05 and 0.05 pixels for VIS. Thismeans that the format-check frame can well be matched with the line list, and the guess solutionsfor order centre location and wavelength calibration are already quite accurate.


Figure 6.3: Part of a prepared VIS order definition frame over-plotted with a region file produced from the ordercentre solution of xsh orderpos. The differences are so small that they are not visible on this scale.

6.1.5 Tracing of echelle orders

The guess order centre solution is refined with xsh orderpos. This recipe makes use of an orderdefinition frame. An order definition frame is an exposure through a pinhole, but now illuminatedby a continuum flatfield lamp (two lamps for the UVB calibrations). It is a very high signal-to-noiseframe that describes the location of the centre of the orders precisely.

Figure 6.3 shows a part of a prepared VIS order definition frame over-plotted with the ordersolution region file. We constructed region plots with quality check script test xsh data orderwhich is provided with the pipeline kit.

The residuals in order location vary between a minimum of −0.05 and a maximum of 0.05 pixelsfor the VIS frames and between a minimum of −0.10 and a maximum of 0.10 pixels for the UVB.

6.1.6 Flat field correction

Now the location of the centre of the orders is precisely known. To determine the edges of the orders,a full slit exposure is needed. Flat field frames are long slit exposures taken with a continuum lamp(the same lamps as used for order definition frames). They need to be taken in the same set-up asthe observations; for example slit width should be equal to the one used for the science observations.For the UVB arm, two lamps are used to guarantee enough illumination in the whole wavelengthrange.

Flat field frames give information on the response of the detector and of the extent of the ordersin the spatial direction. There are variations on three different size scales. On the large scale, itallows to correct for the blaze function. This effect, which is intrinsic to the dispersion in the echelle


Figure 6.4: A zoomed in part of a prepared VIS flat field image over-plotted with a region file created from the orderedge solution with the quality check script test xsh data order. In the overview image in the right upper corner onecan see the effect of the blaze function over the flat field image as a whole.

grating, makes the middle of the orders brighter than the edges. On the intermediate scale wecorrect for the fringing, an effect that occurs in the red and infra-red which is due to interference inthe detector’s coating layer. On the smallest scale, flat fielding corrects for pixel-to-pixel variations.

We stacked five flat field frames with recipe xsh mflat to increase statistics and to reject outlierssuch as cosmic rays. We created a master flat which is also background subtracted to eliminatediffused light from the orders in the inter-order regions.

Figure 6.4 shows zoomed in part of a prepared VIS flat field frame along with a region file ofthe solution for the order edges. On the small overview image in the right upper corner one can seethe effect of the blaze function.

6.1.7 Two dimensional mapping and wavelength calibration

The recipe xsh 2dmap uses multi-pinhole arc lamp frames to determine the 2D transformation frompixel (X,Y) to wavelength and to rectify the spectral format. Multi-pinhole arc lamp frames areexposures of the ThAr lamp through a row of nine pinholes. With this recipe, we calibrated thewavelength and spatial scale simultaneously. Where the xsh predict recipe only calibrated thewavelength for the order centres we are now able to calibrate the wavelength along the full slit.This allows us to correct for the tilt, which varies along the orders.

For both UVB and VIS raw frames, the Y direction is the dispersion direction (see the rawframes in figure 6.1). The standard deviation of the residuals in the Y direction for the UVB framesare around 0.18 pixels, which corresponds to 0.0018 nm. For the VIS frames, the standard devi-


Table 6.2: The resolving power that is measured from slit arc lamp calibration frames with matching slit width (0.8′′

for UVB and 0.9′′ for VIS). The median (column 3) and the root mean square (column 4) of the measured resolvingpower are quality check parameters given by the xsh wavecal recipe. The last column states the theoretical resolvingpower for the used slit width according to the X-Shooter User Manual.

Arm NightResolving PowerR rmsR Rth

13/08/09 6299 343UVB 27/09/09 6283 304 6200

30/09/09 6233 292

13/08/09 7744 317VIS 27/09/09 7878 599 8800

30/09/09 7812 799

ation of the Y residuals are 0.27 pixels on average, corresponding to 0.0027 nm. The accuracy ofthe wavelength calibration is 1.2 km s−1 on average for VIS and UVB. According to the X-ShooterUser Manual, the accuracy of the wavelength calibration typically achieved with this data reductionsoftware is 2 km s−1, so this is an excellent result. The accuracy of the wavelength calibration ismuch higher than the spectral resolution of the instrument, so this small error will not be significantin the final results.

From slit arc lamp frames, which are long slit exposures taken with the ThAr lamp, recipexsh wavecal can refine the wavelength calibration obtained from the multi-pinhole frames. Wedid not perform this refinement because the wavelength calibration obtained with xsh 2dmap wasalready good enough. The slit arc lamp frames can also be used to measure the instrument resolvingpower. In table 6.2, the measured median (R) and root mean square (rmsR) of the resolving poweris listed per night for the UVB an VIS instrument arms. The spectral resolution of the UVBarm matches the theoretical resolving power Rth very well. The results of the VIS arm reach asignificantly lower resolving power than the prediction for this slit width from the X-Shooter UserManual.

6.1.8 Computing efficiency and response, flux calibration

The translation between ADU and entrance flux varies with wavelength in a way that depends onthe telescope, instrument and detectors. One may use a flux standard star observation and a tablecontaining its flux distribution in order to flux calibrate the science frames. At the same time onecorrects for the instrument response and the telescope + instrument + detector efficiency. Withthe pipeline for example, this can be done with the xsh respon stare recipe. This recipe was notyet available in pipeline version 0.9.4. However, for the comparison of emission line strengths insection 10.2.2 we have applied a relative flux calibration.

6.1.9 Telluric correction

The disadvantage of ground-based observations is that the light has to pass through earth’s atmo-sphere. This introduces absorption, continuum (thermal) emission, and line emission. The thermalemission is strongest in the longest wavelength part of the near infra-red spectrum. Telluric features


(emission and absorption), mostly present in the VIS and NIR, can be localised in the spectrum ofa telluric standard star. Such a star has a minimum of stellar absorption features in its spectrumwhere the emission and absorption lines from earth’s atmosphere are located. With the pipeline,the telluric correction can be carried out with the recipe xsh telluric stare, but this was notyet available in Pipeline version 0.9.4. There were observations of some of these standard starsavailable, so we could have done it by hand.

Luckily the telluric features are far less abundant in the spectral range of VIS and UVB armsthan in the NIR. Since we did not use the NIR frames for scientific analysis because of their lowerquality, we decided to not apply a telluric correction to any of the frames. By performing a skysubtraction on the obtained spectra (see section 6.2), we correct for atmospheric emission, but notfor atmospheric absorption.

6.1.10 Processing science frames

The last and most important step performed within the Pipeline is the reduction of the raw scienceframes with help of all calibration files that have been created in the previous steps. For every modein which science frames can be recorded there is a separate recipe. Despite the fact that our datawas taken in nodding mode, it turned out that the xsh science slit nod, especially suited for thismode, was not satisfying our needs.

While executing the recipe xsh science slit nod, two images with different nodding positions(A and B, see section 5.1) of the source in the slit, are subtracted from each other. This is theconventional way to do (near) infra-red spectroscopy. In one image the source is at location A; inthe other there is a patch of sky at the same position, while the source is at B. The result of thesubtraction of A from B is an 2D image with a sky-corrected stellar spectrum at the centre and twonegative sky-corrected spectra above and below it (see figure 6.5). But this will only give a properresult with a relatively homogeneous background; something we did not have. There are otherobjects visible in the neighbourhood of the source: other stars and objects that are even brighterthan our source at specific wavelengths. This will be more thoroughly discussed in chapter 7.

The final step was instead performed with the xsh science slit stare recipe. This recipe isdesigned for observations in stare mode, where the source is located at the centre of the slit. Sincewe only use the 2D result of this recipe, this is not an issue. We performed no automatic skysubtraction in this step. Instead of using the automatically generated 1D result, we extracted the1D spectrum manually from the 2D spectrum. Likewise, we also extracted ‘sky’ spectra, in order todo a proper sky subtraction. The reason for this manual procedure is that it allowed us to specifythe location and size of the extraction windows in a distinct way. The procedure will be describedin the following section.

6.2 1D extraction

A MIDAS script is written and used to extract one dimensional spectra from the generated twodimensional spectrum. Per night, there are 2 or 4 sequential exposures per arm (see table 5.1).The exact location of the source on the two dimensional spectrum, and how far it stretches out dueto seeing is different for each night. The borders in vertical pixel space are therefore determinedmanually per arm and per night. Furthermore, two dark pieces of are also selected manually perarm and per night. Table 6.3 lists the centre locations of the object and two sky patches per nightand per wavelength arm. The two dimensional spectrum is 110 pixels high (1 px = 0.1′′). The size

6.2. 1D EXTRACTION 41

Figure 6.5: The upper image shows a part of the output 2D VIS spectrum of xsh science slit nod between 626and 690 nm (27/09/09). The lower image is the 2D result from xsh science slit stare, without sky subtraction.The bright blurred horizontal line in the middle of the upper image is the sky-corrected stellar continuum; the twodark lines above and below are a result of subtracting the continuum from the sky. As a result of being created fromfour images in two different nodding positions, this 2D frame has more pixels in the spatial direction than the 2Dspectra from the xsh science slit stare recipe. The strong feature at 656.28 nm is Hα. At this line’s wavelength,the sky is not al all homogeneous. Therefore the sky subtraction accomplished with the nodding reduction mode ismeaningless at this wavelength, which can be seen from the distribution of light at the Hα line. This is true for mostBalmer lines and some Helium lines. To have more control over the sky subtraction, we have used the staring reductionmode without sky subtracting, and selected and subtracted the sky by hand. The most right narrow emission lines,visible in the lower panel, are telluric lines. These are homogeneously distributed in the spatial direction and aretherefore neatly subtracted in the nodding reduction mode. They have disappeared in the nodding reduction moderesult (upper panel).

of the windows of the sky patches is 10 pixels. The size of the object window is 10 pixels for theAugust night, and 14 for the September nights. This is necessary for still extracting as much lightas possible while the seeing is higher (see table 5.1).

The script will produce three different spectra for each arm and each night: a raw one, whichwill be the extracted spectrum at the location of the star with no sky correction, a sky spectrum,which is average of the extracted spectra at the two selected sky locations and the obj one, whichis simply the difference between the raw and the sky.

What is meant by extraction is a vertical collapse of the selected part (summing). Optimally,fitting a Gaussian profile in the slit length direction would be the correct way to do this, but in ourcase this was not possible due to the inhomogeneous background. Even very close to the source thereis a difference in background level above and below the source in the slit, especially at wavelengthsof strong hydrogen lines. This suggests that the region is full of emission line objects or nebulae,and we should be very careful with our selection of the sky. We will discuss this more thoroughlyin chapter 7.

One has to keep in mind that since the slit orientation (positioning angle) differs between thenights, the patches of sky we selected are not even exactly the same even if the same borders areused. The final sky-spectrum we produce is in fact an average of at least four different regions: twofor the differences in positioning angle between August and September, times two because we usethe average of two sky patches. The obj spectra per night are produced by only subtracting thesky average of the same night.


Table 6.3: Centre locations (pixels) for object and two patches of sky in two nodding modes A and B. The heightof the 2D spectra is 110 pixels (1 px = 0.1′′). The window size for the sky patches is 10 pixels. The window size forthe object is 10 pixels for August and 14 pixels for September.

Arm Nightobj sky1 sky2

(A,B) (A,B) (A,B)

13/08/09 29,80 12,63 66,56UVB 27/09/09 27,79 10,62 44,96

30/09/09 29,80 12,63 46,97

13/08/09 29,80 12,63 66,56VIS 27/09/09 30,81 13,64 47,98

30/09/09 33,84 16,67 50,101

13/08/09 29,80 12,63 66,56NIR 27/09/09 28,78 11,61 45,95

30/09/09 33,84 16,64 50,101

Figure 6.6: This is a small part (between 429 and 442 nm) of the first 2D UVB spectrum of 13/08/09, noddingposition A. The red boxed region is where the raw spectrum is extracted; the green lines mark the borders of the twosky regions, see table 6.3. From this image, one can already predict that this will work in the continuum, but mightnot be suitable at strong emission lines. The left strong emission line at 434.0 nm is Hγ; the right one is Ca ii at 436.3nm.

6.2. 1D EXTRACTION 43

Sky subtraction is very important to remove the lunar contamination. As can be seen in table 5.1,this contamination is severe especially in the September nights, when most of the observations aredone. The diffuse moonlight will also be seen by the telescope, and since the intended source isrelatively weak, we see a very strong and recognisable spectrum of the sun appearing in in the rawand sky spectra. However, since the moonlight intensity is constant along the slit and there is noradial velocity profile, it is neatly subtracted: in the obj frame there are only very small residuals,see for example figure 7.1 in chapter 7.

In the September nights, where there are 4 exposures, we are able to compute the median ofthe four values for a certain pixel (the instrument or detector and thus the spectral format is notexpected to have changed between the sequential exposures). With this method, cosmic rays areremoved, since the over-illuminated pixels caused by this phenomena will only end up in one of thefour exposures at the same position. When there are only two exposures, the average of the pixelvalue is taken.

After applying barycentric correction to the raw, sky and obj spectra of the different nights,they have been added weighted by their signal to noise. The spectra have been re-binned to stepsize of 0.06 nm, matching the resolution of the instrument.

We will see in the results in figure 7.1 and in the spectra shown in chapter 8, that residualsare left in the obj spectra at the nebular emission line locations. The sky spectra do not onlycontain clean sky light (moonlight; thermal and line emission form earth’s atmosphere), but alsolocal emission lines; something intrinsic to the region where our target is located. Subtracting themcauses the residuals at the emission lines. An option to avoid residuals is to remove the nebularemission lines from the sky spectra, before subtracting it from the raw. This can be done by fittingGaussian functions to them. In that case, the obj spectrum would still show the nebular emissionat that location, but now sky corrected and without the residuals. For our purposes, this procedurewill not be necessary.

The signal-to-noise ratios of the different spectra and the final combined results are given intable 6.4. In chapter 8 we show the final obj spectra (figures 8.1 to 8.3), but we will first have acloser look at the 2D spectra in chapter 7.


Table 6.4: Signal-to-noise ratios (SNR) of different parts of the UVB and VIS spectra per night, and of the finalcombined spectra. The resulting total spectra are the result of adding the different nights weighted by their averageSNR. They have also been re-binned to a step-size of 0.06 nm, matching the instrument resolution. As is shownin table 5.1, the total exposure time per night of the August observations was only half of that of the Septemberobservations. Still, the SNR is comparable. The observational conditions were clearly better in the August night.

