Inuence of model parameters on optical emission-line properties .

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In this work (cloudy)

In this thesis, I adopt the latest version of the photoionization code cloudy (c13.03; described in Ferland et al. 2013)2 to compute the nebular emission from gas ionized by young stars in galaxies. cloudy is an open-source plasma simulation code, which models physical conditions within clouds, over full density and temperature ranges. The goal is to simulate the ionization, level populations, molecular state, and thermal state of material exposed to an external radiation eld or some other source of heating. To compute the synthetic spectrum of the ionized cloud, one has to specify the shape of the incident ionizing continuum { in my case, the spectral energy distribution computed using the galaxev stellar population synthesis code described in Section 2.1.5 { as well as the parameters of the diuse gas. Overall, the main input parameters that have to be specied to compute the radiation emerging from gas ionized by young stars in this approach with cloudy are:
• the nature of the external incident radiation striking the cloud: I will consider ionizing star clusters at the center of Hii regions, corresponding to simple stellar populations computed using the galaxev stellar population synthesis code (Section 2.1.5).
• the chemical composition and dust-grain content of the gas: cloudy handles in detail the lightest 30 elements at all stages of ionization, and I describe in Section 3.2.3 the individual abundances and factors of depletion onto dust grains I adopt in this thesis.
• the geometry of the gas (see Fig. 2.4): I use a closed geometry to compute the nebular emission from Hii regions, which means that the gas is assumed to be spherically distributed around the ionizing star clusters; instead, in Chapter 4, we use an open geometry to investigate the nebular emission from narrow-line emitting regions of AGN.
• the total hydrogen density of the gas: as detailed in Section 3.3.1, I investigate 4 hydrogen densities of the ionized gas, nH = 10, 102, 103 and 104 cm−3, spanning most of the range of observed electronic densities in extragalactic Hii regions (e.g., Hunt & Hirashita, 2009).

Transmission function of the ISM

To compute the transmission function T(t, t) of the ISM in equation (3.1), I follow CL01 (see also Pacifici et al. 2012) and write this as the product of the transmission functions of the ionized gas, T+ (t, t), and the neutral ISM, T0 (t, t), i.e. T(t, t) = T+ (t, t) T0 (t, t) . (3.2).
If the ionized regions are bounded by neutral material, T+ (t, t) will be close to zero at wavelengths blueward of the H-Lyman limit but greater than unity at the wavelengths corresponding
to emission lines. In this work, I focus on the nebular emission from star-forming galaxies, which is controlled primarily by the function T+ . I assume for simplicity that this depends only on the age t of the stars that produce the ionizing photons. Since 99.9 per cent of the H-ionizing photons are produced at ages less than 10Myr by a single stellar generation (e.g. Charlot & Fall, 1993; Binette et al., 1994), as in CL01, I write T+ (t, t) = T+ (t) for t 10 Myr, 1 for t > 10 Myr . (3.3). I do not consider in this chapter the attenuation by dust in the neutral ISM, which is controlled by the function T0 , but it is possible to include it in my models as explained in Section 3.2.4. I use the approach proposed by CL01 to compute the transmission function of the ionized gas in equation (3.3). This consists in describing the ensemble of Hii regions and the diffuse gas ionized by a single stellar generation in a galaxy with a set of effective parameters (which can be regarded as those of an effective Hii region ionized by a typical star cluster) and appealing to a standard photoionization code to compute T+ (t) at ages t 10 Myr for this stellar generation. By construction, the contributions by individual Hii regions and diffuse ionized gas to the total nebular emission are not distinguished in this prescription. This is justified by the fact that diffuse ionized gas, which appears to contribute around 20–50 per cent of the total H-Balmer-line emission in nearby spiral and irregular galaxies, is observed to be spatially correlated with Hii regions and believed to also be ionized by massive stars (e.g., Haffner et al., 2009, and references therein; see also Hunter & Gallagher 1990; Martin 1997; Oey & Kennicutt 1997; Wang, Heckman & Lehnert 1997; Ascasibar et al. 2016). The effective parameters describing a typical Hii region in this approach therefore reflect the global (i.e. galaxy-wide) properties of the gas ionized by a stellar generation throughout the galaxy. Fig. 3.1 illustrates the expansion in this way of a galaxy in a series of Hii regions through the convolution of the evolution of individual star clusters with a star formation history.

