TRACING ABUNDANCES IN GALAXIES WITH THE SPITZER ...

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(ICF) to account for it. A program written by Jeronimo Bernard-Salas solves for the population of the levels and ionic abundances following equations found, for example, in Osterbrock (1989). This program assumes a five level atom which serves as a good approxi- mation because the upper levels are only sparsely populated. The populations of the levels are found by assuming statistical equilibrium for each level i: � j�=i njNeqij � �� collisional (de)excitation rate + � j>i njAji � �� radiative transition rate from all upper levels = � j�=i niNeqij � �� collisional (de)excitation rate + � j
(Iion), the radiative transition rate for the line between the upper and lower levels (Aul), the total population in the upper level (Nion nu where Nion is the total population density of the ion and nu is the ratio of the population density of the upper level from which the line originates to Nion), and the energy of the line (hνul): Iion ∝ Aul Nion nu hνul. (1.4) Additionally we use the relation between the intensity of the extinction corrected Hβ line (IHβ ), the effective recombination coefficient for Hβ (αHβ), the density of protons (NH +), the density of electrons (Ne), and the energy of the line (hνHβ ): IHβ ∝ αHβ NH + Ne hνHβ . (1.5) Taking the ratio of these intensities and re-arranging, ionic abundances with re- spect to hydrogen (Nion/NH +) are given by (Bernard-Salas et al., 2001): Nion NH + = Iion Ne λul αHβ IHβ nu λHβ Aul . (1.6) The bulk of the program for determining abundances is then in determining nu (the fraction of the total population of the ion in the level from which the line originates) according to Equation 1.1. Abundances derived from infrared lines have several advantages over those determined from optical lines (Rubin et al., 1988; Pottasch & Beintema, 1999). (1) There is less extinction in the infrared than the optical, and thus extinction- corrected line fluxes (from which abundances are derived) are more accurate in the infrared than in the optical because uncertainties in the extinction coefficient and law affect the infrared line fluxes less. This is especially true in areas of high extinction, such as the Galactic Bulge. (2) Infrared lines arise from levels close to the ground level and thus abundances determined from them do not depend as much on the adopted electron temperature as abundances determined from optical 18
Page 1 and 2: TRACING ABUNDANCES IN GALAXIES WITH
Page 3 and 4: TRACING ABUNDANCES IN GALAXIES WITH
Page 5 and 6: BIOGRAPHICAL SKETCH Shannon was bor
Page 7 and 8: ACKNOWLEDGEMENTS First and foremost
Page 9 and 10: thank my husband Ryan. The rest of
Page 11 and 12: 3 Chemical Abundances and Dust in P
Page 13 and 14: LIST OF FIGURES 1.1 Diagram of the
Page 15 and 16: ditionally we find that these nebul
Page 17 and 18: seed nuclei relative to the β deca
Page 19 and 20: shape (Carollo et al., 2007). These
Page 21 and 22: Boltzmann velocity distribution cha
Page 23 and 24: etween recapture and photoionizatio
Page 25 and 26: a whole. Additionally, abundances f
Page 27 and 28: at the appropriate temperature and
Page 29: (a) 3.37 eV 1.40 eV 0.10 eV 0.04 eV
Page 33 and 34: lines. For example, changing the ad
Page 35 and 36: Figure 1.4 Starburst99 synthetic sp
Page 37 and 38: Figure 1.5 Example of crystalline s
Page 39 and 40: H H H H H H H H H Anthanthrene C H
Page 41 and 42: year lifetime. Spitzer is in an Ear
Page 43 and 44: CHAPTER 2 THE SPITZER IRS INFRARED
Page 45 and 46: these elements than previous studie
Page 47 and 48: dust is seen in the spectrum. We se
Page 49 and 50: Figure 2.2 Close-ups of emission li
Page 51 and 52: 2.4 Optical and UV Data We compleme
Page 53 and 54: 2.5 Data Analysis Our goal is to ca
Page 55 and 56: 2.5.2 Electron Density We assume an
Page 57 and 58: Table 2.6 Electron temperatures ass
Page 59 and 60: O, N, and C. The helium ionic abund
Page 61 and 62: Barlow (1994) use UV lines for impo
Page 63 and 64: is a spherical nebula with a low ma
Page 65 and 66: CHAPTER 3 CHEMICAL ABUNDANCES AND D
Page 67 and 68: essential data on important ionizat
Page 69 and 70: are members of the Bulge, (Acker et
Page 71 and 72: Table 3.2 Properties of observed GB
Page 73 and 74: nod-averaged line fluxes. Uncertain
Page 75 and 76: Table 3.3 Multiplicative scaling fa
Page 77 and 78: Table 3.4 Observed infrared line fl
Page 79 and 80: 3.5 Data Analysis 3.5.1 Abundances
Page 81 and 82:
Table 3.6 Comparison of the derived
Page 83 and 84:
from the most reliable line(s), and
Page 85 and 86:
Table 3.7 Electron temperatures and
Page 87 and 88:
Table 3.9 Ionic abundances Ion x a
Page 89 and 90:
3.5.2 Crystalline Silicates Crystal
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argon, neon, and sulfur tend to be
Page 93 and 94:
Figure 3.2 Continuum subtracted spe
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Figure 3.3 Profiles of PAH features
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mean argon and sulfur abundances ar
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of two smaller than that of Simpson
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towards the Galactic Center. On the
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Table 3.14 Parameters of linear fit
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models that show how a binary compa
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y an outer dense O-rich torus, indi
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CHAPTER 4 ABUNDANCES IN H II REGION
Page 111 and 112:
expect their abundance ratio to rem
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We extract spectra processed throug
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Figure 4.2 IRS spectra of H II regi
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IRS slits from these WIRC observati
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S IV) we employ is equal to Te(O +2
Page 121 and 122:
Figure 4.3 M51 H II region Ne III/N
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Figure 4.4 shows a plot of the (Ne
Page 125 and 126:
James Webb Space Telescope (JWST);
Page 127 and 128:
Bernard-Salas, J. & Tielens, A. G.
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Edgar, R. G., Nordhaus, J., Blackma
Page 131 and 132:
Hummer, D. G., Berrington, K. A., E
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Matsuura, M., Zijlstra, A. A., Mols
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Pradhan, A. K., Montenegro, M., Nah
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Vassiliadis, E. & Wood, P. R. 1993,
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