Evolution of the dust-to-oxygen ratio

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Morphology

The well known Hubble classification makes use of morphological patterns, from elliptical, early-type galaxies to late-type galaxies. Elliptical galaxies have no or faded structures. They are classified into 7 classes, depending on the degree of ellipticity (earlier galaxies are perfectly round). After that, moving to late-type galaxies, two branches appear according to the presence or absence of a bar. Moving from early to late type, we firstly encounter lenticular galaxies. Their central region is similar to elliptical galaxies, but they have a large flat structure. Further down the sequence, spiral galaxies contains a bulge similar to an elliptic galaxy, and a full featured disk with spiral arms. Late spirals have no bulge, while early type ones have a prominent bulge.
Morphology is strongly affected by the environment as morphology–density (of galaxies) or morphology–radius (in clusters) relations show. Galaxies are believed to evolve from late-type to early-type, in the hierarchical scenario, through merging events and/or secular evolution.

Kinematics

New: Spiral galaxies have a rotating stellar and gaseous disk component. For instance, the Sun has a typical circular speed of 220 km s−1. Near the center, the circular speed increases with radius. Then the rotation curve flattens at high radius. The highest speed is related to the luminosity of the spiral galaxy via the Tully-Fisher relation.
New: For elliptical galaxies as well as bulges of spiral galaxies, the kinematics are described by the dispersion velocity of stars, since there is no global motion of stars. Typical dispersions are about 100 km s−1.

The interstellar medium

Tielens (2005) is a review of chemistry and physic in ISM.
The distance between stars in Solar neighbour is about 1 pc. This space is filled by ISM, mainly gas, with dust as solid state. It is of primary importance for galactic evolution. Indeed, stars born in cloud of collapsing gas and ISM evolves with feedback and metal enrichment from stellar winds and SNe remnant. Many objects reveal the richness of ISM and emphasizes its morphological complexity and the numerous physical processes inside ISM.
HII regions, like M42, are nebulosity ionized by the intense radiation of early-type (OB) and young stars. Their temperature is about 104 K at density > 10 cm−3 and their size are about 1 pc. Warm dust, heated, emits inside these regions.
Reflection nebulae are gaseous clouds illuminated by radiation field of neighbour stars. How-ever it is not sufficiently heated to be ionised and to emit themselves. Dust is also heated. Dark clouds are dense regions where dust could absorbs up to > 10 mag. They have a size of . 10 pc. They usually emit in IR but some can be also dark at these wavelength.
Photodissociation regions are usually at the interface between molecular and ionized phases. Their molecules and atoms receive a sufficiently intense level of UV (or even far ultraviolet) to dissociate molecules and ionize them. They show also an IR continuum due to their dust.
SNe leave materials, ejected by these highly energetic events. With time the surrounding ISM is shocked and it therefore forms SN remnants. These remnants have hot gas at about ∼ 106 K emitting X-rays and a synchrotron emission at radio wavelength.
ISM can be splitted in unmixed phases: cold neutral medium, warm atomic, either neutral or ionized, hot phase and MC.
Neutral medium is observed using 21 cm HI line and absorption of light passing through ISM from bright sources. Cold neutral medium has a temperature of ∼ 100 K in diffuse clouds of ∼ 10 pc diameter at density ∼ 50 cm−3. In the Galaxy, this phase is located in a thin disk having height of 100 pc. Intercloud medium is filled by a warm neutral phase at ∼ 8,000 K at ∼ 0,5 cm−3. It is a little thicker with a scale height of 220 pc and an observable tail unlike gaussian distributions. Although clouds contain 80 % of mass in disk plane between 4 and 8 kpc from center, these two phases, cold and warm neutral, have the same surface density on average. Warm ionized medium, also at ∼ 8,000 K, which is the transition temperature for ionisation of hydrogen, have a more extended altitude with scale height about 1 kpc, and a volume filling factor of 0.25. Hot ionised medium at 105–106 K replenishes the halo (scale height of 3 kpc). It is observed through continuum and lines in UV and X-ray wavelengths. It is also located in SN remnants. All these phases are generally assumed to be in pressure equilibrium.
On the contrary, MC are thought to be gravitationally bound. They are dense objects, > 200 cm−3 and cores in them can exceed 104 cm−3, at very low temperature: 10 K. GMC span a large range of properties. However a typical GMC has a size of 40 pc, a lifetime of about 3 × 107 yr and a mass of 4 × 105 M⊙. Although H2 is the most abundant molecule in them, they are usually observed using CO (1-0) transition line at millimeter wavelength. MC form from HI filaments thanks to thermal/gravitational instabilities and/or shock compression (Fukui & Kawamura 2010). MC are therefore surrounded by HI envelope, and theoretical studies emphasizes on pressure and radiation field for HI conversion to H2. HI filaments are themselves formed by supershells and from density waves in spiral galaxies. GMC mass distribution follows the law: dn ∝ m−α, (2.2) dm .

