Spectral lines stellar

Our multi-level carbon model atom is adapted from D. Kiselman (private communication), with improved atomic data and better sampling of some absorption lines. The statistical equilibrium code MULTI (Carlsson 1986), together with ID MARCS stellar model atmospheres for a grid of 168 late-type stars with varying Tefj, log g, [Fe/H] and [C/Fe], were used in all Cl non-LTE spectral line formation calculations, to solve radiative-transfer and rate equations and to find the non-LTE solution for the multi-level atom. We put particular attention in the study of the permitted Cl lines around 9100 A, used by Akerman et al. (2004). 

Abstract. We present preliminary results of 3D hydrodynamical simulations of surface convection in red giants stars. We investigate the main differences between static ID and 3D time-dependent model stellar atmospheres of red giants for a range of metallicities between solar and [Fe/H] = —3 focusing in particular on the impact of 3D spectral line formation on the derivation of stellar abundances. 

In this contribution we present preliminary estimates of the effect of convection in 3D model stellar atmospheres of red giants on the formation of spectral lines and on the derivation of chemical abundances in stars. 

We then use a Feautrier scheme [4] to perform spectral line formation calculations in local thermodynamic equilibrium approximation (LTE) for the species indicated in table 1. At this stage we consider only rays in the vertical direction and a single snapshot per 3D simulation. Abundance corrections are computed differentially by comparing the predictions from 3D models with the ones from ID MARCS model stellar atmospheres ([2]) generated for the same stellar parameters (a microturbulence = 2.0 km s-1 is applied to calculations with ID models). 

EXO 0748-676, Cottam et al. (2002) have found absorption spectral line features, which they identify as signatures of Fe XXVI (25-time ionized hydrogenlike Fe) and Fe XXV from the n = 2 —> 3 atomic transition, and of O VIII (n = 1 —> 2 transition). All of these lines are redshifted, with a unique value of the redshift z = 0.35. Interpreting the measured redshift as due to the strong gravitational field at the surface of the compact star (thus neglecting general relativistic effects due to stellar rotation on the spectral lines (Oezel Psaltis 2003)), one obtains a relation for the stellar mass-to-radius ratio ... 

Unlike stellar spectroscopy, the analysis of meteoritic grains and inclusions can provide an extremely precise isotopic breakdown. The weak point of this technique, however, is that the exact characteristics of the stars from which the grains formed can only be inferred. When we detect fight, we can deduce its celestial source by extending back its fine of incidence and we can determine the composition of the source from the spectral lines it contains. But we do not know where the meteorite grains came from, and only their composition can tell us anything of their origins. 

When Payne began her work in the 1920s, stellar spectroscopy was a very active area of research. Numerous elemental and molecular lines had been identified in stellar spectra. The lines observed in each star varied with the inferred temperature of the star, which was understood to mean that the elemental abundances varied with temperature. This body of data was the basis for the spectral typing of stars ( , B, A, F, G, , M, L). However, the power source for stars was not understood and it was not clear why the composition of a star should be related to its temperature. In the 1920s, it was also widely believed that the Sun had the same composition as the Earth models considered the Earth to have formed from the outer layers of the Sun. Payne used the new guantum mechanical understanding of atomic structure to show how and why the spectral lines of the different elements varied as a function of stellar spectral type. She demonstrated how the temperature of the stellar surface controls the spectral lines that are observed. Her analysis led to the conclusion that the chemical... 

The last few years have seen a minor revolution in determining solar and stellar abundances (Asplund, 2005). Much of the previous work assumed that the spectral lines originate in local thermodynamic equilibrium (LTE), and the stellar atmosphere has been modeled in a single dimension. Since 2000, improved computing power has permitted three-dimensional modeling of the Sun s atmosphere and non-LTE treatment of line formation. The result has been significant shifts in inferred solar abundances. 

Shapes or profiles of spectral lines are often very important for obtaining information about the physical system under observation. Thus a spectral profile can tell us the state of a stellar atmosphere, or it can be used as diagnostic for a plasma. The line shape of an NMR spectrum may disclose the internal motion of nuclei in a molecular system. The information thus obtained is admittedly not extremely detailed, but it may be very useful particularly when combined with other information, and often it is the only data available. 

X-ray spectroscopy has also been applied to the interpretation of solar spectra, which are emitted by solar flares. Now stellar objects are under investigation by X-ray satellites such as Chandra and XMM. Whereas the present X-ray telescopes are medium resolution devices, the next generation (Constellation-X, XEUS) will provide sufficient spectral resolution for detailed analysis. The spectra from distant object usually suffer from low statistics solar flares have low emission time and the observation time of stellar objects is limited. In addition, the electron distribution is not Maxwellian, in general, and some of the spectral lines may be polarized. Therefore, verified theoretical data are of great importance to interpret solar and stellar spectra, where they provide the only source of information on the plasma state. 

It often takes time for the implications of experimental data to be understood and to be acted upon. Fraunhofer s earlier observation that the solar D-lines coincided with the spectral lines of a sodium lamp eventually prompted further important experiments. In 1849, Jean Bernard Leon Foucault (1819-1868), a Parisian physicist, made an unexpected discovery. He passed sunlight through a vapor of sodium and he found that the solar D-lines were darker. His conclusion was that the sodium vapor presents us with a medium which emits the rays D on its own account, and which absorbs them when they come from another quarter. The consequences of Foucault s experiment, however, were expressed more cogently by Sir William Thomson (later Lord Kelvin). He drew the following explicit conclusion That the double line D, whether bright or dark, is due to the vapor of sodium. . . That Fraunhofer s double dark line D, of solar and stellar spectra, is due to the presence of vapor of sodium in atmospheres surrounding the Sun and those stars in whose spectra it has been observed. ... 