UVBSNR [range (nm)]

Spectrum Night [424:428] [460:465] [505:510]

raw 13/08/08 11.96 22.92 20.5127/09/09 13.89 27.45 23.3730/09/09 9.37 24.27 20.72

total raw re-binned 13.65 37.73 32.46

sky 13/08/09 7.07 14.33 11.2827/09/09 8.78 17.33 14.8930/09/09 7.31 18.45 16.35

total sky re-binned 9.22 24.86 22.34

obj 13/08/09 4.78 9.14 8.7527/09/09 4.22 8.97 7.6830/09/09 2.50 5.56 5.22

total obj re-binned 13.89 24.18 22.34

VISSNR [range (nm)]

Spectrum Night [604:609] [675:680] [811:816] [975:978]

raw 13/08/09 16.19 20.81 21.45 7.6927/09/09 23.21 29.01 31.35 9.9430/09/09 28.55 34.71 34.46 11.63

total raw re-binned 59.66 80.57 64.55 19.31

sky 13/08/09 10.90 11.80 13.92 3.6827/09/09 14.10 16.10 14.76 6.6030/09/09 20.35 22.05 22.30 7.77

total sky re-binned 38.60 49.27 35.85 14.30

obj 13/08/09 5.35 7.75 7.92 2.0227/09/09 5.41 7.60 9.36 1.9530/09/09 3.80 6.30 7.22 1.50

total obj re-binned 13.52 22.02 21.97 5.67

Part III

Observational Results

72D spectra

7.1 2D-1D correspondence

A 2D spectrum is a sequence of projections of the slit in increasing wavelength. The result is animage with on the vertical axis the spatial direction along the slit and on the horizontal axis wave-length. The source appears as a blurred horizontal line. How a 2D spectrum corresponds to themore common 1D spectrum is shown in figure 7.1, which shows a small part of the UVB spectrumbetween 381 and 407 nm. The green line is the resulting spectrum extracted at the source location:what we earlier called a raw spectrum. The red line is the average of spectra of the sky, extractedfrom the object free regions in the slit, see figure 6.6. In this spectrum, we can recognise the solarspectrum which is coming from the moon. In UVB, the moonlight is the main contribution to thesky. For example, the strong Fraunhofer absorption lines Ca ii K and H can be distinguished. Theblack spectrum is the resultant obj spectrum. The spectral features from the moonlight have beenneatly removed.

7.2 Other objects

2D spectra give us information on the surrounding region. Apart from the aimed source, the lightof other objects will be seen by the instrument and will produce a spectrum. For example, thebrightest horizontal line in the 2D spectrum in figure 7.1 is the continuum of our star, but above itthere is another weaker continuum visible.

One has to keep in mind that to examine the off-source objects, one must look at one night’sobservations at the time. This is because the positioning angle is different for the different nights(table 5.1). The slit has a different angle with respect to the stars and other objects will come in toview. This is illustrated in figure 7.2. The slit projection (in this case the VIS 0.9′′ slit) for threedifferent nights is shown on an Hα image of the region of our source. The cyan rectangles show theslit projections in A and B nodding mode for 13 August 2009, the green ones for 27 September andthe white ones for 30 September. It is easy to see that the 2D spectra of August and Septembercan not be compared, except for the intended source of course.

Sometimes, however, a secondary source will be less evident when it shows no continuum emis-sion. This is what occurs in the 2D spectra of 27 and 30 September. Figure 7.3 shows a small partbetween 404 and 416 nm of the first two 2D spectra of 27 September. One is in nodding position Awhile the other is in B. The primary source’s continuum appears brightly in the lower position inthe upper panel, and in the higher position in the one below. At the lower edge of the lower image,we see another continuum source. Between these two continuum emitters, there is another object

7.2. OTHER OBJECTS 47

Figure 7.1: The correlation between a 2D and the 1D spectra. The upper panel shows the final combined andre-binned spectra: green is raw, red is sky and black is obj. The range is 381 to 407 nm. The 2D spectrum in thelower panel is the first exposure of August. The stellar continuum is visible as a blurred horizontal line; the offsetfrom the middle of the slit is due to the nodding mode (this is position A). A few lines are labelled. The Ca ii K andH lines are very strong Fraunhofer absorption lines from the solar spectrum (through the moonlight). One can seethat these are neatly subtracted with only a small residual in the obj spectrum (black). The smaller features we seein the green raw and the red sky spectrum are clearly not noise, but lines in the solar spectrum, since they are similarat different locations in the slit. The emission lines are not homogeneous in luminosity throughout the slit. This isan indication that is a local effect: the emission from the surrounding H ii region(s). As can be seen by the residualsin the obj spectrum (black), it is more difficult to subtract the nebular emission in a proper way. The line from thesurrounding region apparently has a small wavelength shift with respect to the line from region we use to subtractthe sky.

48 CHAPTER 7. 2D SPECTRA

Figure 7.2: This is a zoomed in part of a FORS Hα image of NGC 55 from the ESO archive. The rectangles showthe projections of the VIS slit (0.9′′ × 11′′) in nodding positions A and B. Cyan is 13/08/09, green is 27/09/09, whiteis 30/09/09. 1′′ corresponds to a physical size of approximately 9.7 pc.

emitting only at certain wavelengths: an emission line object. The source emerges when we draw afew contour plots. The figure shows that at the Hδ 410.2 nm line, the emission at the location ofthis object is higher than at our source. To quantify this we computed a spatial profile at this lineand compared it to a spatial profile in a continuum region. The two graphs at the left in figure 7.4show the result of this procedure at Hδ: the black line is the spatial profile at the line, the red oneis an average of the continuum between 410.5 and 433.3 nm, where there are no prominent emissionlines. The upper figure is with the primary source in A position (upper panel in figure 7.3), thelower is in B. Here we see the emission line object appear roughly at −1.5′′, while there is no traceof it in the continuum. This effect is occurring at all hydrogen lines of the Balmer series. The twographs at the right in figure 7.4 display the result of the same procedure, but now at the [Oiii] 500.7nm line. In these plots, the level of the profile at 500.7 nm (black) is divided by 25 with respect tothe red one; its pixel value corresponds to the scale at the top of the graph. We see that the excessof the emission line object relative to the emission at our source is even stronger for this oxygen line.

Figure 7.2 may help to identify the other sources such as the emission line object. The imageis taken with an Hα filter. At this wavelength the same effect occurs, so the emission line sourceshould be bright on this image. Indeed we see a bright spot between the primary source and thebright star or cluster in the South East. However, this spot is weaker in Hα, while it should bebrighter judging from the spatial profiles (figure 7.4). We could not yet identify the emission lineobject with certainty.

7.2. OTHER OBJECTS 49

Figure 7.3: This image shows a part of 2D UVB spectra from 27/09/09 around Hδ (410.2 nm). Upper panel is withthe source in nodding position A, lower with source in B. Contour plots show the borders of pixel value 6, 10, 25 and40. The spectral line Hδ is in emission at the source location, but also in the rest of the slit. The contours show thatthe line is strongest at a position which shows no continuum emission. This is not only true for the hydrogen emissionlines, but for example also for the [O iii] 500.7 nm forbidden line. Probably an emission line object is located here.Around Hδ we see dark line wings. These are the result of the solar spectrum entering the slit via the moon.


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Figure 7.4: The spatial profiles at the lines Hδ (410.2 nm; left) and [Oiii] (500.7 nm; right) in black, and an averageof the the continuum between 410.5 and 433.3 nm in red, where there are no prominent emission lines. The upperfigures are in nodding mode A, whereas the lower figures are in nodding mode B. The profile on [Oiii] (black curve inright panels) is divided by 25 and corresponds to the upper flux scale in the two right graphs. The vertical scale is inpixels along the slit: 1 px = 0.1′′ ∼= 0.97 pc. We see from the profile of Hδ that there probably is a region centred atour source, with a radius of approximately 20− 25 pc. The source that is visible in the lower panels at roughly pixel40 is an emission line object: it has very strong line emission, but is not distinguishable in the continuum.

81D spectra

8.1 Object spectrum

Figures 8.1 to 8.3 show the complete normalised spectrum of the object (obj, see section 6.2). TheUVB part ranges from 300 to 550 nm (figure 8.1) and the VIS from 550 to 1020 nm (figures 8.2;8.3). We do not show the NIR (1000− 2500 nm) since the signal to noise ratio was not high enoughto do a proper analysis (SNR < 6). We re-binned all spectra to a step-size of 0.06 nm to reduce thenoise. At 310 nm, the resolution element ∆λ = 0.05 nm is somewhat smaller than this step-size,but for λ & 372 nm it is higher: ∆λ = 0.1 nm at 780 nm and ∆λ = 0.13 nm at 1020 nm (seetable 6.2 for the obtained resolving power). Binning too many data points, even if the step-size isstill smaller than the spectral resolution, may not optimally sample the line profiles, especially forthe more narrow lines. Binning too few point can make spectral features difficult to detect becauseof the high noise level.

The wavelength scales of the spectra in figures 8.1 to 8.3 are shifted to let them correspond tothe rest frame of the source. Based on measurement of the centres of the photospheric He i lines at414.38, 438.79 and 413.08 nm, the velocity shift is 195± 18 km s−1. For this velocity determinationwe used only photospheric lines that are not contaminated by nebular emission. The nebular emis-sion lines are much more pronounced and their central wavelength can be determined with a higheraccuracy, but the nebula (or nebulae in the line of sight) may have a different radial velocity.

The UVB object spectrum starts at 300 nm, but the first ∼ 40 nm are practically unusable forscientific analysis due to the high noise level. The signal-to-noise ratio in the range 310− 330 nm isonly 3.58. From 340 nm, the SNR becomes > 10 and it increases until ∼ 470 nm (see table 6.4 forSNR for different ranges). The UVB part shows a lot of lines with a peculiar profile: first a narrowpart in absorption, then a narrow part in emission, sometimes superimposed on a broader spectralline. These are the residuals of the subtraction of nebular emission lines. We have seen in the 2Dspectra, that there is nebular emission along the whole slit, and therefore also in our sky spectrum.In every line the residual profile starts with absorption and ends with emission. This suggests thatthe nebular lines at the location of our target have a slightly longer wavelength than the lines of thenebula we subtracted. This can be caused by a real radial velocity difference between those regions,or it can be an instrumental effect. Such an artefact could occur when the flexure is not correctedoptimally.

By sky subtraction, the solar spectrum from the moonlight is removed, as well as the telluricemission. However, atmospheric absorption features are still visible, for instance at 760 nm.


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8.1. OBJECT SPECTRUM 53

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Figure 8.2: Normalised VIS object spectrum from 540 nm to 780 nm.


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Figure 8.3: Normalised VIS object spectrum from 780 nm to 1020 nm. Dips occur at 839.8, 883.8, 932.7 and 987.4nm; the maximal wavelengths of the VIS orders 20, 19, 18 and 17 (see X-Shooter User Manual). These are residualsfrom the order merging.

8.2. NEBULAR SPECTRUM 55

8.1.1 Comparison with other observations

In figures 8.4, 8.5 and 8.6 we focus on a part of the UVB spectrum. This range between 380 and490 nm is traditionally used for spectral classification. In these figures, we compare the observednormalised spectrum with that of spectral standard stars (Walborn & Fitzpatrick, 1990). Theseare observations of galactic O-stars with a resolution of 0.2 nm and a SNR of maximally 80. Fig-ure 8.4 shows spectra of O7 and O8 type stars of various luminosity classes. Figure 8.5 comparesour spectrum with luminosity class V objects: main sequence stars in the temperature range O3- O9.5. Figure 8.6 shows spectra that have been classified with luminosity class I: supergiants ofspectral type O3 until O8.5.

In order to compare our observations with the spectral standard stars, we added artificiallygenerated noise to the standard spectra up to the SNR level of our observations (∼ 18 in thisrange). The result is shown in the lower panel of figure 8.6. Our starting point is the classificationfrom Castro et al. (2008): early O supergiant. In figure 8.6 we see that He iλ4471 grows stronger inabsorption for later subclasses of O-stars, while He iiλ4541 diminishes (see also section 9.2.1). Inour spectrum we see He iλ4471, but no He iiλ4541. A faint component of the He ii line could behidden in the noise. Based on the noise-added spectra, the spectral type of our object must thenbe later than O7.5. According to Martins et al. (2005), the star would be cooler than ∼ 34000 K. Aquantitative estimate of the effective temperature based on line profiles will be made in section 9.2.1.

Another striking difference between our spectrum and the standard spectra is He iiλ4686. Wesee that supergiants of some kinds have emission here, but those lines are not as broad as what weobserve. This will be analysed in detail in section 9.2.5.

8.2 Nebular spectrum

Figures 8.7 to 8.9 again show the complete spectrum from the UVB and the VIS arm, but now withfocus on the nebular emission. This is the raw spectrum (see section 6.2), which shows the nebularlines as they have been observed. The disadvantage is that the solar spectrum and as well as thetelluric emission lines are still visible, because the spectrum is not sky-subtracted. The telluricabsorption is also present as in the object spectra (figures 8.1 to 8.3). Because this spectrum is theresult of the combination of spectra of different nights with different barycentric corrections (seetable 5.1), the telluric lines are broader than in the raw spectra for the separate nights. The maximaldifference in shift due to barycentric correction between the August and September observationsis 18 km s−1; telluric lines do not shift and are therefore slightly displaced in the barycentricallycorrected spectra.

To show the very strong nebular emission lines, we plotted the normalised nebular spectrum ona logarithmic scale. The nebular emission lines are labelled: red for forbidden lines, blue for theBalmer lines, green for the Paschen lines and pink for He i lines. The dashed-dotted cyan lines areemission lines from earth’s atmosphere. This was concluded from their homogeneous distributionof light along the slit in the 2D spectra. We assume that all nebular lines have an inhomogeneousdistribution along the slit. More of the lines which are not labelled as such might also be telluricemission lines. The spectra of figure 8.7 to 8.9 have been shifted such that the nebular lines are attheir rest wavelengths. This radial velocity (145± 5 km s−1) is lower than the one derived for thecentral star (195± 18 km s−1).


3 4 5 6 7 8 9 10 3

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f))

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(f)

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))

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ηH

ζC

aIIH

εH

eI

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HeI

I

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HeI

HeI

I

NII

I HeI

IHeI

Hβ

Figure 8.4: Part of the normalised object UVB spectrum between 380 and 490 (black) compared with Walborn &Fitzpatrick (1990) spectral standard stars of type O7 and O8 with luminosity classes I , II , III and V (red).


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I HeI

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Figure 8.5: Part of the normalised object UVB spectrum between 380 and 490 (black) compared with Walborn &Fitzpatrick (1990) spectral standard stars of type O3 to O9.5 luminosity class V (red).


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CaII

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NIV

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Hβ

Figure 8.6: Part of the normalised object UVB spectrum between 380 and 490 (black) compared with Walborn& Fitzpatrick (1990) spectral standard stars of type O3 to O8.5 luminosity class I (red). In the lower panel weshow the result of adding artificially generated noise to the standard spectra, to match the signal-to-noise ratio of ourobservations (SNR ∼ 18 in this range).


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wavelength (nm)

Figure 8.7: Normalised UVB raw spectrum from 500 nm to 540 nm on log scale, showing the nebular emission, butalso telluric lines and solar spectrum via moonlight. The nebular emission lines are labelled: red for forbidden lines,blue for the Balmer lines, green for the Paschen lines and pink for He i lines.


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Figure 8.8: Normalised VIS raw spectrum from 540 nm to 780 nm on log scale, showing the nebular emission, butalso tellluric lines and solar spectrum via moonlight. The nebular emission lines are labelled: red for forbidden lines,blue for the Balmer lines, green for the Paschen lines and pink for He i lines. The lines that have been identified withcertainty as telluric emission lines based on the spatial profile along the slit are indicated with dashed-dotted cyanlines.