Dwarf galaxies

In order to describe properties of dwarf galaxies, I use two reviews, one from Mateo (1998) and the other from Tolstoy et al. (2009).
Although definition varies with authors, dwarf galaxies are considered as low-luminosity galaxies (MB & −16, MV & −17), more extended than globular clusters (half-light radius r1/2 ≥ 1.6 pc) and having a dark halo. This definition has not a physical meaning because there is not a clear distinction with less massive disk systems. Dwarf galaxies are well studied in the Local Group (Milky Way and M31), where ∼40 galaxies of such type have been found, and other ∼ 20 galaxies are expected and not yet been found because of their low galactic latitude. Well known exemples are the Small and Large Magellanic Cloud (SMC and LMC), neighbours of our Galaxy at 58 and 49 kpc respectively. Dwarf galaxies of the Local Group extend to 1, 600 ± 200 kpc (for IC5152), which is approximatively the size of the Local Group (∼ 1.8 Mpc) when considering dynamical approach using zero velocity surface1. Dwarf galaxies are the most numerous galaxies and own a large fraction of a cluster mass.
Fig. 1 of Tolstoy et al. (2009) allows to distinguish different classes of objects. Note that M32 falls in the definition of dwarf galaxies but looks like a low-luminosity elliptical, indeed it is a compact galaxy. Dwarf spheroidal (dSph) are early-type galaxies and dwarf irregular (dIrr) are late type, while transition type are at an intermediate state. Blue compact disk (BCD), not observed in Local Group, are less extended and have a high SF. There are ultrafaint dwarf, which seems to be low-luminosity dSph, more metal poor, dark matter dominated, thus, probably not globular clusters but galaxies, with low surface brightness. We also observe ultra compact dwarf similar in size to globular clusters, observed in other galaxy clusters.
dIrr are optically dominated by star-forming regions and OB associations with clumpy mor-phology. Background of older stars emits a more extended, smooth, and symmetric component. DSph are smooth, few have nuclei and M32 is believed to have a black hole. Profiles are generally fitted with King model, preferentially for earlier type, or exponential model for late type. dSph appear as gas-free dIrr, due to SF and/or ram-pressure stripping, but discrepancies remain on central surface brightness or structures (Hunter & Gallagher 1985, Bothun et al. 1986, and James 1991). If a dwarf galaxy lives near a massive one, it can have tidal distortions: stars become unbounded to their galaxy. Then the host galaxy elongates and finally can disrupt the dwarf galaxy to form streams and partially populate the halo of the massive galaxy.
SFH of dwarf galaxies is derived from photometry, by using different classes of stars, each being witness for a typical age of SF: Wolf-Rayet for the past 10 Myr, Blue-loop stars for 100– 500 Myr, red supergiants for 10–500 Myr, AGB for ∼1 Gyr, RGB for & 1 Gyr, red-clump and horizontal branch stars for 1–10 Gyr and & 10 Gyr respectively, subgiant branch stars for &2– 4 Gyr, main sequence with the help of subgiant for &1–2 Gyr with a good resolution (∼ 1 Gyr). SFH is also built by comparing synthetic color-magnitude diagram with observed one. This latter method is affected by age-metallicity degeneracy.
Dwarf galaxies have various moderate SF with short quiescent phases. Old population for dIrr are younger than 10 Gyr, recent SF display short bursts (10–500 Myr duration). Due to their gas-poor nature, dSph have not formed stars since at least 100 Myr ago. Transition type have HI gas but no SF. BCD display a recent and high SFR.
dIrr galaxies can have gas-to-total mass ratio from 7% to 50% while transition galaxies typ-ically have 1–10% and spheroidals ≤ 0.1%. New: Since we will study especially dust, we are interested in gas rich galaxies. Although dIrr display HI clumps of 100–300 pc generally located near SF regions, HI emission is generally centered with optical emission and is more extended. Dust as well as CO clouds in dIrr have generally a size of 20–40 pc and a mass of few 100 M⊙, near SF regions, dust is also found near the core of some dSph. HII regions . 200–400 pc are detected in all dIrr, whereas no HII is detected in most galaxies of earlier types.
Chemical abundances are estimated from photometry or spectroscopy of RGB stars or emis-sion lines in HII regions or planetary nebulae. For stellar component, [Fe/H] is underabundant (< −1 dex Tab. 6 of Mateo 1998) and has significant dispersion in a galaxy (from ∼ 0.3 ± 0.1 to ∼ 0.7 ± 0.2 dex). For HII regions, oxygen abundances is well studied, dispersion seems to be insignificant, implying an uniform abundance, thus a good mixing from gaseous hot phase. There is a bimodal correlation (Fig. 7 from Mateo 1998) between luminosity and metallicity, for dIrr on one side and dSph and transition type on other side. The bimodality is expected to be due to SF in dIrr.
From radial scale length and central surface brightness, velocity dispersions should be ≤ 2 km × s−1 while observations give > 7 km × s−1. Moreover, rotating galaxies have kinematic not consistent with mass derived from their luminosity. To be in equilibrium, dwarf galaxies must have an important mass of dark matter.
Oh et al. (2011) derive rotation curves as well as gaseous mass distribution of THINGS dwarf galaxies from HI data. They assume a stellar mass distribution from infrared and visible data. Then they compute dark matter profile and find it is better described by a core model. They therefore choose to use pseudo-isothermal halo model.

