A spatially resolved study of ionized regions in galaxies at different ...

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4 1 • Introduction Object Galaxy Diameter L(Hα) M(Hii) N (O5V) (pc) (erg s −1 ) (M ⊙ ) Orion MWG 5 1.0 × 10 37 50 0.2 NGC 3603 MWG 100 1.5 × 10 39 3.9 × 10 4 20 W49 MWG 150 1.8 × 10 39 4.5 × 10 4 27 NGC 604 M33 400 4.5 × 10 39 7.0 × 10 5 65 30 Doradus LMC 370 1.5 × 10 40 6.0 × 10 5 230 NGC 5455 M101 750 2.5 × 10 40 · · · 380 Table 1.1: Integrated properties of some Hii regions. M(Hii) is the mass of ionized hydrogen and N (O5V) the equivalent number of O5 V stars needed to match the ionizing luminosity of the region. The abbreviations MWG and LMC stand for the Milky Way Galaxy and Large Magellanic Cloud, respectively. Based on Kennicutt (1984). tantly, they are used to measure abundances of various elements through their emission lines. These emission lines can be used to determine some of the physical properties of the ionized gas and how the metallicity changes within a galaxy and among galaxies (McCall et al., 1985). These studies give clues for models of the chemical evolution of galaxies. Furthermore, these emission lines yield kinematic information about the gas in the galaxy, and therefore, allows us to study the dynamics of galaxies. GEHRs have a wide range of metallicities. In principle, the effective temperature of ionizing stars should be lower in regions of higher metallicity due to the increasing opacity of the stellar material (McGaugh, 1991). However, highly evolved massive O stars, Wolf-Rayet (WR) stars, would increase their surface temperature to very high values because of loss of the outer envelope due to the strength of stellar winds (which increases with metallicity). As the limiting mass for a star to enter the WR phase decreases with increasing metallicity, it is expected to find a higher fraction of WR stars in high metallicity environments. Thus, the finding of WR stars in this type of Hii regions would indicate a relatively high excitation. Since the models predict larger equivalent widths and luminosities of the WR “bump” at higher metallicities, the detection of these WR features combined with a detailed analysis of the emission line spectra can be a powerful tandem for the determination of both metallicity and age in these regions. GEHRs are an excellent environment to study large groups of massive young stars such as star clusters. Massive stars play an important role in the structure and evolution of galaxies due to their ionizing radiation, mechanical energy in the form of winds, and chemical enrichment of the ISM. Most of them are born in stellar clusters ranging from those with one or a few massive stars to massive young clusters with hundreds or thousands of massive stars, which can change the global appearance of their host galaxy, even ejecting a good fraction of the ISM in a superwind (Tenorio-Tagle et al., 2006).
1.1. Overview 5 These massive star clusters can have diverse morphologies (Maíz-Apellániz, 2007). Some consist of a single core, others contain double cores, and others are coreless. Also, some are surrounded by a strong halo that contains more stars than the core and some have few stars outside the core. Those without a core are likely to be not real, bound clusters, but simply unbound Scaled OB Associations. Those with a core are also called Super Star Clusters and may last long enough to become globular clusters in the distant future, but their survival depends on whether they can retain their halos. The (partial) halo retention in turn depends on the detailed star-formation history, since mass loss through winds and supernovae in the first 5 Myr of a cluster can make it expand significantly (Bastian and Goodwin, 2006). It has been commonly assumed that massive star clusters are formed instantaneously. However, studies of 30 Doradus have revealed that these objects can have complex starformation histories. In this region it is observed that not all stars were formed simultaneously (Walborn et al., 2002), an assumption that often underlies studies of massive star clusters and that is likely true only for the cluster core but not for the halo. Such multiple populations have been recently revealed in the long-term descendants of SSCs, globular clusters (Mackey et al., 2008). 1.1.2 Hii Galaxies Dwarf galaxies, the most numerous type in the Universe, can be classified into dwarf spheroidal/dwarf elliptical systems, dwarf irregulars (Magellanic systems), and active starforming dwarf galaxies, also known as blue compact dwarf galaxies (BCDs). This last group was first identified as a new class of galaxy by Haro (1956) and Zwicky (1966). Several years later, Sargent and Searle (1970) found that some of these objects have spectra similar to those of Hii regions in spiral galaxies. They called these new objects “isolated extragalactic H II regions”. In the bluest galaxies known before this discovery, as the early type spirals or irregulars, the stellar light was not dominated by the massive stars but by the oldish underlying population. However, the new type of galaxies did not show evidence of old population stars, indicating a high star formation rate. They have radius of less than 2 kpc, surface brightness µ v ≥ 19 mag arcsec −2 and total masses of less than 10 11 M ⊙ . Nowadays, they are referred to by different names, the most common being BCDs and Hii galaxies. Although these two terms are commonly used interchangeably, the term Hii galaxy is used when the objects have been selected for their strong, narrow emission lines in objective prism plates (Terlevich et al., 1991). On the other hand, BCDs have been defined according to different observational criteria: they are selected mainly for their blue colors and compactness. At any rate, both terms are used to design the same objects, since their observed properties in the optical are dominated by their Hii region spectrum. Nevertheless, in general the optical properties of Hii galaxies are dominated by one or a few Hii regions and their starbursts tend to be younger (Telles and Terlevich, 1995). Strictly speaking, not
Page 1: Facultad de Ciencias Departamento d
Page 5: A Abigaíl, mi ut A mi familia, por
Page 8 and 9: La gente del grupo de Astrofísica
Page 11 and 12: Contents List of Figures List of Ta
Page 13: CONTENTS iii B •Physical conditio
Page 16 and 17: vi LIST OF FIGURES 3.12 Example of
Page 19 and 20: List of Tables 1.1 Integrated prope
Page 21 and 22: Chapter 1 Introduction 1.1 Overview
Page 23: 1.1. Overview 3 Figure 1.1: Sample
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Page 60 and 61: 40 3 • IFS of a GEHR in NGC 6946
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54 3 • IFS of a GEHR in NGC 6946
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β[H 84 3 • IFS of a GEHR in NGC
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106 4 • Long-slit spectrophotomet
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Hδ 116 4 • Long-slit spectrophot
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Table 4.4 continued SDSS J1657 Knot
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150 Conclusions and future work has
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152 Conclusions and future work oth
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154 Conclusions and future work The
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156 Conclusiones Las abundancias to
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158 Conclusiones consistentes con l
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160 Appendix A where c=0.434C. This
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162 Appendix B where R S2 = I(6717
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164 Appendix B Sulphur As in the ca
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166 Appendix B Argon For argon, the
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Appendix C Empirical calibrators In
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C • Empirical calibrators 171 One
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Appendix D Stellar photometry resul
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D • Stellar photometry results of
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190 REFERENCES Bertelli, G., Bressa
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192 REFERENCES Girardi, L. & Bertel
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194 REFERENCES Kunth, D. & Sargent,
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196 REFERENCES Pérez-Montero, E.,
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198 REFERENCES Terlevich, R., Melni
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