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wavelength (nm)

Figure 8.9: Normalised VIS raw spectrum from 780 nm to 1020 nm on log scale, showing the nebular emission, butalso telluric lines and solar spectrum via moonlight. The nebular emission lines are labelled: red for forbidden lines,blue for the Balmer lines, green for the Paschen lines and pink for He i lines. The lines that have been identified withcertainty as telluric emission lines based on the spatial profile along the slit are indicated with dashed-dotted cyanlines.

Part IV

Scientific Analysis

9Central star NGC 55 C1 31

In this part we will derive constraints on properties of NGC 55 C1 31 and its surroundings . Insection 9.1 we discuss the luminosity of the source. In section 9.2 we compare line profiles frommodel atmospheres with our observed photospheric spectrum to derive constraints on effectivetemperature, mass loss and rotational velocity. The possibility of having a composition of more thanone source is worked out section 9.2.6. The conclusions on the central source will be summarised insection 9.3.

In chapter 10 we discuss the surrounding nebula, and derive properties such as electron tem-perature and electron density from emission line ratios. Finally, in chapter 11 the properties ofthe nebula will be compared with models, taking into account the properties of the central ionisingsource. We will combine the results in an attempt to create a consistent picture of the central sourceand it surroundings.

9.1 Luminosity

The V apparent magnitude mV of NGC 55 C1 31 is 18.523 (Castro et al., 2008). There is extinctionfrom our own galaxy as well as from NGC 55. The foreground reddening is low due to the highgalactic latitude (b = −75◦) of NGC 55: E (B − V )mw = 0.013 (Schlegel et al., 1998). However,since we see the NGC 55 galaxy almost edge-on, the internal reddening is expected to be muchhigher: E (B − V )int = 0.127 (Gieren et al., 2008).

The difference between apparent magnitude mV and absolute magnitude MV is caused by dis-tance and by extinction AV = E (B − V ) RV :

(m−M)V = 5 log(

d

10 pc

)+ E (B − V )mw RV,mw + E (B − V )int RV,int

�� 9.1

Taking RV,mw = 3.10 and RV,int = 3.24 (Gieren et al., 2008), and a distance d of 2.0 Mpc (seesection 2.1), we find an absolute visual magnitude MV of −8.43.

According to Martins et al. (2005), even supergiants do not have absolute visual magnitudeslower than MV = −6.35. The obtained value of MV for this source suggests that we are looking atmultiple sources.

MV is translated into the bolometric magnitude applying the bolometric correction BC, whichis −2.9 for objects with Teff = 30000 K1. This yields a luminosity log L/ L� = 6.43. This either

1This choice of temperature is based on our comparison with spectral standard stars is section 8.1.1. It will turnout in section 9.2.1 that this is indeed a good estimate for the average effective temperature that we measure.

64 CHAPTER 9. CENTRAL STAR NGC55 C1 31

implies an extremely luminous object, or a small cluster of stars. To put this luminosity in per-spective, the most luminous star in the Milky Way is η Carinae with a luminosity log L/ L� = 6.7.According to Martins et al. (2005), O-type supergiants have log L/ L� = 5.5− 6.0. The projectedwidth of the slit used in the observations corresponds to a region with a physical diameter of 7.8pc (UVB slit width 0.8′′). This would, for example, be large enough to contain the entire OrionTrapezium Cluster. In this galactic cluster, at least five bright massive stars are concentrated ina region of only half a parsec. Most, if not all stars are formed in clusters (Lada & Lada, 2003)and especially massive stars are usually found in clusters. This too supports the hypothesis that wehave observed multiple sources instead of a single star.

If we have indeed observed multiple stars this severely complicates the determination of stellarparameters from line profiles, because there is an almost infinite number of free parameters. There-fore, we are first going to see whether we can reproduce most of the line profiles with a single starmodel, while we ignore V (sections 9.2.1 to 9.2.3). After that, we see if we can modify the bestfitting model to let it reproduce V . Then in section 9.2.5 we will have a closer look at He iiλ4686,which is a very peculiar line in our observations and very difficult to fit. It will turn out that weindeed need multiple stars. The procedure of this exercise is described in section 9.2.6. We willsummarise our conclusion on the central source in section 9.3

9.2 Modeling line profiles: Teff, M and vrot sin(i)

The profiles of spectral lines are responsive to various stellar parameters such as effective surfacetemperature, mass loss, surface gravity, chemical abundances, rotation speed, etcetera. The effectsof parameters on the line profiles can be explored by modelling stellar atmospheres and examiningthe predicted profiles.

We use FASTWIND (Puls et al., 2005) to model stellar atmospheres and to compute the profilesof spectral lines formed in these atmospheres. FASTWIND is able to model non-LTE line blanketedstellar atmospheres and is thus suited to model stars with strong winds. In the first grid of models(MOD01 to MOD27 in table A.1) we vary three parameters: effective temperature Teff, luminosityL, and mass loss rate M , with three different values each, resulting in 27 models. Teff = 30000,32500 and 35000 K, log L/ L� = 5.4, 5.6 and 5.8 and M = 1.0×10−6, 3.0×10−6 and 1.0×10−5 M�yr−1. All models in this grid have M = 40 M�

2 The stellar radius R follows from Teff and L usingL = 4πR2σT 4

eff, with σ the Stefan-Boltzmann constant. The surface gravity g is computed fromR and M using R = GM/R2, with G being the gravitational constant. Furthermore, the rate ofacceleration of the wind is characterised by β = 1 and v∞ = 2.6vesc with vesc =

√2GM/R the

escape velocity at the stellar surface. We ignore the correction factor due to the electron scattering(see section 1.2).

The abundances for the models computed in this study except MOD60, MOD61, MOD70 andMOD77 are X = 0.710 and Y = 0.284. This implies Z = 0.006 = 0.5 Z�. A turbulent velocityfield was added to the atmosphere, characterised by vturb = 10 km s−1. No clumping of the wind istaken into account.

Models MOD28 to MOD36 vary the temperature, and the luminosities as in the default grid, but

2This is a canonical value which we do not vary to restrict the number of free parameters. According to Martinset al. (2005), the spectroscopic masses of O-type supergiants vary between 30 and 67 M�. The mass will onlyinfluence the line profiles via g and v∞.

9.2. MODELING LINE PROFILES: TEFF, M AND VROT SIN(I) 65

now with a fixed mass loss of M = 6.0× 10−6 M� yr−1. Models MOD37 to MOD40 are variationson MOD34, providing an extension of the range of the effective temperature down to 25000 K andsome steps in between. MOD61 to MOD69 vary the wind parameters β and v∞. MOD60, MOD61and MOD70 are used to probe other abundances (X = 0.400 and Y = 0.600) in combination withlarger luminosities and mass loss rates. All relevant model parameters can be found in the tablesin appendix A.

The computed line profiles that will be shown in the following section have been convolvedwith appropriate functions to simulate the resolving power of the instrument and the rotationalvelocity of the star. Unless indicated otherwise, the projected rotational velocity vrot sin(i), with ithe inclination angle, is 150 km s−1. For the resolving power we have taken the theoretical estimatesdescribed in the instrument manual: 6200 for the UVB arm with slit width 0.8” and 8800 for the VISarm with slit width 0.9”. The convolutions were carried out with help of IDL procedures especiallydesigned for FASTWIND output.

9.2.1 Constraints on effective temperature

The ratio of equivalent widths of the absorption lines He iλ4471 and He iiλ4541 defines the sub-classification of O-stars, see table 9.1 (Conti & Alschuler, 1971). The He i line will be stronger incooler, late O-stars while the He ii line is more pronounced in the hotter, early O-stars. This isbecause a larger fraction of the Helium in the atmosphere will be ionized. This effect is also clearlyvisible in figure 8.5: we see He iλ4471 increase and He iλ4541 decrease in the spectra from earlydown to late type O-stars of luminosity class V. The same effect is visible for the luminosity class Iobjects in figure 8.6 though one has to keep in mind that a supergiant of a certain spectral type is∼ 2000 degrees cooler than a main sequence star of the same spectral type (Martins et al., 2005).Although this pair of lines is the primary criterion for the O-type sub-classification, there are morephotospheric lines that are sensitive to effective temperature.

He iiλ4541 is not visible in the NGC 55 C1 31 spectrum, but a faint component of this line couldbe hidden in the noise: the signal to noise ratio is ∼ 24 in this part of the spectrum (see table 6.4).When the spectrum is integrated over the interval where the line should be located, one measuresan equivalent width of 0.016 ± 0.013 nm, which implies that it can indeed be absent. The error isdue to the uncertainty in determining the noisy continuum level.

Table 9.1: Spectral types based on He iλ4471/He iiλ4541 equivalent width ratio W ′. Table from Conti & Alschuler(1971).

log W ′ Limit Type Limit log W ′ Limit Type Limit

O9.5 ≥ + 0.45 0 > O7 ≥ − 0.10+ 0.45 > O9 ≥ + 0.30 − 0.10 > O6.5 ≥ − 0.20+ 0.30 > O8.5 ≥ + 0.20 − 0.20 > O6 ≥ − 0.30+ 0.20 > O8 ≥ + 0.10 − 0.30 > O5.5 ≥ − 0.50+ 0.10 > O7.5 ≥ 0 − 0.50 > O5 ≥ − 0.70

− 0.70 > O4 ...


Figure 9.1 shows, along with the observed spectrum, four computed line profiles. These areMOD31, MOD37, MOD32 and MOD33 (see tables in appendix A), which only differ in temper-ature: Teff is respectively 30000, 31000, 32500 and 35000 K. The line becomes stronger in hotteratmospheres. The computed equivalent width is respectively 0.021, 0.029, 0.040 and 0.054 nm.MOD37 (Teff = 31000 K) is just weak enough to fall within the error bars of the observed value of0.016± 0.013 nm, which means that, based on this line, the effective temperature can not be higherthan 31000 K.

The strength of the He iλ4471 line is also dependent on temperature, but it does not varystrongly in the range 30000 - 35000 K. It is, however, useful to examine the line profiles of thesame models as used for He iiλ4541, and to compare them with the observed spectrum. The lowerpanel of figure 9.1 shows the profile of He i λ4471 in the observed spectrum together with modelsMOD31, MOD32 and MOD33 (30000, 31000 and 35000 K). We see that the line is slightly strongerin cooler atmospheres, but the difference is small. The shape of the profile in the observed spectrumis uncertain, because we have a residual of a subtracted nebular emission line. The equivalent widthis measured 0.12± 0.02 nm, where we assume that the residual adds positively as well as negativelyto the integral of the line. These three models all fit the observed spectrum and no additionalconclusions can be drawn. It is in agreement with the hypothesis that Teff ≤ 31000 K.

We can divide the equivalent widths of the He i and He ii line to obtain an estimate of the equiva-lent width ratio W ′. Taking into account the errors in the measurements, we find log W ′ = 0.88+0.26

−0.77.According to table 9.1, the star is an O8 star or cooler.

He iλ4387 is clearly visible in the observed spectrum, and has the advantage of not being con-taminated by nebular emission. The line is not present in the spectra of early O-stars, but becomesstronger in later type stars. Figure 9.2 shows the computed line profiles for this line, along withthe observed spectrum. The models MOD40, MOD39, MOD31, MOD32 and MOD33 have effectivetemperatures of respectively 25000, 27500, 30000, 32500, 35000 K. The strength of the absorptionline varies significantly with temperature in this range. The line profiles with Teff ≤ 30000 K allfit the observed spectrum. This confirms our earlier hypothesis that the effective temperature isnot higher than 31000 K. Because the line is only marginally sensitive to temperature at lowertemperature, we can not use it to obtain a lower limit on the temperature.

He iiλ4200 is, like He iiλ4541, apparently absent. When we measure the equivalent width atthe line location, we find 0.049 ± 0.300 nm. The computed line profiles are plotted in figure 9.3.The absorption lines produced from models MOD31, MOD37, MOD32 and MOD33 have equivalentwidths of respectively 0.019, 0.025, 0.033 and 0.046 nm. These values all fall within the error estimateof the measured equivalent width in the observed spectrum. No conclusions on the temperature canbe drawn from this line.

9.2.2 Constraints on rotational velocity

Because the observed line profile of He iλ4387 is very pronounced and uncontaminated, we can alsouse it to obtain an estimate of the rotational velocity of the star. The line profiles of MOD31 has beenconvolved with appropriate functions to simulate three different rotational velocities: vrot sin(i) =50, 150 and 250 km s−1 (measured at the surface of the star).The lines with vrot sin(i) = 50 km s−1 are clearly not broad enough to fit the observed profile, whilethe lines with vrot sin(i) = 250 km s−1 appear to be too broad. This confirms that our originallyadopted default value for vrot sin(i) was a good estimate; the projected rotational velocity of the


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data

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MOD33: T=35000 KMOD32: T=32500 KMOD31: T=30000 K

observed spectrum

Figure 9.1: Upper panel: The observed line profile of He iiλ4541 along with four predicted line profiles from models.The models have equal parameters except for Teff, which takes values of 30000, 31000, 32500 and 35000 K. The otherparameters of the models can be found in table A.1. The equivalent width of the apparently absent line is measured0.016± 0.013 nm, the equivalent widths of the models MOD31, MOD37, MOD32 and MOD33 are respectively 0.021,0.029, 0.040 and 0.054 nm. These diagnostics imply that the source must have Teff ≤ 31000 K. Lower panel: Theobserved line profile of He iλ4471 along with the predicted profiles from MOD31, MOD32 and MOD33 which differin effective temperature (30000, 32500 and 35000 K respectively). The strength of the absorption line is a function oftemperature, but not very strongly in the range 30000 − 35000 K. The observed absorption line is contaminated bynebular emission, which makes it difficult to examine the profile. These three model profiles are all consistent withthe observations.


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MOD31: T=30000 K vsini= 50 km/sMOD31: T=30000 K vsini=150 km/sMOD31: T=30000 K vsini=250 km/s

observed spectrum

Figure 9.2: The observed and computed profiles of He iλ4387. MOD40, MOD39, MOD31, MOD32 and MOD33respectively have effective temperatures of 25000, 27500, 30000, 32500 and 35000 K. This line does not suffer fromnebular contamination. In the upper panel we see that the strength of the line is a function of temperature, but thatsensitivity is almost lost at temperatures below 30000 K. Therefore we can conclude that Teff ≤ 30000 K, which is inagreement with the conclusion based on He iiλ4541 (Teff ≤ 31000 K). In the lower panel we show the profile of MOD31(30000 K) convolved with appropriate functions to simulate different projected rotational velocities. It appears thatv sin(i) = 150 km s−1 is the best fit for the observations.


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MOD31: T=30000 KMOD37: T=31000 KMOD32: T=32500 KMOD33: T=35000 K

observed spectrum

Figure 9.3: The observed spectrum at the interval where He iiλ4200 should be located, along with profiles frommodels which only differ in temperature. MOD31, MOD37, MOD32 and MOD33 respectively have effective temper-atures of 30000, 31000, 32500 and 35000 K. The computed equivalent widths are 0.019, 0.025, 0.033 and 0.046 nm,whereas the measured value is 0.049 ± 0.30 nm. All predicted values fall within the error estimate of the measuredvalue, therefore no additional conclusions can be drawn based upon this line profile.


star must be close to 150±30 km s−1. The error estimate is based on the comparison of the profilesfor more different rotational velocities.