Interstellar abundances and depletion factors

The chemical composition of the ISM has a primary in uence on the transmission function T+ (t0) of equation (3.3). In this work, I adopt the same metallicity for the ISM, noted Zism, as for the ionizing stars, i.e., I set Zism=Z. A main feature of my model is that I take special care in rigorously parametrizing the abundances of metals and their depletion onto dust grains in the ISM, to be able to model in a self-consistent way the in uence of `gas-phase’ and `interstellar’ (i.e., total gas+dust-phase) abundances on the emission-line properties of a star-forming galaxy.

Optical emission-line properties

I have built a comprehensive grid of models spanning wide ranges of interstellar parameters, for stellar populations with a Chabrier (2003) stellar IMF with upper mass cutos 100 and 300M. I show that these models can reproduce available observations of starforming galaxies in several line-ratio diagrams at optical ([Oii]3726; 3729, H, [O iii]5007, H, [Nii]6584, [S ii]6717; 6731) and ultraviolet (Nv 1240, Civ 1548; 1551, He ii 1640, Oiii]1666, [Ciii]1907+Ciii]1909, [Si iii]1883+Si iii]1892) wavelengths.

Grid of photoionization models

My motivation is to build a grid of photoionization models that should be adequate for the purpose of investigating the emission-line properties of star-forming galaxies at all cosmic epochs. To this end, I adopt the following sampling of the main adjustable model parameters described in Section 3.2. I emphasize that these must be regarded as eective parameters describing the ensemble of Hii regions and the diused gas ionized by young stars in a galaxy: Interstellar metallicity, Zism: I consider 14 values of Zism between 0.6 per cent of and 2.6 times solar, corresponding to metallicities at which stellar population models are available (Section 3.2.1). These are Zism = 0:0001, 0.0002, 0.0005, 0.001, 0.002, 0.004, 0.006, 0.008, 0.010, 0.014, 0.017, 0.020, 0.030 and 0.040.


In uence of model parameters on optical emission-line properties

I now brie y describe the in uence of the main adjustable parameters of my model on the predicted optical emission-line properties of star-forming galaxies (see also CL01). In this description, I explore the eect of varying a single parameter at a time, keeping the other main adjustable parameters xed:
Interstellar metallicity. Fig. 3.4 (solid line) shows that increasing Zism at xed other parameters makes the [Oiii]/H ratio rise to a maximum (around Zism 0:006) and then decrease again. This is because gas cooling through collisionally excited optical transitions rst increases as the abundance of metal coolants rises, until the electronic temperature drops low enough for cooling to become dominated by infrared ne-structure transitions  (e.g., Spitzer, 1978). Ecient ne-structure cooling by doubly-ionized species n the inner parts of Hii regions also makes the [Oii]/[Oiii] ratio rise in Fig. 3.4d (Stasinska, 1980). The [S ii]/H ratio behaves in a similar way to the [Oiii]/H ratio (Fig. 3.4b). In contrast, my inclusion of secondary nitrogen production causes the [Nii]/H and [Nii]/[Oii] to rise steadily with metallicity (Figs 3.4a and 3.4c). I note that, because of my adoption of the same metallicity for the stars and the ISM, lowering Zism also leads to a harder ionizing spectrum (since metal-poor stars evolve at higher eective temperatures than metal-rich ones; e.g., g. 15 of Bressan et al. 2012). This has little in uence on the results of Fig. 3.4, which are largely dominated by the other eects described above.
Zero-age ionization parameter at the Stromgren radius. Fig. 3.4 shows that increasing US at xed other parameters makes the [Oiii]/H ratio rise and the [O ii]/[Oiii], [Nii]/H and [S ii]/H ratios drop. This is because increasing US at xed density nH and ionizing photon rate Q(0) in my model amounts to increasing the eective gas lling factor (equation 3.7 of Section 3.2.2), causing the Hii regions to be more compact and concentrated close to the ionizing star clusters. This strengthens the highionization [Oiii]5007 line relative to the lower-ionization [Oii]3727, [Nii]6584 and [S ii]6717; 6731 lines.
Dust-to-metal mass ratio. The eects of changes in d at xed other parameters are shown in Fig. 3.5. For clarity, I plot models for only a subset of 4 interstellar metallicities and 3 associated zero-age ionization parameters (using the dependence of US on Zism identied by Carton et al., in preparation; see equation 4.2 of Chevallard & Charlot 2016). Increasing d depletes metal coolants from the gas phase. The electronic temperature increases, as does cooling through collisionally excited optical transitions (e.g. Shields & Kennicutt, 1995). The implied rise in [Nii]/H and [S ii]/H ratios is signicantly stronger than that in [Oiii]/H ratio (Figs 3.5a and 3.5b), because oxygen is a refractory element strongly depleted from the gas-phase (making the [O iii]/H ratio drop as d increases), while S and N are both non-refractory elements (Table 3.2). Fig. 3.5 further shows that, not surprisingly, the eect changing d is more pronounced at high than at low Zism.

