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Part V
Stellar spectroscopy
53
Chapter 10
Classification of stellar spectra
Goal-of-the-Day
To classify a sample of stars using a number of temperature-sensitive spectral lines.
10.1
The concept of spectral classification
Early in the 19th century, the German physicist Joseph von Fraunhofer observed the solar
spectrum and realised that there was a clear pattern of absorption lines superimposed on
the continuum. By the end of that century, astronomers were able to examine the spectra
of stars in large numbers and realised that stars could be divided into groups according
to the general appearance of their spectra. Classification schemes were developed that
grouped together stars depending on the prominence of particular spectral lines: hydrogen
lines, helium lines and lines of some metallic ions. Astronomers at Harvard Observatory
further developed and refined these early classification schemes and spectral types were
defined to reflect a smooth change in the strength of representative spectral lines. The
order of the spectral classes became O, B, A, F, G, K, and M; even though these letter
designations no longer have specific meaning the names are still in use today. Each
spectral class is divided into ten sub-classes, so that for instance a B0 star follows an O9
star.
This classification scheme was based simply on the appearance of the spectra and
the physical reason underlying these properties was not understood until the 1930s. Even
though there are some genuine differences in chemical composition between stars, the main
property that determines the observed spectrum of a star is its effective temperature. Thus
spectral classification provides a measurement for the surface temperature of a star. For
instance, O-type stars are the hottest in the classification scheme. They exhibit strong
lines of ionised helium because at very high atmospheric temperatures (∼ 40 000 K) such
ions are present in large enough numbers to produce a detectable absorption line. At the
other temperature extreme, M-type stars show absorption bands due to molecules that
are not present in stars with hotter atmospheres. This is because at higher temperatures
these molecules become dissociated.
The spectral classification scheme in use today is a refinement of these earlier efforts.
Firstly it was necessary to account for the fact that stars have very different radii at
different stages in their evolution (i.e. their surface gravity is very different), with larger
stars being more luminous. The luminosity together with the temperature completely
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CHAPTER 10. CLASSIFICATION OF STELLAR SPECTRA
define a star’s position in the HR diagram. Certain spectral lines are not only temperaturesensitive but also gravity-sensitive. Such lines are used to further classify stars into
luminosity classes (using roman numerals): supergiants (type I), giants (III), and mainsequence stars (V). The Sun for instance, a typical main-sequence star, is classified as
G2V. In the lab session we will limit ourselves to classifying main-sequence stars, thus we
are not concerned with the luminosity class.
More recently, as astronomers became able to observe fainter stars, it was realised that
the classical spectral classification scheme did not reach low enough temperatures. Two
more spectral types were introduced, L and T, and astronomers are currently trying to
identify the first elusive Y-type stars, the coolest stars expected to exist.
And then there are various classes of peculiar stars, that show differences in their
spectra with respect to those of normal stars of similar temperature and luminosity. The
chemically most-exotic of these have spectral types of their own: notably the cool carbon
stars, and the hot Wolf-Rayet stars that lack hydrogen and only have spectral lines in
emission.
The table below lists the spectral types and most important properties (in terms of
spectral lines) that are still routinely used to classify new stars.
Type
O
B
A
F
G
K
M
Colour
Blue
Blue
Blue
Blue/White
White/Yellow
Orange/Red
Red
Temperature (K)
> 30 000
10 000–30 000
7 500–10 000
6 000–7 500
5 000–6 000
3 500–5 000
< 3 500
Main characteristics
singly-ionised helium lines (He ii)
neutral helium lines (He i)
hydrogen lines (H i); maximum strength A0 stars
metallic lines
strong lines of metallic atoms and ions (e.g., Ca i)
metallic lines dominate
molecular bands of titanium oxide (TiO)
Exercise 10.1
(a) The light we receive from a star in a strong spectral line was emitted higher up in
the star’s atmosphere than the light emitted in the spectral continuum. What, therefore,
can you say about the temperature structure of the star’s atmosphere if the spectral line
is (i) in absorption, or (ii) in emission?
10.2
Measuring the strength of spectral lines
The strength of a spectral feature is often measured by its equivalent width, Wλ :
Z
Wλ =
Fλc − Fλ
dλ,
Fλc
(10.1)
where Fλc is the level of the continuum emission at wavelength λ. The unit of equivalent
width is Ångström (Å), where 1 Å = 0.1 nm. The equivalent width is a measure of the
area of the spectral line on a plot of intensity versus wavelength. It is found by forming
a rectangle with a height equal to that of the continuum emission, and finding the width
such that the area of the rectangle is equal to the area in the spectral line.
10.2. MEASURING THE STRENGTH OF SPECTRAL LINES
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Exercise 10.2
(a) Consider two stars of identical spectral type. One of the stars is twice as bright as the
other. How do the values of the equivalent widths of the spectral lines compare between
these two stars?
Iraf offers a set of tools to analyse spectra — these can be found in the onedspec package
which itself is found in the noao package. The task we shall use here is called splot. If
you type splot [spectrum] (where [spectrum] is the name for the FITS file with the
spectrum you want to measure), the spectrum will be displayed in a separate graphics
window (click on it to make it active).
To read the cursor position, press the spacebar. To zoom in, place the cursor first at
one side of the region of interest, and press a. Then, place the cursor at the other side of
the region of interest, and press a again. To zoom out, press c. To measure the equivalent
width of a spectral line, place the cursor first on the point where the spectral line is first
seen to go into absorption, and press e. Then, place the cursor on the point where the
spectrum has returned to the level of the continuum emission, and press e again. To quit
splot, press q.
Exercise 10.2
(b) Create and change to a new directory called project4, and make Iraf — do not forget
to update the login.cl file!
(c) Retrieve the data from ftp.astro.keele.ac.uk, directory /pub/astrolab/project4;
(d) Measure the wavelength and equivalent width of a spectral line of your choice in one
of the spectra.
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CHAPTER 10. CLASSIFICATION OF STELLAR SPECTRA
10.3
Classifying a sample of spectra
The diagram below shows how the equivalent widths of several spectral lines vary amongst
stars of different spectral type (remember that the x-axis is a temperature scale):
Exercise 10.3
(a) Fill in the table below with your measurements of the equivalent widths of as many
of these spectral lines as you can identify;
(b) By comparison with the above diagram, determine the spectral types of the stars.
Star Ca ii K
Fe i
Ca i
Hγ
He i
He ii Spectral
3933 Å 4045 Å 4227 Å 4340 Å 4471 Å 4541 Å Type
1
2
3
4
5
6
Exercise 10.3
(c) What could be the reason for the presence of the Ca ii line in the spectrum of the
hottest star in the sample?