9.2.3 Constraints on mass loss

In our observed raw spectrum (see figure 8.8), Hα (λ0 = 656.28 nm) is very strong in emission.This is most probably due to the surrounding H ii region, and possibly ionised regions in the line ofsight in front of the star. With sky-correction we perfectly subtract the solar spectrum, which hasan absorption profile at this wavelength. As a result, the Hα line wings, which are probably formedin the wind/atmosphere of the star, become visible. Still, the centre of Hα (∆λ < 1 nm) can notbe used for scientific analysis, since this only shows the residual of the subtraction of the nebularemission line.

The profile of Hα appears to be extremely sensitive to mass loss flux M/R2. In the upper panelin figure 9.4 we plot the object spectrum together with three line profiles from models: MOD12,MOD31, MOD11 and MOD10. These models only differ in mass loss rate: 1 × 10−5, 6 × 10−6,3 × 10−6, and 1 × 10−6 M� yr−1 respectively. Teff = 30000 K and R = 23.40 R� for all threemodels. Depending on the mass loss, the line can be either in emission or in absorption. Sincethe computed profile for MOD31 provides the best fit we conclude that the mass loss of the starmust be close to 6 × 10−6 M� yr−1 for a star with Teff = 30000 K and a radius of 23.40 R�. Topreserve the profile for a model with a higher luminosity at fixed Teff and therefore a larger radius,the mass loss rate should be assigned a higher value. Mass flux M/R2 is proportional to M/L, soin order to preserve the shape of the profile, the mass loss rate should be equally increased by thesame fraction as the luminosity. The luminosity and the mass loss rate of MOD72 are increased bythe same factor (6.8) with respect to the parameters of MOD31, in order to preserve the mass flux.MOD72 has a luminosity of 106.43 L�, which is in agreement the observed V , and a mass loss rateof 4.07 × 10−5 M� yr−1, whereas MOD31 has L = 105.6 L� and M = 1 × 10−6 M� yr−1. Bothhave Teff = 30000 K. In the lower panel of figure 9.4 it can be seen that they have roughly the sameprofile.

Concluding, the mass loss flux of the source has a value of 1.10× 10−8 M� yr−1 R−2� . For our

best fit model MOD31, with R = 23.40 R�, the mass loss rate should be 6× 10−6 M� yr−1.

9.2.4 Accordance with observed magnitude

Up to here, most of the observed line profiles are well reproduced with MOD31, except for theluminosity, which is not in agreement with V (see section 9.1). MOD31 represents an objectwith an effective temperature of 30000 K, a luminosity of L = 105.6 L�, a mass loss rate ofM = 6 × 10−6 M� yr−1 and a projected rotational velocity of 150 km s−1 (measured at thesurface). To reproduce V , we need a luminosity L = 106.43 L� for an object with effective tem-perature of 30000 K. Let us modify MOD31 such that we conserve Teff, M/R2 and g, the surfacegravity3, while the luminosity reproduces V . This would probably result in an artificial object, butwe can still model it to see what the line profiles would be.

We will use the effective temperature of MOD31, the best fit model, which is 30000 K. In orderto emit a luminosity L = 106.43 L�, the star needs to have a radius of 60.84 R�. Preserving surfacegravity g is achieved by setting the stellar mass to 270 M�. To keep the mass flux equal to that of

3We have not probed g individually because the quality of our data is too low, but preserve it to be sure we seeno effect from changing it.


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MOD12: 1 * 10-5 Msun yr-1

MOD31: 6 * 10-6 Msun yr-1

MOD11: 3 * 10-6 Msun yr-1

MOD10: 1 * 10-6 Msun yr-1

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MOD72: 4.07 * 10-5 Msun yr-1, R=60.84 Rsun L=106.43 LsunMOD31: 6.00 * 10-6, Msun yr-1,R=23.40 Rsun, L=105.6 Lsun

observed spectrum

Figure 9.4: Upper panel: The line profile of Hα in the observed spectrum compared with the computed profilesfor MOD12, MOD31, MOD11 and MOD10. These models have mass loss rates 1 × 10−5, 6 × 10−6, 3 × 10−6, and1 × 10−6 M� yr−1, respectively. All models are equal in effective temperature (Teff = 30000 K) and luminosity(L = 105.6 L�); therefore they also have the same radius (R = 23.4 R�). A mass loss rate of 6 × 10−6 M� yr−1

provides the best fitting line profile. Lower panel: The comparison of two line profiles of models that have the samemass loss flux M/R2 ∝ M/L. MOD72 has a luminosity of 106.43 L� and a mass loss rate of 4.07 × 10−5 M� yr−1,whereas MOD31 has L = 105.6 L� and M = 6 × 10−6 M� yr−1. The line profiles, especially the line wings, aresimilar.


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Hα MOD31, MOD72, MOD73

MOD31 T=30000 K L=105.6 LsunMOD72 T=30000 K L=106.43 LsunMOD73 T=25000 K L=106.33 Lsun

observed spectrum

Figure 9.5: The observed line profile of Hα together with the predicted profiles for MOD31, MOD72 and MOD73.MOD31 provided the best fitting profile. MOD72 mimics this model in mass flux, surface gravity and temperature buthas a luminosity L = 106.43 L� that would result in the measured V . MOD73 does the same, but at a temperatureof 25000 K instead of 30000 K. Despite the fact that the mass flux is conserved, the line profile for MOD73 does notreproduce the observed profile. Apparently, the strength of the line is not only a function of mass flux but also oftemperature.

MOD31, M will be 4.07× 10−5 M� yr−1. This will be MOD72.

The line profiles of Hα from MOD72 is in agreement with the observations, though the profileis a bit different than that of MOD31; see figure 9.5. The lines He iiλ4541 and He iiλ4200 of thismodel have equivalent width of respectively 0.025 and 0.021 nm, which are inside the error bars ofthe apparently absent lines in the observed spectrum. He iλ4471 fits the observed profiles as goodas MOD31 (upper panel figure 9.6). The profile of He iλ4387 is too shallow to fit the data (lowerpanel figure 9.6). The trend in figure 9.2 suggests that we need a lower effective temperature forthis high-luminosity model. Therefore we introduce MOD73. This model has an effective temper-ature of 25000 K. With a slightly lower bolometric correction of −2.65 for this temperature, weneed a luminosity of L = 106.33 L� to reproduce the observed V . R = 78.08 R�, and to preserveg, we need M = 440 M�. The mass loss needs to be 6.71×10−5 M� yr−1 to preserve the mass flux.

With MOD73 we see indeed a better fit for He iλ4387, see lower panel figure 9.6. However,despite the conservation of the mass flux, Hα does not fit anymore for MOD73 (figure 9.5). Appar-ently, the strength of the line is not only a function of mass flux but also of temperature. We couldwork around this with lowering the mass loss rate. This might result in model that reproduces allline profiles that we analysed up till now, but we have not yet treated He iiλ4686.

9.2.5 He ii λ4686

There is another emission line visible in our observed spectrum: He iiλ4686, which has such a pe-culiar shape that we analyse it separately. Figure 9.7 shows that this profile is not al all reproduced


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Figure 9.6: The observed and predicted line profiles of He iλ4471 (upper panel) and He iλ4387 for MOD31, MOD72and MOD73. MOD31 provided the best fitting profile. MOD72 mimics this model in mass flux, surface gravity andtemperature but has a luminosity L = 106.43 L� that would result in the measured V . MOD73 does the same, but ata temperature of 25000 K instead of 30000 K. He iλ4471 is not significantly different for these three models; all modelsare in agreement with the observed profile. Based on He iλ4387 (lower panel) we earlier concluded Teff ≤ 30000 K,but for the high luminosity models, we see that we need a lower temperature: MOD73, with Teff ≤ 25000 K, providesa better fit than MOD72 with Teff = 30000 K.


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HeII λ4686 MOD31, MOD72, MOD73


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Figure 9.7: The observed line profile of He iiλ4686 together with the predicted line profiles for MOD31, the bestfitting model for all other lines and MOD72 and MOD73 that mimic MOD31, but do also reproduce the observed V .None of these models even slightly reproduces the strength or broadness of the emission profile.

by MOD31, MOD72 or MOD73, or actually by any other model computed so far. The observedemission line is very broad (∼ 3000 km s−1) but has an equivalent width of only −0.36± 0.04 nm.Figure 9.8 shows the observed spectrum around this feature. We see that the feature is visible in theraw (green) and obj (black) spectrum, but not in the sky. In this interval, the sky contribution isonly moon light and there are no nebular lines, therefore we can be sure that the feature is real (i.e.not the result of the subtraction of sky from raw to produce obj, see section 6.2). Furthermore, thefeature is also visible in this source’s spectrum in Castro et al. (2008). This strengthens the ideathat this is not some instrumental artefact or an unintended result of the data reduction.

Stars with strong stellar winds like early-O type supergiants, show He iiλ4686 emission lines, ascan be seen in figure 8.6. Figure 8.4 shows that the line turns from a sharp emission feature for aluminosity class I object to an absorption line for a main sequence star (luminosity class V). AlsoWolf-Rayet stars are known to show strong He iiλ4686 emission, including WN/O stars that arelikely very massive main sequence stars that already suffer from strong mass loss.

To better undtand the observed profile we are going to analyse which parameters have thestrongest influence on this line, and what we need to change in order to reproduce this profile. Infigures 9.9 to 9.12 we show the observed profile of He iiλ4686 along with models in which we varyone parameter at the time in order to see the effect on the shape and strength of the line. The redline in these plots indicates the result for MOD31 with a rotational velocity vrot sin(i) = 150 km s−1,which was the best fitting model for all the other spectral lines.


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total raw rebinnedtotal sky rebinned

total object rebinned

Figure 9.8: This figure shows a part of the UVB spectrum between 455 and 480 nm. The green spectrum is extractedat the source (raw), the red spectrum is an average of various patches of sky, black is the resultant obj spectrum,created from subtracting the sky from the raw spectrum. He iiλ4686 emerges as a broad (∼ 3000 km s−1) feature inthe obj spectrum, while there is nothing remarkable in the sky spectrum. While other lines are heavily contaminatedby nebular emission, this one is not. Therefore we can be sure this feature really belongs to the source’s spectrum.This figure also shows that the feature at the left side of He iiλ4686 probably is a residual of the subtraction of raw

and sky, since it is stronger in those spectra.


Mass loss and effective temperature

Figure 9.9 shows that the line strength increases with mass loss rate and temperature. The line isformed in the wind and will be stronger for higher mass loss rate. Higher effective temperaturesincrease the line strength because a larger fraction of the helium in ionised. However, the highvalues, for which He iiλ4686 is strong enough, are not in agreement with the results from the otherlines we examined in the previous sections.

Helium abundance

We could also alter the helium abundance in the star in order to simulate a Wolf-Rayet star. Thesestars usually show strong emission at He ii λ4686. MOD60 and MOD61 and MOD70 are differentfrom the rest of our grid, since they have X=0.398 and Y=0.596 while the others all have X=0.710and Y=0.284 (solar values). MOD70 is exactly like the preferred MOD31, only with enhanced he-lium abundance. Furthermore, MOD60 has M = 40 M�, M = 1.5 × 10−5 M� yr−1, L = 106 L�and v∞ = 2000 km s−1, whereas MOD61 has M = 100 M�, M = 3× 10−5 M�yr−1, L = 106 L�and v∞ = 2637.98 km s−1. The computed profiles for He iiλ4686 are shown in figure 9.10. We seethat just enhancing the helium abundance (MOD70) does not change the strength and shape ofthe line very much, though there is a small effect. The reason why the profiles of models MOD60and MOD61 are stronger than MOD31 and MOD70 is probably due to the higher mass loss rates,an effect we have already seen in figure 9.9. Increasing the helium abundance will make all heliumlines, created by neutral as well as ionised helium, stronger. For solar abundances we were able tofind a temperature (30000 K) at which all helium lines more or less match the observed profiles.With enhanced helium abundance, this is not possible. Only increasing the helium abundance doesnot seem to be the right step.

Wind parameters

Another option is to change wind parameters: v∞, the terminal wind velocity and β, the parametermeasuring the rate of acceleration of the wind. The models up to MOD40 all use v∞ = 2.6 vesc

and β = 1. We are probably able to change the strength and profile of wind line He iiλ4686 with-out affecting the photospheric lines too much. The upper panel of figure 9.11 shows the result ofchanging β. It is clear that increasing β makes the line stronger. In terms of equivalent width, avalue of β = 3 would be best: the computed equivalent width for MOD69 is 0.358 nm, matchingour measured value of 0.36± 0.04 nm. All other parameters of MOD69 are equal to MOD31.In the lower panel of figure 9.11 we show the effect of combining two different values of β: 1 and2, with two different values for v∞: 2100.4 (which is v∞ = 2.6 vesc) and 3000.0 km s−1. The linestrength decreases for higher values of v∞ because the density in the line forming part of the windis decreased.

Rotational velocity

Though the equivalent width of some models may fit the measured value, it is obvious that we donot reproduce the broadness of our observed line. In an attempt to resolve this we simulate a higherrotational velocity. The upper panel of figure 9.12 shows the computed line profile of MOD69,which had a matching equivalent width, but now with four different projected rotational velocities:vrot sin(i) = 150, 600 and 900 and 1200 km s−1. To reproduce the measured profile, we would need


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HeII λ4686 varying mass loss

MOD12: 1 * 10-5 Msun yr-1

MOD31: 6 * 10-6 Msun yr-1

MOD11: 3 * 10-6 Msun yr-1

MOD10: 1 * 10-6 Msun yr-1

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MOD33: T=35000 K, 6 * 10-6 Msun yr-1

MOD32: T=32500 K, 6 * 10-6 Msun yr-1

MOD37: T=31000 K, 6 * 10-6 Msun yr-1

MOD31: T=30000 K, 6 * 10-6 Msun yr-1

observed spectrum

Figure 9.9: The observed line profile of He iiλ4686 together with the predicted line profiles for models with differentmass loss rates (upper panel) and temperatures (lower panel). The red profile is from MOD31, our favoured modelfor all other lines. The mass loss rate appears to have an effect on whether the line is in emission or absorption.Fortunately, the value for which the profile of Hα had the best fit: M = 6.0× 10−6 M� yr−1 is already in emission,though it is much too weak and narrow. At higher temperature, more helium will be ionised resulting in strongerHe ii lines. Although higher temperatures seem to provide a better fit, at least in terms of equivalent width, this isnot in agreement with the temperatures derived from the strength of the other He i and He ii lines, see section 9.2.1.


