Lecture 3: Detailed Stellar Abundances

Lecture 3: Detailed Stellar Abundances

Eline TolstoyKapteyn Astronomical Institute,

University of Groningen

Abundances and Kinematics from High-Resolution Spectroscopic Surveys

Eline Tolstoy Third Chilean Advanced School of Astrophysics Concepción, 8-12 January 2007

I have a spectrum, and colours...

D=79 kpcrcore=5.8'

130 candidates


Part 3a: Measuring AbundancesMeasurements of equivalent widths can be used to derive the abundance ofelements for which we can detect and measure absorption lines in the gasmaking up the stellar atmosphere.

Letarte 2007, PhD thesis

The Boltzmann and Saha equations are applied and combined with thepressure and temperature of the gas to derive an abundance of the element.


Obtaining information from spectraChemical analysis usually proceeds using a curve-of-growth technique eitherexplicitly or implicitly. The alternative is to synthesize the whole region.

Literature: Alot of the detailed work presented here comes from the Phdof Bruno Letarte, University of Groningen, 2007 - you can ask for a copyif you are interested in all the details mentioned here, and somethingsfrom yesterday: [email protected]


No. of absorbing atoms per unit areaTo find the number of absorbing atoms per unit area, Na, that have electrons in theproper orbital to absorb a photon at the wavelength of the spectral line - the T, p areused in the Boltzmann and Saha equations to calculate the excitation and ionisation.

This task is complicated by the fact that not all transitions betweem atomic statesare equally likely. Each transition has a relative probably, or f-value (also calledoscillator strength).These can be calculated theoretically or measured in a lab, and they are defined sothat the f-values for transitions from the same orbital add up to the number ofelectrons in the atom or ion, i.e., the effective number of electrons per atomparticipating in a transition.Multiplying the number of absorbing atoms per unit area by the f-value gives thenumber of atoms lying above each cm2 of photosphere.


Voigt Profiles of K line of Ca II

For varying numbers of absorbing Ca ions


Curve of GrowthThis an important tool to determine Na, the number of absorbing atoms, and thusthe abundances of elements in stellar atmospheres, because EW varies with Na.Itis a log-log plot of EW as a function of the number of absorbing atoms.


So, practically - Low Mass Stars (RGB):

1. Spectral type 2. Atmospheric properties (line formation)

Effective TemperatureSurface GravityMetallicity

What we need to start:

Then:

1. Interpolate published model atmosphere grids (ormake own)

2. Compute theoretical (synthetic) spectrum for assumedabundances

3. Compare with observations and iterate onabundances

TO GET AN ABUNDANCE


Stellar Parameters

Surface Gravity, log10g

Metallicity, Z, or [Fe/H]

Effective Temperature, TeffTemperature of Black-body with the same L and R as the real star


How to determine TeffPhotometry: An empirical method, calibrated on the InfraRed Flux Method,(Blackwell & Lynas-Gray 1998), gives a relation between photometriccolours (like V I, V K) and Teff. The general method and correctionpolynomials are described in the series of papers by Ramirez & Melendez(2005); Alonso et al. (1999a,b, 2001), and references therein.

Excitation Equilibrium: define Teff such that the abundance of an element isindependent of the excitation potential (χex) of the individual lines. to usethis method, we need many lines of a single element sampling a range ofχex. The precision with which Teff can be determined depends upon theresolution, the choice and number of lines and signal to noise of eachspectrum used.

Teff +400K


Microturbulence Velocity, vt

Can be important in strong lines, by broadening and hence desaturating them. It iscaused by small cells of motion in the photosphere and is treated like an additionalthermal velocity in the line absorption coefficient. Weak lines are not affected, asthey are gaussian, and so broadening them also makes them shallower and henceEW is preserved. Typical values are 1-2 km/s. It can be determined by ensuring thatfor individual elements EW is independent of line.


Measurement of log10 g

• Balmer jump can be used in A and F stars to measure Pe

• A comparison of lines formed by neutral atoms to the lines of ions alsogives a measure of electron pressure

• The lines strong enough to show wings (including H) often have the wingstrength dependent on the pressure through (e.g.) Stark broadening.

