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 55 56 CHAPTER 10. CLASSIFICATION OF STELLAR SPECTRA define a star’s position in the HR diagram. Certain spectral lines are not only temperature- sensitive but also gravity-sensitive. Such lines are used to further classify stars into luminosity classes (using roman numerals): supergiants (type I), giants (III), and main- sequence 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 Colour Temperature (K) Main characteristics O Blue > 30 000 singly-ionised helium lines (He ii) B Blue 10 000–30 000 neutral helium lines (He i) A Blue 7 500–10 000 hydrogen lines (H i); maximum strength A0 stars F Blue/White 6 000–7 500 metallic lines G White/Yellow 5 000–6 000 strong lines of metallic atoms and ions (e.g., Ca i) K Orange/Red 3 500–5 000 metallic lines dominate M Red < 3 500 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: c Fλ Fλ Wλ = −c dλ, (10.1) Z Fλ c where Fλ is the level of the continuum emission at wavelength λ. The unit of equivalent width is Angstr¨om˚ (A),˚ where 1 A˚ = 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 57 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. 58 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 A˚ 4045 A˚ 4227 A˚ 4340 A˚ 4471 A˚ 4541 A˚ 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?.