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HeII λ4686 for enhanced helium abundance

MOD61: X=0.398, Y=0.596, M= 40, 1.5 * 10-5 Msunyr-1

MOD60: X=0.398, Y=0.596, M=100, 3.0 * 10-5 Msunyr-1

MOD70: X=0.398, Y=0.596, M= 40, 6.0 * 10-6 Msunyr-1

MOD31: X=0.710, Y=0.284, M= 40, 6.0 * 10-6 Msunyr-1

observed spectrum

Figure 9.10: In this figure we show the line profiles of He iiλ4686 of the three models with enhanced heliumabundance, along with the observations. The red profile is from MOD31, our favoured model for all other lines.MOD70 shows the effect of only enhancing the helium abundance, which appears to be not major. The fact thatMOD60 and MOD61 have stronger He iiλ4686 is most probably due to their high mass loss rates.

vrot sin(i) = 1200 km s−1. This is more than the escape velocity vesc ∼ 800 km s−1 of this modelstar, resulting in an unstable situation. This problem could be solved by increasing the star’s massto 90 M� while keeping the radius equal. A more important problem is that the other line profilesmight not fit any more with this extreme rotation speed. Let us check how the other line profilesrespond to a rotational velocity of 1200 km s−1.

The lower panel of figure 9.12 shows the profile of Hα for MOD69 with vrot sin(i) = 150 and1200 km s−1. The centre of this line is less reliable than the wings due to the subtraction of anarrow nebular line. Therefore, the 150 km s−1 version may be called a rough fit. But the profilefor 1200 km s−1 certainly does not reproduce the observations.

Figure 9.13 shows that it was a mistake to presume that changing only β would not affect photo-spheric lines. Even the normal 150 km s−1 ‘version’ of MOD69 does not reproduce the He iλ4387 ab-sorption line, while MOD31 did. These two models only vary in β: MOD31 has β = 1 while MOD69has β = 3. It is not surprising that the convolved profile to simulate vrot sin(i) = 1200 km s−1 doesnot leave a trace of this very weak line. Apparently the stellar wind does affect this spectral line.

We conclude that a high rotational velocity is not the explanation for the broad He iiλ4686 line,because the other lines would then have to be broad as well. That this is not the case becomes clearfrom figure 9.12, lower panel, and figure 9.2, lower panel.


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MOD69: β=3.0MOD67: β=2.0MOD65: β=1.2MOD31: β=1.0MOD62: β=0.7

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HeII λ4686 varying β and v∞

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observed spectrum

Figure 9.11: The effect of changing the wind parameters β and v∞ on the profile of He iiλ4686. Upper panel: theeffects of changing β only. All shown models have parameters equal to those of MOD31, except for β. The line profileof MOD69 (β = 3.0) matches the observations in terms of equivalent width: the computed value 0.358 nm, matchingour measured value of 0.36± 0.04 nm. However, the broadness of the line can not be reproduced by just increasing β.Lower panel: the effects of changing wind parameter β as well as terminal wind velocity v∞. Increasing the terminalwind velocity decreases the overall density of the wind, which results in a weaker emission line.


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Figure 9.12: MOD69 is like MOD31, only β = 3.0 instead of 1.0. This model reproduced the observed He iiλ4686emission line in terms of equivalent width, but not in terms of broadness. This can be resolved by increasing therotational velocity. The upper figure shows the profile of He iiλ4686 from MOD69 with rotational velocities ofvrot sin(i) = 150, 600, 900 and 1200 km s−1. The profile for 1200 km s−1 matches the observations, but this does notcorrespond to a physically stable situation: this rotational velocity is higher than the escape velocity at the surface(∼ 800 km s−1 for this model). This problem can be evaded by increasing the star’s mass to 90 M� while keeping theradius equal. However, with this extreme rotational velocity, the other profiles might not fit any more, as is shown inthe lower panel. In this figure the observed line profile of Hα is shown along with the predicted profiles for MOD69.We plotted the results for two different rotational velocities: the default value of 150 km s−1 and the value whichreproduced the profile of He iiλ4686: 1200 km s−1. While we may call the profile for 150 km s−1 a rough fit (since thecentre of the line is less reliable due to nebular line subtraction), the result for 1200 km s−1 is in strong disagreementwith the observations.


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Figure 9.13: The observed profile of He iλ4387, along with the predicted profile of MOD69 with two differentrotational velocities: the default value of 150 km s−1 and the value which reproduced the profile of He iiλ4686:1200 km s−1. MOD69 has the same values for all its parameters as MOD31, the favoured model, except for β, whichis 3. MOD31 with 150 km s−1 reproduced the profile very well (see figure 9.2, lower panel). This figure shows,however, that changing β does not leave the profile of He iλ4387 intact: the line has almost disappeared.

9.2.6 Multiple sources producing blended lines.

We already mentioned that, due to the high visual brightness we observe, it is very probable thatwe look at multiple sources. If this is indeed the case, the resulting spectrum would also be acomposition of the various spectra of the different sources, weighted by their intrinsic luminosities.

The attempts in section 9.2.5 show that it is impossible to reproduce the profile of He ii λ4686with a single object while keeping rotational velocities in a reasonable range compared to the escapevelocity. Furthermore, if we include high rotational velocities, the other line profiles are smearedtoo much and do not match the data any more. If we allow the profile to be composed of differenttypes of stars, it will be easier to produce a data-matching line profile. Besides, we can choose morereasonable values for L, because we assume there is more than one star.

Naturally, when we allow multiple components with various parameters, we should be able toproduce any line profile we like. But there are a few restrictions:

� Not just He ii λ4686, but all profiles should match the observations in the same ratio.

� The combination of models should, with the chosen ratio, reproduce V .

� We do not allow the different components to have different radial velocities which would shiftthem in wavelength. If they belong to a compact cluster, there is no reason for them to haveradial velocities that differ in the order of the broadness of the line (3000 km s−1)4.

4According to the Virial Theorem, a cluster with a mass of a few hundred solar masses and a radius of ∼5 parsechas a velocity dispersion of a few hundred km s−1.


� We want the composition to contain the lowest possible number of different components (i.e.the simplest solution).

The best fit profile is obtained via the following procedure. We constructed a model whichwould produce a very strong He ii λ4686 emission line: a high temperature of 40000 K and massloss rate of 2.02 × 10−5 M� yr−1. In order to keep the Hα emission modest, we also altered thechemical composition: X = 0.237 and Y = 0.757. Lower hydrogen abundance would result in aweaker Hα line than with solar abundances. The model resembles a Wolf-Rayet star of the nitro-gen WN sequence. (The parameters are similar to the model parameters for observed WN starsin LMC in Koesterke et al. (1991).) For simplicity, we kept the mass 40 M�. Because we willneed roughly 10% of this line profile for He iiλ4686, we choose the luminosity L = 105.43 L�, whichwould provide 10% of the observed visual magnitude. This is MOD81, of which all parameters canbe found in table A.2 in the appendix. The resulting line profiles are indicated in green in figure 9.14.

If we let this profile count for ten percent, and let the remaining light come from a modelthat does not produce very strong absorption or emission at this wavelength, we can perfectly re-produce He iiλ4686. For this latter model we introduce MOD80, which has Teff = 30000 K andL = 105.43 L� and therefore R = 19.24 R�. In order reproduce He ii λ4686 and Hα the best weneed a mass loss rate M = 3.3 × 10−6 M� yr−1 (mass loss flux is 8.93−9 M� yr−1 R−2

� ). Thismodel has default abundances X = 0.710, Y = 0.284. In figure 9.14 we show with the blue dashedprofiles the result of 10% MOD81 and 90% MOD80, which appears to fit for all six analysed lines.The red profile is for MOD80 alone. In the He i lines the WN model does hardly contribute any-thing. MOD80 has the same temperature as MOD31, which was the best fit for these lines. Sincemass loss does not influence this line, MOD80 still provides a good fit. The other He ii lines do havea significant feature for MOD81, but when we only count it for 10%, the combined profile is hardlyinfluenced. The composed profile of Hα is also just strong enough to match the observations. Fromthe Hα profile we conclude that the 10% contribution of the hot, high mass loss source needs tocome from a helium rich star. Otherwise Hα would be much too strong. MOD79 is like MOD81but uses X = 0.398, Y = 0.596. The equivalent width of Hα for MOD79 is roughly 1.5 times asstrong as that of MOD81, while the equivalent widths of He iiλ4686 are similar. With MOD79 itis much more difficult to make a combined profile which is in agreement with the observations inboth He iiλ4686 and Hα.

If we allow the group of stars to have three different components there are of course morepossibilities. For instance, figure 9.15 shows the resulting line profiles when we take 10% of MOD77,40% of MOD10 and 50% of MOD31. MOD10 and MOD31 respectively have M = 1 × 10−6 andM = 6 × 10−6 M� yr−1. The average profile of He iiλ4686 form these two models resembles thatof MOD80. The other parameters of these models can be found in appendix A. The resulting lineprofiles for the other lines only differ significantly in the centre of Hα, but this is a region we cannotmeasure properly due to the nebular emission line subtraction. With this example we show thatthe solution is not unique, especially not when we allow more than two different components.


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416 417 418 419 420 421 422 423 424

wavelength (nm)

HeII λ4200

Figure 9.14: The resulting line profiles for a composition of MOD80 (30000 K, L = 105.43 L�, M/R2 = 8.93−9 M�yr−1 R−2

� , X = 0.710, Y = 0.284) counting for 90% and MOD81 (L = 105.43 L�, M/R2 = 1.73−7 M� yr−1 R−2� ,

X = 0.237, Y = 0.757) counting for 10%. MOD80’s parameters resemble those of a late O type supergiant (O9.5 I)while the properties of MOD81 are similar to a WN star (Koesterke et al., 1991). A physical cluster producing theseblended line profiles could for instance consist of 9 late O supergiants and one WN star.


0.6

0.8

1

1.2

1.4

1.6

464 465 466 467 468 469 470 471 472 473

nor

mal

ised

flu

x

HeII λ4686

MOD10MOD31MOD77

0.4*MOD10 0.5*MOD31 0.1*MOD77observed spectrum

0.6

0.8

1

1.2

1.4

1.6

650 652 654 656 658 660 662

H α

0.7

0.8

0.9

1

1.1

1.2

1.3

437.5 438 438.5 439 439.5 440

nor

mal

ised

flu

x

HeI λ4387

0.7

0.8

0.9

1

1.1

1.2

1.3

443 444 445 446 447 448 449 450 451

HeI λ4471

0.8

0.85

0.9

0.95

1

1.05

1.1

1.15

1.2

452 452.5 453 453.5 454 454.5 455 455.5 456

nor

mal

ised

flu

x

wavelength (nm)

HeII λ4541

0.8

0.85

0.9

0.95

1

1.05

1.1

1.15

1.2

416 417 418 419 420 421 422 423 424

wavelength (nm)

HeII λ4200

Figure 9.15: The resulting line profiles for a cluster composed of 3 different components: 10% of MOD77, 40%of MOD10 and 50% of MOD31. The parameter values can be found in appendix A. This example shows that thecomposition of the cluster we derive is not a unique solution; this combination of models fits all profiles as well.

9.3. CONCLUSION 85

9.3 Conclusion

If we ignore the fact that the luminosity, derived from the observed V , of our source is too high fora single object and we treat the observed spectrum as a single star we can conclude the following.The best fitting model for all usable lines except He iiλ4686 is MOD31, which represents a starof 40 M�, Teff = 30000 K, L = 105.6 L�, M = 6 × 10−6 M� yr−1 and vrot sin(i) = 150 km s−1.The effective temperature is an upper boundary. The luminosity of this model is a factor 7 too lowto be in agreement with the observed V magnitude. Furthermore, this model does not reproducethe observed profile for He iiλ4686. In order to achieve the same equivalent width as observed forthis line, either Teff, M or β should be dramatically increased. If we do so, we must also rise therotational velocity to a value of 1200 km s−1 in order to reproduce the broadness of the observedfeature. But then, the other spectral lines like Hα and He iλ4387 do not fit any more. We couldnot find a set of stellar parameters that reproduces all measured spectral lines in a single model.

If we allow the spectrum to be the result of the combined light of different types of stars itis possible to reproduce all the line profiles we examined in section 9.2 with one consistent sce-nario. Allowing multiple stars is also in agreement with the high observed visual magnitude of thesource, see section 9.1. The simplest scenario that reproduces the observed profiles, and which isalso in agreement with the observed V magnitude, consist of 10% light star with a temperatureTeff = 40000 K, a luminosity L = 105.43 L�, a mass loss rate of 2.02× 10−5 M� yr−1 and a helium-rich composition: X = 0.237 and Y = 0.757 (for all parameters see table A.2, MOD81). Thismodel resembles a Wolf-Rayet star of the sequence WN. The remaining 90% of the light wouldcome from stars that have an average temperature of 30000 K, an average mass loss flux of around8.93 × 10−9 M� yr−1 R−2

� and solar abundances X = 0.710 and Y = 0.284. This contribution ishere simply represented as one stellar model, MOD80, that counts for 90% in luminosity (the modeluses L = 105.43 L� for simplicity, so that we would need exactly nine of them). The parameters ofMOD80 resemble those of an O9.5 supergiant (Martins et al., 2005). We neglect the fact that therelative contributions of the light from stars of two different temperatures are not exactly equal atdifferent wavelengths. The wavelength range we analyse and the temperature range we modelledare both small enough to permit this.

It is very probable that this is not a unique solution; there might be various combinations oftwo or more different stars that, in a specific ratio, provide fitting profiles for all examined lines.Even the simplest scenario that we sketched may correspond to many different physical situations.The luminosities are now chosen such that if we take one WN star (MOD81) and 9 O9.5I (MOD80)stars, the observed V is reproduced. But it is very likely that the 90% component consists of manydifferent kinds of stars. Nevertheless, the contribution of a hot, helium rich star with a high massloss appears to be needed to reproduce He iiλ4686, while keeping the Hα wings modest. For theobjects that produce the remaining 90% of the light we can only derive average values of parameters(Teff = 30000 K; M/R2 = 8.9× 10−9 M� yr−1 R−2

� ). Nonetheless, one WN star and 9 late O-typesupergiants is not an implausible composition for the most luminous components of the cluster. Wecan not draw quantitative conclusions on the lower-luminosity components (e.g. less massive starsthat are still main sequence stars), because we do not see them. But, according to initial massfunctions, it is very likely that these are also there.

10Circumstellar nebula

In this chapter we will analyse the nebular spectrum. This was unavoidably observed simultane-ously, and it is at the same time an obstacle as a bonus. We have already seen that when the strongnebular lines are subtracted from the object spectrum the residual due to the difference in radialvelocity makes many lines difficult to analyse. On the other hand, since our star, or star cluster, isthe most luminous object in this environment, it is most probably also the ionising source of thisregion. Therefore, it might be possible to derive properties of this source from the nebular spectrum.

In section 10.1 we will give a phenomenological description of the nebular spectrum. In sec-tion 10.2 we will apply a quantitative analysis on this nebular spectrum to derive properties suchas electron density and electron temperature of the region. Finally, in chapter 11 will create aconsistent picture of the central source and its surrounding which is in agreement with as manyderived properties as possible.

