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

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96 3 • IFS of a GEHR in NGC 6946 Object ID N(WR) WR/O Knot A 125 ± 10 0.047 ± 0.008 Knot B 5 ± 2 0.002 ± 0.001 Knot C 22 ± 3 0.008 ± 0.001 Table 3.11: Derived number of WR stars and ratios between WR and O stars. of star formation. The only exception is knot B, where the equivalent width is slightly higher than predicted by the models. However, the observed relative intensities are inside in the expected range of theoretical values. From both quantities, an age around 4 Myr fits well for the three knots, indicating that these bursts are very young. As already told in section 3.4.7, a red bump has been marginally detected in knot A, indicating the existence of WC stars, but no accurate measurements were made due the low signal-to-noise, so only the total number of WR (WN+WC) can be estimated. Table 3.11 shows the number of O and WR stars. We have derived these numbers by comparing the dereddened luminosities of Hα and of the WR blue bump with the predictions of the instantaneous model of Starburst 99. As it can be seen, knot A contains more than one hundred WR stars, as expected from the strong blue bump feature from its integrated spectra. Interestingly, Knot B has a few WR stars, despite its weak contribution in the continuum map of Figure 3.15. Other way to estimate the number of WR stars is by means of calibration of WR stars. Smith (1991) estimates that a single WN7 star emits 3.2 × 10 36 ergs s −1 at 4650Å, which yields a total number of 120 ± 10 WR stars, very close to the value derived from the Starburst 99 model. On the other hand, Vacca and Conti (1992) estimate that one WN7 star emits 1.7 × 10 36 ergs s −1 , rising the number of WR stars almost by a factor of 2. 3.5.3 Metal content We derived oxygen, sulphur, nitrogen, neon, argon and iron chemical abundances in the integrated spectrum of the PPak-field and in three of the four knots, where auroral lines of [Oiii], [Oii] and [Siii] have been measured and thus their respective electron temperature derived. Due to the quality of the data, the electron temperatures derived and thus, the metallicities, have associated errors comparable to the dispersion found in empirical calibrations. The derived temperatures (see Table 3.7) for the three knots and the integrated spectrum of the PPak-field show values very similar for all of them, centered in 8000K for t([Oiii]). Knot A shows a higher temperature, as expected from its Hα luminosity, but similar to the other measurements within the errors. Similarly, all the values for t([Oii]) are very close. The only exception is the one estimated for knot B from a relation with t([Oiii]) based on
3.5. Discussion 97 photoionization models. This value is a little bit higher than the mean value of the other direct measurements. It is known that the relation between t([Oii]) and t([Oiii]) does not present a clear trend, showing a large scatter. The [Oii] temperature derived from relations based on the photoionization models described in Pérez-Montero and Díaz (2003) covers the range of values for a sample of Hii galaxies, GEHRs and diffuse Hii regions in the Galaxy and the Magellanic Clouds, as shown in Hägele et al. (2006). Higher density models show lower values of t([Oii]) for a given t([Oiii]), being this effect more notable at high electron temperatures (see Figure 4.12 of Chapter §4). Our objects lie in the low temperature region, where N e = 100 and N e = 500 models are very close. The derivation from t([Oiii]) gives a slightly higher temperature than the directly derived, but still within the errors for our sample. A worse situation is presented for the relation based on the photoionization models by Stasińska (1990), which could yield a higher [Oii] temperature by 2000 K for the regime of our data. These differences may translate into lower O + /H + ratios, and therefore lower total abundances by an amount of 0.35 dex, taking as an example knot A. Nevertheless, the differences between the derived temperatures from Pérez-Montero and Díaz (2003) and the directly measured give abundances compatible within the errors. Regarding the [Siii] temperature, the only direct temperature available we have is for Knot A, 7300 ± 1200 K. The value derived from the empirical relation found by Hägele et al. (2006) using t([Oiii]) is 7400 ± 1300 K, which matches perfectly the direct measurement. In this case, the relation found by Garnett (1992) gives a [Siii] temperature 1400 K higher. As it has been told, we have used the former relation to estimate t([Siii]) for the other knots, in which the [Siii] λ 6312 Å auroral line could not be measured. The abundances derived for the integrated spectra show values near the solar ones. The abundances of various elements in Hii regions in NGC 6946 have been studied by McCall et al. (1985) and Ferguson et al. (1998). At the distance of the Hii complex discussed in this paper studies find oxygen abundances very similar to solar. The galaxy shows a clear abundance gradient as a function of distance from the center. Belley and Roy (1992) estimated the global oxygen abundance gradient of NGC 6946 by means of imaging spectrophotometry in the nebular lines Hα, Hβ, [Nii], and [Oiii]. Using the empirical calibration of Edmunds and Pagel (1984), they calculated O/H abundances for Hii distributed throughout the disc of the galaxy. For NGC 6946, they found that the global gradient is given by ∆ log (O/H) / ∆R = -0.089 ± 0.003 dex kpc −1 , and extrapolated central abundances of 9.36 ± 0.02. This value is the same as the one estimated by Ferguson et al. (1998) using emission line spectra and the empirical relation of McGaugh (1991). At the projected galactocentric distance of our Hii complex (6.85 kpc), this corresponds to 12 + log(O/H) = 8.75, 0.1 dex higher than our mean direct measurement, but within the errors. Larsen et al. (2006), by means of H and K band spectra for a young luminous super star cluster at 4.8 kpc, derived an (direct) abundance of 12 + log(O/H) = 8.76, about 0.2 dex lower values of Hii regions at their galactocentric distance. To our knowledge, there is no other direct abundance determinations of our Hii
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Facultad de Ciencias Departamento d
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A Abigaíl, mi ut A mi familia, por
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La gente del grupo de Astrofísica
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Contents List of Figures List of Ta
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CONTENTS iii B •Physical conditio
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vi LIST OF FIGURES 3.12 Example of
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List of Tables 1.1 Integrated prope
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Chapter 1 Introduction 1.1 Overview
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1.1. Overview 3 Figure 1.1: Sample
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1.1. Overview 5 These massive star
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1.1. Overview 7 Figure 1.3: Color-m
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1.2. Aims and structure of this wor
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12 2 • The star formation history
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40 3 • IFS of a GEHR in NGC 6946
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146 4 • Long-slit spectrophotomet
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148 4 • Long-slit spectrophotomet
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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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