Table of contents :

1.1 The early Universe
1.1.1 First steps
1.1.2 Stars and galaxies
1.1.3 Chemical composition of galaxies
1.2 Nebular emission from ionized gas
1.2.1 Basic properties of HII regions
1.2.2 Photoionization and recombination processes
1.2.3 Spectra: lines and continuum
1.3 Emission-line diagnostics
1.3.1 Abundance determination in HII regions
1.3.2 Emission line-ratio diagrams
1.3.3 Other use of emission lines
1.4 Outline
2 Tools 
2.1 Stellar population synthesis codes
2.1.1 Generalities
2.1.2 Stellar evolutionary tracks
2.1.3 Stellar Initial Mass Function
2.1.4 Library of stellar spectra
2.1.5 In this work (galaxev)
2.2 Photoionization codes
2.2.1 Generalities
2.2.2 In this work (cloudy)
2.3 Conclusion
3 Modelling the nebular emission from star-forming galaxies 
3.1 Introduction
3.2 Modelling
3.2.1 Stellar emission
3.2.2 Transmission function of the ISM
3.2.3 Interstellar abundances and depletion factors
3.2.4 Dust in ISM
3.3 Optical emission-line properties
3.3.1 Grid of photoionization models
3.3.2 Comparison with observations
3.3.3 Inuence of model parameters on optical emission-line properties .
3.4 Ultraviolet emission-line properties
3.5 Limitations of standard methods of abundance measurements
3.5.1 The `direct-Te’ method
3.5.2 A case study: the C/O ratio
3.6 Conclusions
4 Comparison between star-forming galaxies and AGN 
4.1 Introduction
4.2 Photoionization models from AGN
4.2.1 Narrow-line regions of AGN
4.2.2 Dierences between AGN and SF models
4.3 Optical emission lines and standard AGN/star-formation diagnostics .
4.3.1 SDSS observational sample
4.3.2 [Oiii]5007 H versus [Nii]6584 H diagram
4.3.3 Other AGN/star-formation diagnostic diagrams
4.4 Ultraviolet emission lines and new AGN/star-formation diagnostics
4.4.1 Ultraviolet observational samples
4.4.2 Diagnostics based on the Civ1550, Heii1640 and Ciii]1908 emission lines
4.4.3 Nv1240-based diagnostics
4.4.4 Heii1640-based diagnostics
4.4.5 O-based diagnostics in the far and near ultraviolet
4.4.6 Ne-based diagnostics in the near ultraviolet
4.4.7 Distinguishing active from inactive galaxies in emission line-ratio diagrams
4.5 Ultraviolet line-ratio diagnostic diagrams of active and inactive galaxies
4.6 Conclusions
5 Linking my nebular emission modelling with observations 
5.1 Ultraviolet emission lines in young low-mass galaxies at z ‘ 2
5.1.1 Introduction
5.1.2 Observational sample
5.1.3 Photoionization modelling
5.1.4 Conclusion
5.2 Spectroscopic detections of CIII]1909 at z ‘ 6 􀀀 7
5.2.1 Introduction
5.2.2 Observational sample
5.2.3 Modelling the continuum and emission lines of A383-5.2
5.2.4 Conclusion
5.3 Spectroscopic detection of CIV1548 in a galaxy at z = 7:045
5.3.1 Introduction
5.3.2 Observational sample
5.3.3 A hard ionizing spectrum at z = 7
5.3.4 Conclusion
5.4 Ly and CIII] emission in z = 7 􀀀 9 galaxies
5.4.1 Introduction
5.4.2 Observational sample
5.4.3 Photoionization modelling
5.4.4 Conclusion
5.5 Conclusion
6 Conclusions 
A Nebular emission les 


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