10.1 Phenomenological description of nebular spectrum

In figures 8.7, 8.8 and 8.9 we plot the observed normalised raw spectrum on a logarithmic scalein order to show the nebular emission lines. Please keep in mind that these spectra also includetelluric lines and the solar spectrum which enters the slit via the moonlight. The lines that havebeen identified with certainty as telluric emission lines are labelled with dashed cyan lines. Skysubtraction is not possible, because there is not a region in the slit without nebular emission. Thiscan be seen from the spatial profiles of Hδ and [O iii] in figure 7.4. On the other hand, we canwork around it as long as we can tell telluric lines from nebular lines. The solar spectrum is muchweaker than the nebular lines. We have to keep in mind that when we for example look at Hα, wesee an ensemble of the following: a sharp emission profile centred at 655.96 nm (145 km s−1) fromthe surrounding nebula and possibly nebulae in the line of sight, a broad emission profile at 655.85nm (195 km s−1) from the central source and a broad absorption profile at approximately the restwavelenth (656.279 nm) from the solar spectrum via the moonlight. The relative radial velocitiesof the different features are small enough for the total line profile to be blended.

We have identified the hydrogen lines of the Balmer series (n = 3, 4, 5 ... → n = 2) from Hα(n = 3 → n = 2) in the VIS at 656.279 nm to H15 (n = 15 → n = 2) at 371.198 nm in theUVB spectrum. Every subsequent line in the series is weaker than the previous. In the VIS, alsothe hydrogen lines of the Paschen series (n = 4, 5, 6 ... → n = 3) can be distinguished. From thestrongest in our range, Paschen-7 (n = 7→ n = 3) at 1000.498 nm to Paschen-30 (n = 30→ n = 3)at 828.642 nm which is the last and weakest we can see. There might be a few more of this seriespresent in the spectrum, but is is not visible because of other nebular and telluric emission lines inthe range 835-830 nm.

10.2. PROPERTIES OF THE SURROUNDING REGION 87

We see a lot of He i lines in emission, but no He ii lines. Like the hydrogen lines, the He iemission lines contaminate the photosheric line profiles when the sky spectrum is subtracted tocreate the object spectrum.

Furthermore we also observe ’forbidden’ emission lines in the nebular spectrum. Those areusually very strong in spectra of gaseous nebula (Osterbrock & Ferland, 2006). We see [N ii], [O ii],[O iii], [S ii], [S iii], [Ar iii], [Ar iv], [Arv] and [Ne iii] forbidden lines. Though there are transitionsof [Ne iv] and [Ne v] in our wavelength range, those emission lines are not observed.

10.2 Properties of the surrounding region

Measurements of intensity ratios of pairs or trios of forbidden emission lines can be used to deriveproperties of the nebula such as electron temperature and electron density. The forbidden linesare created by the collisional excitation of a bound electron, which emits its radiation during de-excitation. Therefore, the temperature and density of the free (colliding) electrons influence thestrength of these lines. Of course the abundance of the ion which forms the line determines theintensity, but since this diagnostics make use of intensity ratios of lines that are all created by thesame ion, the abundances cancel out. Since forbidden lines are usually very strong in emission,many of them are in the optical and ultraviolet, and there is no contamination from photosphericlines, this is an easy and useful method to determine properties of gaseous nebula (Osterbrock &Ferland, 2006).

We do not have a clean nebular spectrum but we consider an ensemble of stellar spectra, nebularspectra, the solar spectrum via the moonlight and spectral features from earth’s atmosphere (theraw spectrum, see section 6.2). We make use of the fact that these emission lines are very strongcompared to those other contributions. Furthermore, forbidden lines can only be formed in a lowdensity region. A stellar photosphere is so dense that the levels will immediately be depopulatedby collisions rather than de-excite radiatively via the forbidden transition. Since gaseous nebulaeusually have a low density, these lines form here.

We have no absolute flux calibrated spectrum. This means we can not measure absolute lineintensities. However, for our analysis we only need line intensity ratios. For line pairs that arevery close in wavelength, like [O ii] and [S ii], the instrument response and continuum level may beassumed constant and we may use the equivalent width ratio (see section 10.2.1). If the lines arenot close in wavelength we do a relative calibration of the flux. This we have done in the case ofthe line trio of [O iii] (see section 10.2.2).

10.2.1 Electron density

Pairs of forbidden lines that are used to determine the electron density in gaseous nebula come fromenergy levels in a single ion with similar energies but different radiative-transition probabilities.The most suitable line pairs for our purpose are [O ii]λ3729/λ3726 and [S ii]λ6716/λ6731, with theenergy level diagrams shown in figure 10.1, left panel.

Let us consider the level populations in the low density limit (ne → 0) for [O ii]. Every collisionalexcitation is followed by a photon. The relative excitation rates of the levels 2D5/2 and 2D3/2 isproportional to their statistical weights.

88 CHAPTER 10. CIRCUMSTELLAR NEBULA

Figure 10.1: Left panel: energy-level diagram of the lowest levels of the ions O ii and S ii that result in the forbiddentransitions used for measuring electron density. Right panel: line intensity ratio as a function of electron density forthe line pairs [O ii]λ3729/λ3726 (solid line) and [S ii]λ6716/λ6731 (dashed line) at T = 10000 K. At other temperaturesthe plotted curves are very nearly correct by taking the horizontal scale to be ne(104/T )1/2. Figures from Osterbrock& Ferland (2006)

Therefore ratio of the strength of the lines will be jλ3729/jλ3726 = 3/2. In the high density limit(ne →∞), collisional de-excitation will dominate and the population ratio will follow a Boltzmanndistribution. The relative population of the levels will be the ratio of their statistical weights.The line strength ratio will be this population ratio times the ratio of transition probabilities.(Osterbrock & Ferland, 2006)

[O ii]jλ3729

jλ3726=

n(2D5/2)n(2D3/2

Aλ3729

Aλ3726=

32

3.6× 10−5

1.6× 10−4= 0.34

�� 10.1

In the right panel of figure 10.1 the variation of the intensity ratio as a function of electrondensity is shown for [O ii]λ3729/λ3726 and [S ii]λ6716/λ6731. We see that the behaviour of [S ii] issimilar, with slightly different ratio limits and critical densities. The same treatment can be carriedout for this ion. In the optical region, the other ions producing pairs of lines suitable for densitymeasurement are [N i], [Cl iii], [Ar iv] and [K v]. Of these we only distinguish [Ar iv], but those areweaker than [O ii] and [S ii], and one of them is blended with a telluric line. Therefore we will focuson [O ii] and [S ii].

Table 10.1 gives the equivalent widths of [O ii]λ3726, [O ii]λ3729, and [S ii]λ6716 and [S ii]λ6731.Because the continua are considered constant over the short wavelength ranges that these line pairsspan, the ratio of equivalent widths is equal to the line intensity ratio. Taking into account theerrors on the measured ratios, the electron density based on the [O ii] pair is between 1 and 20 cm−3

and based on [S ii] between 1 and 80 cm−3. This range is large because both graphs are relativelyflat in this region: we are still in the low density regime. For our model nebulae (see chapter 11)we will adopt a canonical value of ne = 10 cm−3.

10.2.2 Electron temperature

There are a few ions that are suitable for electron temperature measurement. They have energystructures such that there are two different upper levels with considerable different excitation ener-gies. However, the wavelengths of the photons that are emitted during radiative de-excitation should


Table 10.1: The measured equivalent widths of the [O ii] and [S ii] line pairs used for electron density measurement

Ion Line EW (nm) Intensity ratio

[O ii]λ3726 1.66± 0.07 jλ3729

jλ3726= 1.52± 0.08

λ3729 2.53± 0.09

[S ii]λ6716 1.60± 0.02 jλ6716

jλ6731= 1.40± 0.03

λ6731 1.14± 0.02

Figure 10.2: Left panel: energy-level diagram for the lowest terms of the [O iii] and [N ii] ions. Right panel: lineintensity ratio as a function of temperature for the line pairs of four temperature sensitive forbidden line ratios. O[i](solid line) nearly coincides with N[ii] (dashed line) because they have similar excitation potentials. The curves areshown in the low density limit ne = 1 cm−3. Figures from Osterbrock & Ferland (2006).

90 CHAPTER 10. CIRCUMSTELLAR NEBULA

be close enough in wavelength to be both in an observable wavelength interval. The best examplesare [O iii] and [N ii], of which we show the energy-level diagram in the left panel of figure 10.2.This diagram shows that the [O iii]λ4363 line is formed in the transition of an electron from the1S0 level down to the 1D2 level, while both [O iii]λ5007 and λ4959 result from the transition of theintermediate 1D2 down to the degenerate ground level. Every excitation to 1D2 results in either aλ5007 or a λ4959 photon. The ratio, given by the transition probabilities, is close to 3 to 1. Everyexcitation to 1S0 is followed by λ4363 or λ2321, again in a fixed ratio given by the probabilities.The transition down to 1D2 is again followed by the emission of a λ5007 or a λ4959 photon, butthis effect is negligible in comparison with the emission directly from 1D2. Equation 10.2 showsthe temperature dependence of the line intensity ratio with filled in numerical values for collisionstrengths, transition probabilities, frequencies and energies.

[O iii]jλ4959 + jλ5007

jλ4363=

7.90 exp(3.29× 104/T

)1 + 4.5× 10−4ne/T 1/2

�� 10.2

This results from an analytical first order approximation in exp(∆E × 104/T

), with ∆E the

energy difference between the 1D2 and 1S0 levels, and ne. A similar treatment of the line trios of[N ii], [Ne iii] and [S iii] gives the following equations:

[N ii]jλ6848 + jλ6583

jλ5755=

8.23 exp(2.50× 104/T

)1 + 4.4× 10−3ne/T 1/2

,�� 10.3

[Ne iii]jλ3869 + jλ3968

jλ3343=

13.7 exp(4.30× 104/T

)1 + 3.8× 10−5ne/T 1/2

,�� 10.4

[S iii]jλ9532 + jλ9069

jλ6312=

5.44 exp(2.28× 104/T

)1 + 3.5× 10−4ne/T 1/2

.�� 10.5

The temperature dependencies of the different line ratios are shown in the right panel of fig-ure 10.2.

For all line trios, the line in the denominator is the weakest. For [O i], [N ii] and [Ne iii], thisshortest-wavelength line is practically absent or very weak, introducing a very large error in derivedtemperature. It appears that [S iii]λ9532 is blended with another, possibly telluric, line. Thereforewe will only use [O iii].

[Oiii] λ4959 and [Oiii] λ5007 are relatively close in wavelength, but [Oiii] λ4363 is not. To com-pare the line intensities, we have to carry out a relative flux calibration with help of a photometricstandard star.

We compare the flux in ADU in the non-flat-fielded spectrum with the theoretical flux of thestandard star at the wavelengths 436.3 and 497.0 nm (in between the two strongest lines). Weobtain a ‘sensitivity ratio’ η(497.0nm)/η(436.3nm) = 0.73 ± 0.02 for non-flat-fielded spectra. Thissensitivity is mainly determined by the location on the order. We have to use non-flat-fieldedspectra because the flux calibration observation was done with a broad slit, instead of matching slitwidth. We integrate the observed emission lines on a non-flat-fielded, non-sky corrected image (seetable 10.2. The found ratio of line strengths is multiplied by the sensitivity ratio in order to obtainthe real intensity ratio:

jλ4959 + jλ5007

jλ4363= (204± 18)× (0.73± 0.02) = 149± 14.

�� 10.6


Table 10.2: The measured integrated values of the [O iii] line trio on the non-flat fielded, non-sky corrected spectrum.The error estimate is a combination of the error in measured itegral (dominant for the weak line) and the photonnoise (dominant for the strong lines).

Ion Line Integral (ADU) Ratio

[O iii]λ4363 260± 23λ4959 14000± 180 204± 18λ5007 39130± 250

Following equation 10.2, this ratio corresponds to Te = 11200+390−330 K.

11A consistent picture

Judging from the spatial profiles at the nebular emission wavelengths (see figure 7.4) we see thatthe region with strong emission at the location of our source, is also centred at our source, thoughthere are other strong emission peaks in the environment. This surrounding region has a diameterof 40−50 pc, though it is difficult to measure because the borders are not clear. The central locationof our source in this ionised region is an indication that our cluster is the most important ionisingsource of this region. Under this assumption, it must be possible to bring the conclusions on thecentral source and the surrounding nebula in agreement with each other.

11.1 Procedure

We have used CLOUDY version 08.00 (Ferland et al., 1998) to model the surrounding nebula.CLOUDY is designed to simulate clouds of various densities. The program predicts the thermal,ionisation and chemical structure and computes the emitted spectrum. A number of spectral energydistributions can be put in as central ionising source. Though there are many options in CLOUDY,we will keep our model simple.

The geometry is spherical: the inner radius of the cloud is 0.1 pc. In order to obtain the spec-trum from a pencil beam directed toward the centre of the cloud, we add the command aperturebeam. This allows us to compare the emission spectra to what we observe: the nebular spectrum atthe location of the central source. The initial total (ionic, atomic and molecular) hydrogen numberdensity nH of the cloud is set to 10 cm−3 for the standard grid, which is in agreement with thederived electron density (section 10.2.1). In an ionised region, the hydrogen density will be roughlyequal to the electron density. We will check whether this choice of nH yields a correct result bytrying other values that are within the error bar of the number density measurement.

From the spectral energy distributions produced by FASTWIND we constructed a grid thatcan be called as ionising source in CLOUDY. This grid consists of MOD80 (the 30000 K O9.5Icomponent) and MOD81 (the 40000 K WN component) with seven different metallicities (Z = 0.1−0.7 Z�). The ionising source consists of 9 times MOD80 plus once MOD81: our most simple clustercomposition derived in section 9.2.6. The upper panel of figure 11.1 shows the flux distribution atthe inner radius of the cloud for the cluster and for the two different components (Z = 0.2 Z�).

We make seven cloud models that only differ in metallicity (that is assumed to be identical forthe members of the cluster and the surrounding medium). We will use the stored abundance set‘HII region’, with abundances that are measured from the Orion Nebula (see table B.1). Withthe metals command, we scale all elements heavier than helium. The derived metallicities are thuswith respect to the Orion nebula, which is close to solar. The cloud models do not contain dust.

11.2. RESULTS 93

We add commands to print the continuum and line emission of the clouds, and the temperatureprofile as a function of radius.

Since the cloud is ionised by the central source, its properties will be strongly dependent on q0,the flux of photons above the hydrogen ionisation energy (λ ≤ 91.2 nm or ν ≥ 3.29 × 1015 Hz). Itis defined as

q0(r) =∫ ∞

ν0

Fν(r)hν

dν.�� 11.1

11.2 Results

CLOUDY computes the temperature of different particles as a function of radius. We force it tostop when the electron temperature falls below 4000 K, i.e. when hydrogen strongly recombines.Figure 11.2 shows the electron temperature as a function of r, the distance to the centre of the cloud,for different input values for Z and nH. From this, CLOUDY computes radius and volume-averagesof the temperature. However, to be consistent in our way of deriving the electron temperature, wewill examine the output emission line spectrum, to which we apply the same measurement as wedid on the observations. A bias resulting from this way of measuring electron temperature will beincluded in both the model spectrum and the observed spectrum, and is possibly cancelled out. Asa check, we also measure the line ratios used for the electron density. They should reproduce ourinserted nH. Table 11.1 lists our results.