Production and destruction

The model is fully exposed in Sect. 3.2. We present here an introduction of the processes included to compute the evolution of dust mass in simulations (see Fig. 2.1).
As we have seen in the previous section, three main production processes have to be included: production from AGB stars, from SNe, and growth in ISM. Destruction occurs through shocks produced by SNe. On contrary to the semi-analytic model, in order to design our model, we have paid attention to have local description of all phenomena thanks to numerical simulations. AGB stars highly enrich the ISM. Indeed, when thermally pulsing, they produce winds suf-ficiently cold, where materials condense and form dust grains. SNe are also thought to produce dust. In sufficiently dense ISM like molecular clouds where grains are also protected from de-structive radiation, growth of grains, by accreting atoms and molecules on surface, allows for dust production. Destruction of grains is theoretically known. The main process is the destruction through shocks. These shocks, produced by highly energetic events like exploding SNe, submit grains to sputtering (gas-grain interaction) process and grain-grain collisions. Finally, when stars form in cores of molecular clouds, dust is incorporated in forming stellar system and/or destroyed by the radiation of new stars. Thus, dust is likely to be destroyed with star formation. Since dust production in ISM depends on the density of gas, we expect variations of abundances according to the phases of ISM. Then, there is also dust transfers between these phases.

Table of contents :

1 Introduction 
1.1 Historical review
1.2 Dust
2 Galaxies and dust 
2.1 Galaxies
2.1.1 Introduction
2.1.2 Morphology
2.1.3 Kinematics
2.1.4 Stars
2.1.5 The interstellar medium
2.2 The Milky Way
2.2.1 The stellar component
2.2.2 The gaseous component
2.3 Dwarf galaxies
2.4 Dust
2.4.1 Effects on radiation
2.4.2 Composition
2.4.3 Box model
2.4.4 Production and destruction
2.4.5 Observations
3 Chemodynamical Code 
3.1 Overview
3.1.1 Interaction of gas and stars
3.2 Dust Mass Evolution
3.2.1 The model
3.2.2 Production
3.2.3 Accretion
3.2.4 Destruction
4 First Result: The Milky Way 
4.1 Simulation
4.2 Analysis
4.2.1 Radial distribution of dust
4.2.2 Vertical distribution of dust
4.2.3 Evolution of the dust-to-gas ratio
4.2.4 Evolution of the dust-to-oxygen ratio
4.2.5 Dust–metallicity diagram
4.2.6 Other diagrams
5 Second Result: The Dwarf Galaxies 
5.1 Simulations
5.1.1 Dark matter
5.1.2 Stellar and gaseous components
5.1.3 Summary
5.2 Analysis
5.2.1 Evolution of the dust-to-oxygen ratio
5.2.2 Maps of the dust-to-oxygen ratio
5.2.3 Local correlation
5.2.4 Global correlation
6 Conclusion 
A Improvement 
B Parallelization 

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