• In all cases there is a temperature dependence - either the temperaturemust be securely established before executing a pressure analysis or asimulataneous pressure-temperature solution must be made.

A precise spectroscopic measurement of photospheric pressure is difficultto achieve as there are no spectral features that are strikingly sensitive topressure as the lines and continuum are to temperature.


Balmer Lines

Balmer lines, like Hγ, are pressure indicators for Teff > 7500K


The different methodsThere are several methods (isochrones, pressure broadening in the wings ofstrong lines) to estimate the gravity of a star (log10 g).

Photometry:We can use photometry to estimate the surface gravity of a star if we knowthe mass and the Bolometric magnitude and the Teff.

Spectroscopy:Measure elements in two ionization states (e.g., Fe I & Fe II). By definition,gravity is related to the gas pressure (Pg g2/3) and the electronic pressure(Pe g1/3). From Saha’s equation in the cool stars case, where the numberof atoms of Fe I FeII; Fe I, the dominant species, will depend on 1/Pe andFe II, the minority species, with the majority of atoms in the state i 1 = 1, will depend on 1/Pe

2. This is what makes the ionization equilibrium agood tool to constrain gravity.


MetallicityElemental abundance properties contribute to the continuous absorptionproperties of the atmosphere.


What is a stellar model - what comes out


Stellar ModelsChose which models to use, or create


Scaling RelationsWhen we have weak lines dominated by Doppler broadening

Changes in log A are equivalent to changes in log gfλ, or θxχ or log κν

wavelength integral of the atomicabsorption coefficient, α. Constantfor a given star & ion

Partition fn, used to calculateexcitation & ionisation

A, the number abundance forelement E relative to H, A = NE / NH

Excitation potential for the level(Saha)

Normalizes Dopplerdependent phenomena

Integral over theline profile

Transition probability opacity

The equation tells us that for a give star, curves of growth for lines of the samespecies, where A is constant, will differ only in displacements along the abscissaaccording to their individual values of log gfλ, or θxχ or log κν. So we can vary A togenerate the curve of growth for this line.


Curve of Growth


Line ListA proper line list is a critical part of the analysis, and building one needs some care

Lines need to be chosen carefully, making sure they have reliable gf-values and aresufficiently isolated from their neighbours at the resolution of the observations and ofcourse to lie within the wavelength coverage of the instrument.


Looking for crowding

Letarte 2007, PhD thesis

Spectral Synthesis


Part 3b: Some Results


David Malin UKST

60’

7.107M

MV=-13.2140kpc

The Fornax dwarf spheroidal:


Alpha Elements


Heavy Elements: s-process


Heavy Elements: r-process


Part 3c: What do we learn from Abundances?


We seeks to account for the production of the chemicalelements that we see in the Universe, its time dependenceand for many of the features of galaxies that we observe.Understanding stellar evolution, the birth and death of starsand how they interact with their environments is central tounderstanding the evolution of galaxies.


Light Elements – e.g., O Na Mg Al tracers of deep mixing abundances patterns (globular clusters versus field stars) α- Elements – e.g., O Mg Si Ca Ti dominated by products of Supernovae II Iron-peak Elements e.g., V Cr Mn Co Ni Cu Zn explosive nucleosynthesis (supernovae I)

Heavy Elements ( Z > 30 ) mix of r- and s- process elements e.g., s-process e.g., Ba, La (stellar winds) r-process e.g., Eu

e.g., McWilliam 1997

Chemical Tagging


Cluster 1Letarte et al. 2006

Letarte et al. 2006

VLT/UVES spectroscopy

Deep Mixing Pattern

Fornax Globular Clusters: deep mixing


Origin of the r-process

It has been argused that the r-process production occurs predominantly in low massSNe II (Matthews et al. 1992). The high [Eu/Fe] in Fornax suggests an importantcontrubution of low mass stars to the enrichment of Fornax. This is compatible withlow [α/Fe], for which we also need low mass SNII. However. Plotting [Eu/α] it is clearthat the sites and relative contribution of α- and r- process elements differ in Fornaxand the Milky Way. If r- and a- elements were created in the same way the ratio of[Eu/α] should be constant.

Lecture 3: Detailed Stellar Abundances

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