The line strength ratios found for [O ii] and [S ii] all are in prefect agreement with the inputvalue of nH and with our observations. However, there appears to be a slight trend in the [S ii] ratio:the ratios for higher metallicity are higher. A high ratio corresponds to a low electron density (seefigure 10.1, right panel). So the lower the metallicity, the higher the electron density. Because thistrend is not present in the [O ii] ratio, while this is more density dependent in this limit, we regardit as a fluctuation or a bias in the density determination based on the [S ii] ratio.

The upper panel of figure 11.2 shows the electron temperature gradient in the cloud, for differentinput values for Z. Te(r) is relatively flat up to 80% of its radius. Then the temperature dropsroughly 1000 K, and after a rise, of which the relative strength appears to be dependent on Z, thetemperature drops very quickly down to 4000 K. This point is defined as the edge of the region. Te

in table 11.1 is derived from the emission spectrum from a pencil beam through the centre of thecloud, so this is an average over r. We see in table 11.1, as well as in the upper panel of figure 11.2,that the electron temperature of the nebula appears to be very sensitive to the metallicity of thenebula. The lower panel of figure 11.1 shows that the spectral energy distribution and especiallythe total number of ionising photons is not significantly sensitive to Z. If we keep the central sourceconstant while varying only the nebular metal content, we see a similar trend in Te. The nebula cancool more efficiently with a higher metal abundance. The forbidden line radiation resulting fromthe collisionally excited low-lying energy levels of ions such as O ii, O iii, and N ii is by far the mostimportant source of radiative cooling in diffuse nebulae (Osterbrock & Ferland, 2006).

Figure 11.3 shows the electron temperature that is derived from the [O iii] lines as a func-tion of the metallicity. A logarithmic function appears to be a good fit of this trend: Te =a log (Z/Zorion) + b, with a = 3565.66± 20.23 K and b = 5508.67± 25.42 K.

94 CHAPTER 11. A CONSISTENT PICTURE

1e-20

1e-18

1e-16

1e-14

1e-12

1e-10

1e+14 1e+15 1e+16

Fν(

0.1

pc)

(er

g cm

-2 s

-1 H

z-1)

ν (Hz)

Input spectra of the cluster and its two components for Z = 0.2 Zsun

cluster spectrum9 x SG spectrum

1 x WN spectrum

1e-20

1e-18

1e-16

1e-14

1e-12

1e-10

1e+15 1e+16

Fν(

0.1

pc)

(er

g cm

-2 s

-1 H

z-1)

ν (Hz)

Input spectra of the cluster for different metallicities

Z=0.1ZsunZ=0.2ZsunZ=0.3ZsunZ=0.4ZsunZ=0.5ZsunZ=0.6ZsunZ=0.7Zsun

Figure 11.1: The flux as a function of frequency at 0.1 pc, the inner radius of the cloud. Upper panel: the fluxes ofthe separate components of the cluster. The orange line corresponds to the 9 O9.5 supergiant component (MOD80:Teff = 30000 K, log g = 3.47 and L = 105.43 L� multiplied by 9). The green line corresponds to the WN component(MOD81: Teff = 40000 K, log g = 3.97 and L = 105.43 L�). The blue line is the combined ionising spectrumilluminating the inner radius of the cloud. The metallicity of the cluster components is Z = 0.2 Z�. We see that bothcomponents are important for the ionisation of the cloud. Lower panel: the combined FASTWIND spectral energydistributions of the whole cluster with varying metallicity (Z = 0.1− 0.7 Z�). We see that the high frequency end ofthe spectrum is sensitive to the metallicity, however, the differences in total number of ionising photons are all within2% for the different metallicities.

11.2. RESULTS 95

4000

6000

8000

10000

12000

14000

16000

0 5 10 15 20 25

Te

(K)

r (pc)

Te(r) for modelled clouds of different metallicity

Z = 0.1 ZorionZ = 0.2 ZorionZ = 0.3 ZorionZ = 0.4 ZorionZ = 0.5 ZorionZ = 0.6 ZorionZ = 0.7 Zorion

4000

6000

8000

10000

12000

14000

16000

0 20 40 60 80 100 120

Te

(K)

r (pc)

Te(r) for modelled clouds of different density (Z=0.2Zorion)

nH = 1 cm-3

nH = 10 cm-3

nH = 20 cm-3

Figure 11.2: The predicted electron temperature Te as a function of distance from the centre of the cloud r. Upperpanel: temperature gradient for models with different metallicities. These models have nH = 10 cm−3. The lowestmetallicity cloud starts the hottest, but the high metallicity models have the temperature increased relatively more forlarger r. The final radii are lower for higher metallicity models. Lower panel: temperature gradient for three modelswith a metallicity of 0.2 Zorion, but now with different densities. The input nH is respectively 1, 10 and 20 cm−3

(values that are all within the error-bars of our measurement, see section 10.2.1). We see that the starting and averagetemperature are not significantly different. However, the low density model cloud is much larger. The final radius fornH = 10 cm−1 is in best agreement with the physical size of the nebula we measure from figure 7.4.


Table 11.1: The results of the grid of cloud models. Columns 1 and 2 list nH and Z, the adopted total hydrogendensity and metallicity. Column 3 gives the final radius of the cloud. This is where the temperature falls below4000 K. Columns 4 to 6 give the forbidden line intensity ratios that we computed from the emission lines (in a pencilbeam though the centre of the cloud) predicted by these models. Column 7 gives the electron temperature that isderived from the [O iii] ratio using equation 10.2, following the same method as we used on our observed nebularspectrum. nH = 10 cm−3 is the standard density for our grid, which probes metallicities of Z/Zorion = 0.1 to 0.7.nH = 1 cm−3, and nH = 20 cm−3 are values that are also in agreement with the measured ne. Calculations with thesevalues are performed to see if the obtained Te is significantly different. This is apparently not the case.

nH Z Rout [Oii] ratio [Sii] ratio [O iii] ratio Te

(cm−3) ( Zorion) (pc) (K)

10 0.10 24.74 1.48 1.41 84.47 1375810 0.20 23.67 1.48 1.42 145.03 1122210 0.30 23.04 1.48 1.42 224.63 976410 0.40 22.53 1.48 1.43 331.49 875410 0.50 22.14 1.48 1.43 479.15 797210 0.60 21.82 1.48 1.43 679.94 734910 0.70 21.54 1.48 1.44 964.87 6816

1 0.10 115.21 1.50 1.42 84.42 137611 0.20 110.64 1.50 1.43 142.47 112901 0.30 107.82 1.50 1.44 219.14 9837

20 0.10 15.61 1.46 1.40 84.62 1374820 0.20 14.92 1.46 1.40 145.12 1121920 0.30 14.52 1.46 1.41 225.15 9758

11.3. CONCLUSIONS 97

6000

8000

10000

12000

14000

0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8

Te

[OII

I] (

K)

Z/Zorion

Derived Te in model nebula with different Z

Model nebulae ionised by clusterTe = -3565.66 * log(Z/Zorion) + 5508.67

derived Te

Figure 11.3: The derived electron temperature for the cloud model from the [O iii] line strength ratios as a functionof metallicity. The results can be fitted with a logarithmic function Te = a log (Z/Zorion)+ b, with a = 3565.66±20.23K and b = 5508.67± 25.42 K. It appears that the temperature is proportional to the logarithm of the metallicity. Forour measured electron temperature Te = 11200+390

−330 K (green point) we find a metallicity of 0.20± 0.02 Zorion for thisregion.

The electron temperature Te = 11200+390−330 K that is derived from our observations (section 10.2.2)

would, according to this modelling exercise, be produced by our cluster if the metallicity of the sur-rounding nebula has a value of Z = 0.20±0.02 Zorion. This small formal error bar is the result fromour measuring error in Te. The systematic errors that may result from our method are discussed insection 11.3.

The derived Te is not significantly affected by nH, for values that are within our measurederror estimates. We see in the lower panel of figure 11.2 and in table 11.1, however, that a lowdensity cloud stretches out much further than a higher density cloud. The physical radius of thenebula we measure from figure 7.4 is 20 to 25 pc. This would mean that nH = 10 cm−3 is the bestchoice. But for our derivation of the temperature, and subsequently the metallicity, we can takeany nH ≤ 20 cm−3.

11.3 Conclusions

We assume that the part of the nebula we observe is ionised by the observed source NGC 55 CI 31.From the visual magnitude and the line profiles of He i, He ii and Hα in our observed spectrumwe derive that NGC 55 CI 31 is a cluster with several stars of different temperature, mass loss, andchemical composition contributing to the integrated light. From the forbidden emission lines of thenebula we derive an electron density and an electron temperature. When we let the spectrum of thiscluster ionise a model cloud which has the measured electron density, we reproduce the measuredelectron temperature only if the region has a metallicity of Z = 0.20 ± 0.02 Zorion. Earlier metal-licity estimates of this region are based on the oxygen abundance. Our oxygen abundance (based


on table B.1) would be 12 + log O/H = 7.90± 0.04. Using Asplund et al. (2004) solar abundances,this corresponds to Z = 0.17± 0.02 Z�.

This metallicity is lower than that has been measured previously for H ii regions in the disk ofNGC 55 (see section 2.2), so this is quite a strong claim. The formal error bar in our estimate isbased on the errors in the line integration of the [O iii] emission lines. However, we make severalapproximations and assumptions that could affect our error estimate. Below we make a few remarksconcerning our method and conclusion.

� All FASTWIND models used for determining a fitting composition for the line profiles (seesection 9.2.6) use a metallicity of Z = 0.5 Z�. In this exercise we alter the metallicity of theionising source. We verified that the line profiles we examine, and thus the derived componentsof the cluster will not be heavily influenced by the metallicity. Figure 11.4 shows how thecombined line profiles would look if we change the metallicity of the preferred models fromZ = 0.5 Z� to Z = 0.2 Z�. It is clear that these combined line profiles are still in agreementwith our (noisy) observations.

� One could argue that the inner radius of our ionised region is too small (0.1 pc). Not only mustour cluster fit into this region, the strong stellar wind could have swept up the surroundingmaterial forcing a larger inner radius. The UVB slit width (0.8′′) allows our cluster to havea diameter of 7.8 pc. Figure 7.2 shows that it can not be much larger than this. We checkedwhether changing the inner radius would affect other properties strongly. Changing the innerradius to 1 pc does not result in significantly different line ratios. When we increase the innerradius to 5 pc, we derive a 1% lower temperature. The electron density that is derived fromthe lines, as well as the outer radius remain unaffected.

� It could also be possible that the region is ionised by additional sources, and not onlyNGC 55 CI 31. Figure 7.2 as well as 7.4 show that this is a crowded region, with nebularemission along the whole slit. It could be that the region that is investigated here is part ofa larger H ii region ionised by multiple sources and clusters.

� Again it needs to be pointed out that the cluster composition derived in chapter 9 is probablynot a unique solution. When other compositions can be found that have significantly differenttotal numbers of ionising photons, probably a different value for Z is found. The picturesketched in this chapter is internally consistent, but not a unique solution to this problem.

11.3. CONCLUSIONS 99

0.6

0.8

1

1.2

1.4

1.6

464 465 466 467 468 469 470 471 472 473

nor

mal

ised

flu

x

HeII λ4686

MOD80 Z=0.2ZsunMOD81 Z=0.2Zsun

0.9*MOD80+0.1*MOD81observed spectrum

0.6

0.8

1

1.2

1.4

1.6

650 652 654 656 658 660 662

H α

0.7

0.8

0.9

1

1.1

1.2

1.3

437.5 438 438.5 439 439.5 440

nor

mal

ised

flu

x

HeI λ4387

0.7

0.8

0.9

1

1.1

1.2

1.3

443 444 445 446 447 448 449 450 451

HeI λ4471

0.8

0.85

0.9

0.95

1

1.05

1.1

1.15

1.2

452 452.5 453 453.5 454 454.5 455 455.5 456

nor

mal

ised

flu

x

wavelength (nm)

HeII λ4541

0.8

0.85

0.9

0.95

1

1.05

1.1

1.15

1.2

416 417 418 419 420 421 422 423 424

wavelength (nm)

HeII λ4200

Figure 11.4: This is the same figure as figure 9.14, only all models have Z = 0.2 Z� instead of Z = 0.5 Z�, whichwas the default value for all other models. All six line profiles are reproduced with the same composition of models,even if they include a much lower metallicity.

Part V

Discussion and Conclusions

12Discussion

12.1 A comparison with former studies

12.1.1 Central source

In section 2.3 we pointed out that our source was selected as a candidate most massive star inthis region based on the observations and classification of Castro et al. (2008). Now that we haveobserved and analysed the source ourselves, a few comments on their preliminary spectrum andclassification can be made. The spectrum of NGC55 C1 31 obtained by Castro et al. is shown infigure 2.3 and the properties of this source according to Castro et al. can be found in table 2.3.

1. The star is classified as an early O-type supergiant, implying it to be hot. But we do not seethe He iiλ4541 in their spectrum. This line usually is in absorption in hot stellar atmospheres.On the other hand, we do see He iλ4471, which increases in strength towards later subclassesof O stars. Even without the thorough analysis of our newly obtained X-Shooter spectrum,one can deduce from this low-resolution spectrum that the object can not be an early O-typesupergiant.

2. The star is possibly classified as an early O-type supergiant because of the He iiλ4686 emission,which is an indication for a strong wind. However, though the low-resolution spectrum displaysthe same unusual shape of this line (broad but not very strong) as we have observed, thisfeature is not discussed by Castro et al..

3. Castro et al. comment that this source suffers from nebular contamination. However, theirlow-resolution spectrum shows pronounced absorption features at the Balmer lines (see fig-ure 2.3). The Balmer lines of the nebula are in emission. It is not described in the article howthey correct for the nebular emission. It looks as if the nebular emission lines have just beensubtracted from the photospheric spectrum, because these hydrogen lines could never be thisstrong in absorption for an O-type star with a high mass loss.

4. The V magnitude we use originates from Castro et al., but apparently they did not computea luminosity based on the derived temperature and the distance. Otherwise they would havenoticed that the source would be strongly over-luminous if it was just a single object.

Mainly based on the first point, we have revisited the classification of NGC 33 C1 31 proposed byCastro et al. (2008). Additionally, Castro et al. did not take into account the possibility of hav-ing several objects in the slit. Yet, the overall properties the spectrum (no He iiλ4541 but broadHe iiλ4686 in emission) can only be reproduced using a combination of spectra.Apart from the Araucaria Project, there are no other published observations of this source. We cantherefore not compare our conclusions on the central source with former research.

102 CHAPTER 12. DISCUSSION

12.1.2 Surrounding nebula

We also have derived properties on the H ii region that is surrounding our source. We conclude thatthe electron density ne ≤ 20 cm−3 and the electron temperature Te = 11200+390

−330 K. To bring thesevalues in agreement with our derived cluster properties, we must assume Z = 0.20± 0.02 Zorion forthis region. This yields an oxygen abundance 12 + log O/H = 7.90 ± 0.04 (see table B.1). Theseparameters can be compared with former observations.

In section 2.2 we discuss several papers concerning emission line spectroscopy on H ii regions inthe disk of NGC 55. The exact region that is probed in each of these papers is different; thereforewe allow for small discrepancies between our derived values and these measurements.

