Abundances and Kinematics from HighResolution Spectroscopic Surveys
Lecture 3:
Detailed Stellar Abundances
Eline Tolstoy
Kapteyn Astronomical Institute,
University of Groningen
I have a spectrum, and colours...
130 candidates
D=79 kpc
rcore=5.8'
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Part 3a: Measuring Abundances
Measurements of equivalent widths can be used to derive the abundance of
elements for which we can detect and measure absorption lines in the gas
making up the stellar atmosphere.
The Boltzmann and Saha equations are applied and combined with the
pressure and temperature of the gas to derive an abundance of the element.
Letarte 2007, PhD thesis
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Obtaining information from spectra
Chemical analysis usually proceeds using a curve-of-growth technique either
explicitly or implicitly. The alternative is to synthesize the whole region.
Literature: Alot of the detailed work presented here comes from the Phd
of Bruno Letarte, University of Groningen, 2007 - you can ask for a copy
if you are interested in all the details mentioned here, and somethings
from yesterday:
[email protected]
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
No. of absorbing atoms per unit area
To find the number of absorbing atoms per unit area, Na, that have electrons in the
proper orbital to absorb a photon at the wavelength of the spectral line - the T, p are
used 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 states
are equally likely. Each transition has a relative probably, or f-value (also called
oscillator strength).
These can be calculated theoretically or measured in a lab, and they are defined so
that the f-values for transitions from the same orbital add up to the number of
electrons in the atom or ion, i.e., the effective number of electrons per atom
participating in a transition.
Multiplying the number of absorbing atoms per unit area by the f-value gives the
number of atoms lying above each cm2 of photosphere.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Voigt Profiles of K line of Ca II
For varying numbers of absorbing Ca ions
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Curve of Growth
This an important tool to determine Na, the number of absorbing atoms, and thus
the abundances of elements in stellar atmospheres, because EW varies with Na.It
is a log-log plot of EW as a function of the number of absorbing atoms.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
So, practically - Low Mass Stars (RGB):
What we need to start:
1. Spectral type
2. Atmospheric properties (line formation)
Effective Temperature
Surface Gravity
Metallicity
Then:
1. Interpolate published model atmosphere grids (or
make own)
2. Compute theoretical (synthetic) spectrum for assumed
abundances
3. Compare with observations and iterate on
abundances
TO GET AN ABUNDANCE
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Stellar Parameters
Effective Temperature, Teff
Temperature of Black-body with the same L and R as the real star
Surface Gravity, log10g
Metallicity, Z, or [Fe/H]
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
How to determine Teff
Photometry: An empirical method, calibrated on the InfraRed Flux Method,
(Blackwell & Lynas-Gray 1998), gives a relation between photometric
colours (like V 䌝㻃I, V 䌝㻃K) and Teff. The general method and correction
polynomials 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 is
independent of the excitation potential (χex) of the individual lines. to use
this method, we need many lines of a single element sampling a range of
χex. The precision with which Teff can be determined depends upon the
resolution, the choice and number of lines and signal to noise of each
spectrum used.
Teff +400K
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Microturbulence Velocity, vt
Can be important in strong lines, by broadening and hence desaturating them. It is
caused by small cells of motion in the photosphere and is treated like an additional
thermal velocity in the line absorption coefficient. Weak lines are not affected, as
they are gaussian, and so broadening them also makes them shallower and hence
EW is preserved. Typical values are 1-2 km/s. It can be determined by ensuring that
for individual elements EW is independent of line.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Measurement of log10 g
A precise spectroscopic measurement of photospheric pressure is difficult
to achieve as there are no spectral features that are strikingly sensitive to
pressure as the lines and continuum are to temperature.
• 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 also
gives a measure of electron pressure
• The lines strong enough to show wings (including H) often have the wing
strength dependent on the pressure through (e.g.) Stark broadening.
• In all cases there is a temperature dependence - either the temperature
must be securely established before executing a pressure analysis or a
simulataneous pressure-temperature solution must be made.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Balmer Lines
Balmer lines, like Hγ, are pressure indicators for Teff > 7500K
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
The different methods
There are several methods (isochrones, pressure broadening in the wings of
strong 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 know
the 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 䌨㻃g2/3) and the electronic pressure
(Pe 䌨㻃g1/3). From Saha’s equation in the cool stars case, where the number
of atoms of Fe I FeII; Fe I, the dominant species, will depend on 1/Pe and
Fe II, the minority species, with the majority of atoms in the state
i䌝
1 = 1, will depend on 1/Pe2. This is what makes the ionization equilibrium a
good tool to constrain gravity.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Metallicity
Elemental abundance properties contribute to the continuous absorption
properties of the atmosphere.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
What is a stellar model - what comes out
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Stellar Models
Chose which models to use, or create
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Scaling Relations
When we have weak lines dominated by Doppler broadening
Transition probability
opacity
wavelength integral of the atomic
Excitation potential for the level
Normalizes Doppler
(Saha)
dependent phenomena absorption coefficient, α. Constant
for a given star & ion
A, the number abundance for
Integral over the
element E relative to H, A = NE / NH
line profile
Partition fn, used to calculate
excitation & ionisation
Changes in log A are equivalent to changes in log gfλ, or θxχ or log κν
The equation tells us that for a give star, curves of growth for lines of the same
species, where A is constant, will differ only in displacements along the abscissa
according to their individual values of log gfλ, or θxχ or log κν. So we can vary A to
generate the curve of growth for this line.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Curve of Growth
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Line List
A 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 are
sufficiently isolated from their neighbours at the resolution of the observations and of
course to lie within the wavelength coverage of the instrument.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Looking for crowding
Spectral Synthesis
Letarte 2007, PhD thesis
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Part 3b: Some Results
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
The Fornax dwarf spheroidal:
7.107M
MV=-13.2
140kpc
60’
David Malin UKST
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Alpha Elements
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Heavy Elements: s-process
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Heavy Elements: r-process
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Part 3c: What do we learn from Abundances?
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
We seeks to account for the production of the chemical
elements that we see in the Universe, its time dependence
and for many of the features of galaxies that we observe.
Understanding stellar evolution, the birth and death of stars
and how they interact with their environments is central to
understanding the evolution of galaxies.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Chemical Tagging
 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
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Fornax Globular Clusters: deep mixing
Letarte et al. 2006
Letarte et al. 2006
Cluster 1
Deep Mixing Pattern
VLT/UVES spectroscopy
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007
Origin of the r-process
It has been argused that the r-process production occurs predominantly in low mass
SNe II (Matthews et al. 1992). The high [Eu/Fe] in Fornax suggests an important
contrubution of low mass stars to the enrichment of Fornax. This is compatible with
low [α/Fe], for which we also need low mass SNII. However. Plotting [Eu/α] it is clear
that the sites and relative contribution of α- and r- process elements differ in Fornax
and the Milky Way. If r- and a- elements were created in the same way the ratio of
[Eu/α] should be constant.
Eline Tolstoy
Third Chilean Advanced School of Astrophysics
Concepción, 8-12 January 2007