As mentioned, stellar spectra with right circular polarization obtained in the first exposure and left circular polarization in the second exposure are projeeted in turn on the same section of the CCD detector. Thus, errors in the flat-fielding procedure for two spectra with opposite eireular polarization are practically the same and do not affect the calculation of GMF in the ease of weak magnetic fields. Additionally, this observational technique automatieally allows us to rale out shifts of spectral lines caused by inaccurate adjustment of the CCD plane to the focal plane of the spectrograph and instrumental drift of contours of spectral lines during the second exposure relative to the first one. 

Figure 4. Mean longitudinal magnetic field strength (Be) of p CrB vs. phases of stellar rotational period. Solid circles Be measurements by Wade [47] using the bulk of spectral lines. Open triangles up are Crimean measurements obtained using Fe I 6136.615 A spectral line. Open circles are our5j. measurements obtained using Fe I 6137.692 A spectral line. Open triangles down are our Be measurements obtained using Ca I 6162.173 A spectral line.

G. KirohhofE, J. N. Lockyer, and F. MoClean, reported that the spectral lines of chromium appear in the solar or in stellar spectra. H. Deslandies also found chromium lines in the ultra-violet spectrum of the corona. 

The intensities of spectral lines depend not only on the population density of the molecules in the absorbing or emitting level but also on the transition probabilities of the corresponding molecular transitions. If these probabilities are known, the population density can be obtained from measurements of line intensities. This is very important, for example, in astrophysics, where spectral lines represent the main source of information from the extraterrestrial world. Intensity measurements of absorption and emission lines allow the concentration of the elements in stellar atmospheres or in interstellar space to be determined. Comparing the intensities of different lines of the same element (e.g., on the transitions Ei Ek and Ee -> Ek from different upper levels Ei, Ee to the same lower level Ek) furthermore enables us to derive the temperature of the radiation source from the relative population densities A/, Ne in the levels Ei and Ee at thermal equilibrium according to (2.18). All these experiments, however, demand a knowledge of the corresponding transition probabilities. 

Stellar spectra normally consist of absorption lines which occur when the intense continuum radiation from the hot interior of the star is filtered on passing through the cooler outer stellar atmosphere. The strength of the absorption line is a measure of the abundance of the element. As a measure in the determination of the number of absorbing atoms the so-called equivalent width of the spectral line is used. The equivalent width is defined as the width of a square box that covers the same surface as the actual absorption profile of the line, as illustrated in Fig.6.79. An example of an experimentally recorded stellar line is included in the figure. 

The Doppler effect causes displaced spectral lines from objects in motion. The rotation of a star around its axis causes an extra broadening of the spectral lines, as does turbulence in the stellar atmosphere. A detailed analysis of the line profiles can clearly yield interesting information on the physical condition of the object. 

Spectroscopy is the key to unlocking the information in starlight. Stellar spectra show a variety of absorption lines which allow a rapid classification of stars in a spectral sequence. This sequence reflects the variations in physical conditions (density, temperature, pressure, size, luminosity) between different stars. The strength of stellar absorption lines relative to the continuum can also be used in a simple way to determine the abundances of the elements in the stellar photosphere and thereby to probe the chemical evolution of the galaxy. Further, the precise wavelength position of spectral lines is a measure of the dynamics of stars and this has been used in recent years to establish the presence of a massive black hole in the centre of our galaxy and the presence of planets around other stars than the Sun. 

In addition to this spectral class, stars are also characterized by a luminosity parameter. This luminosity classification is made on the basis of the width of spectral lines. Table 2 summarizes this classification. The width of spectral lines increases as the gas pressure increases. This so-called pressure broadening is due to the perturbation of atomic energy levels by other, nearby species. The physically largest stars have the lowest surface densities and pressures. Lines from these stars are therefore broader than from smaller stars (Figure 1 and Table 2). This difference in size, which results in a difference in stellar luminosity, has led to the naming scheme from supergiants to dwarfs. 

A final way of measuring stellar fields is by looking at the overall broadening of spectral lines of different Lande factor and Zeeman structure to discover broadening in excess... 

The physical nature of the turbulent component of a spectral line in a molecular cloud is currently a source of considerable debate. Physical processes that have been suggested as sources of the turbulence in molecular clouds are expanding HII regions, supernova remnants, cloud-cloud collisions, galactic differential rotation, and stellar winds. Unfortunately, for all of these processes there are theoretical problems with couphng the energy produced into turbulence. 

We have shown in previous chapters that the -values of spectral lines are important fundamental data which must be known before detailed calculations of the behaviour of gas discharges, plasmas, or stellar atmospheres can be undertaken. Since it is difficult, in many cases, to make theoretical calculations of f-values to an accuracy of better than 20 per cent, experimental measurements of these quantities are essential. A considerable number of different techniques have been developed for this purpose, many of them involving the determination of radiative lifetimes. In this chapter we discuss two such techniques, namely the beam-foil and the delayed-coincidence methods. In Chapter 8 we shall discuss the determination of the f-values of resonance lines by studies of the profiles of spectral lines and in Chapters 15 and 16 the use of the Hanle effect and optical double resonance methods. 

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