The electron densities for these H ii regions are estimated to be < 100 cm−3 (Webster & Smith,1983; Stasinska et al., 1986) and 300 ± 90 (Tullmann et al., 2003). Electron temperatures are de-rived to be 10180 K (Webster & Smith, 1983), 9200± 800 (Stasinska et al., 1986) and 11400± 300(Tullmann et al., 2003) based on the [Oiii] line ratio, using the same method we have used. Thederived oxygen abundances are listed in table 2.2.

Our electron density is similar to those derived by Webster & Smith and Stasinska et al.. Becausethey use the [Sii]λ6716/λ6730 ratio, they also only derive an upper limit. The [Oii]λ3729/λ3926ratio can not be measured in their data because the resolution of their spectra is too low to distin-guish these two lines. The electron density from Tullmann et al. is significantly higher; they use adifferent method in which ne depends on the derived Te. It appears that our electron temperatureis an acceptable value for an H ii region in the disk of NGC 55.

Our electron temperature is comparable to the values determined in earlier research, and agreesbest with Tullmann et al.. In section 2.2 we have summarised the oxygen abundance measurements.These estimates are very sensitive to the measured temperature, especially for those derived by theR23 method. In table 2.2 we see that our value for the oxygen abundance is just within the error barof the lowest estimate: 12 + log O/H = 8.05± 0.10 by Tullmann et al.. They explain the differencesin oxygen abundance by the fact that different regions in the disk are probed. They state thatthe so called ‘central’ H ii region actually consists of three blobs of emission within a common halo(visible on their Hα image). Their slit cuts through the south-eastern one which has apparently thehighest temperature and hence the lowest oxygen abundance.

If the electron temperature and therefore the oxygen abundance can indeed vary so stronglybetween different H ii regions that are relatively close together (up to a kpc at most), the regionaround our cluster could have an oxygen abundance as low as what we derive. The remarks insection 11.3 suggest that the real error on our derived oxygen abundance is probably higher thanthe formal error bar suggests. These are just based on the error in measuring Te.

12.2 Future investigation

This project served as a pilot project for follow-up observation of similar objects with X-Shooter. Astart of these observations has been made in September 2010. The list of observed targets consistsof five additional massive star candidates in NGC 55 and five in IC1613. The latter is an irregulardwarf galaxy that probably has an even lower metallicity. It is also closer than NGC 55, at 0.73±0.2Mpc (Karachentsev et al., 2004).

12.2. FUTURE INVESTIGATION 103

Apart from building expertise in the data reduction and analysis, we also learned a few usefulthings from this pilot study. These new insights were taken into account in the selection of sourcesand the planning of the follow-up observations.

� Dark time is necessary for this kind of observations. Diffuse moonlight resulted in a high skylevel, especially in the UVB.

� Good weather conditions (seeing) also turned out to be very important. The first night of ourobservations (August) had the lowest seeing and was darkest. Despite that it only had halfthe exposure time of the September nights, the signal-to-noise ratio of the August spectrumis comparable to that of the other nights (see table 6.4).

� In the selection of the sources, the massive star candidates with the least nebular contamina-tion should get priority. The nebular spectrum might be used as well, but it is better to havea clear photospheric spectrum if we want to constrain the properties of the star.

It would be very interesting to see what the properties of the other sources in NGC 55 are. Maybe itis possible to draw conclusions on the metallicity of these stars and their environments. This wouldfor example be interesting if the mass loss rates can also be well constrained. It would furthermorebe very valuable if quantitative conclusions on the masses of these objects can be drawn: then wewould know whether these candidates are really the most massive stars in these regions.The observations of the objects in IC1613 also have high potential. The quality of these spectra will,in similar observing conditions, probably be better because the sources are much closer. Maybe evenmore exotic properties will appear in this low-metallicity environment. Expanding our knowledgeof stars and stellar evolution to the lowest metallicities can provide good tests of various Z-relatedhypotheses.

13Conclusions

Based on the classification as early O-supergiant (Castro et al., 2008) we selected NGC 55 C1 31 asa massive star candidate in NGC 55. We obtained a wide band intermediate-resolution spectrum ofthis star using X-Shooter.

We acquired a photospheric spectrum of the object as well as an emission line spectrum of theH ii region that surrounds our source. Based on spatial profiles along the slit at emission line wave-lengths, we conclude that our object is likely the central ionising source of this region.

We constructed a grid of model atmospheres and predicted line profiles for several H i andH ii lines and Hα. We compared our object spectrum with these model line profiles to constrainparameters of the source.

If we restrict ourselves to a single object, the best fitting model represents a star with amass of 40 M�, Teff = 30000 K, R = 23.40 R� L = 105.6 L�, M = 6 × 10−6 M� yr−1 andvrot sin(i) = 150 km s−1. This model star is not luminous enough to reproduce the high visualbrightness of the observed object. Scaling the stellar radius, preserving at the same time the sur-face gravity and mass loss flux, such that the correct luminosity is achieved, resulted in a spectrumin which He iλ4387 and Hα could not be reproduced simultaneously in a model with solar composi-tion. Apart from this issue, no single model could reproduce a matching line profile for He iiλ4686with a rotational velocity that also matched the other line profiles.

It turned out that we could not satisfy all these conditions (all line profiles and the visual bright-ness) with a single model. Therefore we combined different models. Our simplest solution is thatthe object is a small cluster, of which we only observe the most luminous components. In the spec-trum of this cluster, 10% of the (UV-visual) light is produced by a hot (Teff = 40000 K), helium rich(X = 0.237, Y = 0.757) star with a high mass loss (M = 2.02×10−5 M� yr−1). The parameters ofthis model resemble those of a Wolf-Rayet star of the WN sequence. The remaining 90% of the lightcomes from objects of Teff = 30000 K with solar compositions (X = 0.710, Y = 0.284) and massloss fluxes 8.93 × 10−9 M� yr−1 R−2

� . The simplest realistic physical situation producing such aspectrum is a cluster that consists of one WN star and 9 O9.5 supergiant stars. Additionally, theremight be many main-sequence stars in this cluster, but since they would be far less luminous, wedo not detect them.

The forbidden lines that are produced in the surrounding H ii region reveal that the region hasa low electron density ne ≤ 20 cm−3 and a relatively high electron temperature Te = 11200+390

−330 K.We constructed a model nebula of low density in which we use our derived cluster as ionising source.In order to reproduce our measured electron temperature, the metallicity of this cloud must be low:Z = 0.17 ± 0.02 Z�. This low metallicity is just within error bars of earlier measurements of themetallicity of H ii regions in the disk of NGC 55.

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AFASTWIND model grid

In the tables in this appendix, we list the parameter values of the FASTWIND models that have beenused in the analysis of the line profiles (see section 9.2). Table A.1 lists the models of the originalgrid (MOD01 to MOD27) integrated with the intermediate mass loss value of 6.00× 10−6 M�yr−1

(MOD28 to MOD36). Table A.2 gives the parameter values of all models that were used to see theeffect of changing one or more of the parameters.

110 APPENDIX A. FASTWIND MODEL GRID

Table A.1: Main grid of models varying three parameters: effective temperature Teff = 30000, 32500 and 35000 K,log L/L� = 5.4, 5.6 and 5.8 and M = 1.00 × 10−6, 3.00 × 10−6, 6.00 × 10−6 and 1.00 × 10−5 M� yr−1. ID is thename of the model which we refer to in the analysis, M is the star’s mass, Teff effective surface temperature, L is theluminosity, g is the surface gravity, R is the stellar radius, M is the mass loss rate, v∞ is the terminal wind velocity,β is the rate of acceleration of the outflow, X is the mass fraction of hydrogen and Y is the mass fraction of helium.

ID M Teff L log g R M v∞ β X Y(M�) (K) (L�) (cm s−2) (R�) (M� yr−1) (km s−1)

MOD01 40 30000 105.4 3.50 18.58 1.00× 10−6 2356.69 1.0 0.710 0.284MOD02 40 30000 105.4 3.50 18.58 3.00× 10−6 2356.69 1.0 0.710 0.284MOD28 40 30000 105.4 3.50 18.58 6.00× 10−6 2356.69 1.0 0.710 0.284MOD03 40 30000 105.4 3.50 18.58 1.00× 10−5 2356.69 1.0 0.710 0.284MOD04 40 32500 105.4 3.64 15.84 1.00× 10−6 2553.08 1.0 0.710 0.284MOD05 40 32500 105.4 3.64 15.84 3.00× 10−6 2553.08 1.0 0.710 0.284MOD29 40 32500 105.4 3.64 15.84 6.00× 10−6 2553.08 1.0 0.710 0.284MOD06 40 32500 105.4 3.64 15.84 1.00× 10−5 2553.08 1.0 0.710 0.284MOD07 40 35000 105.4 3.77 13.65 1.00× 10−6 2749.47 1.0 0.710 0.284MOD08 40 35000 105.4 3.77 13.65 3.00× 10−6 2749.47 1.0 0.710 0.284MOD30 40 35000 105.4 3.77 13.65 6.00× 10−6 2749.47 1.0 0.710 0.284MOD09 40 35000 105.4 3.77 13.65 1.00× 10−5 2749.47 1.0 0.710 0.284



111

Table A.2: Parameter values of additional models. The first block is an extension of the range of Teff for the modelwith the best fitting mass loss rate. All are variations on MOD34 with different temperatures. Parameters g, R andv∞ have been changed accordingly. The second block contains models to probe different β and v∞. The third blockcontains models that were used to see the effect of changing the helium abundance. MOD70 is like MOD31 exceptfor the abundances. MOD60 and MOD61 are models to simulate Wolf-Rayet stars, with increased mass, mass lossand luminosity. The fourth block contains the models with luminosities that reproduce V . MOD72 preserves thetemperature, the surface gravity and the mass flux of MOD31. MOD73 also does this, but now at a lower Teff. Thefifth block are Teff = 40000 K models with different mass loss rates, luminosities and abundances. The last block onlycontains the model representing the O9.5I star. ID is the name of the model which we refer to in the analysis, M isthe star’s mass, Teff effective surface temperature, L is the luminosity, g is the surface gravity, R is the stellar radius,M is the mass loss rate, v∞ is the terminal wind velocity, β is the rate of acceleration of the outflow, X is the massfraction of hydrogen and Y is the mass fraction of helium.

ID M Teff L log g R M v∞ β X Y(M�) (K) (L�) (cm s−2) (R�) (M� yr−1) (km s−1)

MOD40 40 25000 105.6 2.99 33.69 6.00× 10−6 1750.33 1.0 0.710 0.284MOD39 40 27500 105.6 3.15 27.84 6.00× 10−6 1925.37 1.0 0.710 0.284MOD38 40 29000 105.6 3.24 25.04 6.00× 10−6 2030.39 1.0 0.710 0.284MOD37 40 31000 105.6 3.36 21.91 6.00× 10−6 2170.41 1.0 0.710 0.284

MOD62 40 30000 105.6 3.30 23.40 6.00× 10−6 2100.40 0.7 0.710 0.284MOD65 40 30000 105.6 3.30 23.40 6.00× 10−6 2100.40 1.2 0.710 0.284MOD67 40 30000 105.6 3.30 23.40 6.00× 10−6 2100.40 2.0 0.710 0.284MOD69 40 30000 105.6 3.30 23.40 6.00× 10−6 2100.40 3.0 0.710 0.284MOD64 40 30000 105.6 3.30 23.40 6.00× 10−6 3000.00 0.7 0.710 0.284MOD63 40 30000 105.6 3.30 23.40 6.00× 10−6 3000.00 1.0 0.710 0.284MOD66 40 30000 105.6 3.30 23.40 6.00× 10−6 3000.00 1.2 0.710 0.284MOD68 40 30000 105.6 3.30 23.40 6.00× 10−6 3000.00 2.0 0.710 0.284

MOD70 40 30000 105.6 2.99 23.40 6.00× 10−6 1750.33 1.0 0.398 0.596MOD60 40 30000 106.0 2.90 37.08 1.50× 10−5 2000.00 1.0 0.398 0.596MOD61 100 30000 106.0 3.30 37.08 1.00× 10−4 2637.98 1.0 0.398 0.596

MOD72 270 30000 106.43 3.30 60.84 4.07× 10−5 3384.18 1.0 0.710 0.284MOD73 440 25000 106.33 3.30 78.08 6.71× 10−5 3813.44 1.0 0.710 0.284

MOD77 40 40000 105.6 3.80 13.60 3.00× 10−5 2800.53 1.0 0.398 0.596MOD79 40 40000 105.43 3.97 10.82 2.02× 10−5 3088.45 1.0 0.398 0.596MOD81 40 40000 105.43 3.97 10.82 2.02× 10−5 3088.45 1.0 0.237 0.757

MOD80 40 30000 105.43 3.47 19.24 3.30× 10−6 2316.34 1.0 0.710 0.284

BCLOUDY abundances

Table B.1: The element abundances in the stored abundance set HII region, compared to the default solar compos-tition in CLOUDY. The HII region abundances are a subjective mean of the Orion Nebula abundances determinedby Baldwin et al. (1991), Rubin et al. (1991) and Osterbrock et al. (1992). Abundances of some rare species weretaken from the ISM mix of Savage & Sembach (1996). Abundances are displayed as number fractions of hydrogen.References for the solar compositions are: Allende Prieto et al. (2001, 2002) for C and O; Holweger (2001) for N, Ne,Mg, Si, and Fe; Grevesse & Sauval (1998) for the remainder of the first thirty elements.

Element ‘HII region’ Solar Element ‘HII region’ Solar

Hydrogen 1.0 1.00 Sulphur 1.0× 10−5 1.84× 10−5

Helium 9.5× 10−2 1.00× 10−1 Chlorine 1.0× 10−7 1.91× 10−7

Lithium 5.4× 10−11 2.04× 10−9 Argon 3.0× 10−6 2.51× 10−6

Beryllium 1.0× 10−20 2.63× 10−11 Potassium 1.1× 10−8 1.32× 10−7

Boron 8.9× 10−11 6.17× 10−10 Calcium 2.0× 10−8 2.29× 10−6

Carbon 3.0× 10−4 2.45× 10−4 Scandium 1.0× 10−20 1.48× 10−9

Nitrogen 7.0× 10−5 8.51× 10−5 Titanium 5.8× 10−10 1.05× 10−7

Oxygen 4.0× 10−4 4.90× 10−4 Vanadium 1.0× 10−10 1.00× 10−8

Fluorine 1.0× 10−20 3.02× 10−8 Chromium 1.0× 10−8 4.68× 10−7

Neon 6.0× 10−5 1.00× 10−4 Manganese 2.3× 10−8 2.88× 10−7

Sodium 3.0× 10−7 2.14× 10−6 Iron 3.0× 10−6 2.82× 10−5

Magnesium 3.0× 10−6 3.47× 10−5 Cobalt 1.0× 10−20 8.32× 10−8

Aluminium 2.0× 10−7 2.95× 10−6 Nickel 1.0× 10−7 1.78× 10−6

Silicon 4.0× 10−6 3.47× 10−5 Copper 1.5× 10−9 1.62× 10−8

Phosphorus 1.6× 10−7 3.20× 10−7 Zinc 2.0× 10−8 3.98× 10−8

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