Measuring Spectra

The most important methods of forming a spectrum are by means of an objective prism or a slit spectrograph. In the former case one obtains a photograph, where each stellar image has been spread into a spectrum. Up to several hundred spectra can be photographed on a single plate and used for spectral classification. The amount of detail that can be seen in a spectrum depends on its dispersion, the range of wavelengths per millimetre on the

Fig. 8.1a-g. Typical stellar spectra. The spectrum of n Pegasi (f) is very similar to that of the Sun. The continuous spectrum is brightest at about 550 nm and gets fainter towards smaller and larger wavelengths. Dark absorption lines are superimposed on the continuum. See also Exercise 8.1. (Mt. Wilson Observatory)

Fig. 8.1a-g. Typical stellar spectra. The spectrum of n Pegasi (f) is very similar to that of the Sun. The continuous spectrum is brightest at about 550 nm and gets fainter towards smaller and larger wavelengths. Dark absorption lines are superimposed on the continuum. See also Exercise 8.1. (Mt. Wilson Observatory)

Hannu Karttunen et al. (Eds.), Stellar Spectra.

In: Hannu Karttunen et al. (Eds.), Fundamental Astronomy, 5th Edition. pp. 207-219 (2007) DOI: 11685739_8 © Springer-Verlag Berlin Heidelberg 2007

plate (or per pixel on a CCD). The dispersion of an objective prism is a few tens of nanometres per millimetre. More detailed observations require a slit spectrograph, which can reach a dispersion 1-0.01 nm/mm. The detailed shape of individual spectral lines can then be studied.

The photograph of the spectrum is converted to an intensity tracing showing the flux density as a function of wavelength. This is done by means of a microden-sitometer, measuring the amount of light transmitted by the recorded spectrum. Since the blackening of a photographic plate is not linearly related to the amount of energy it has received, the measured blackening has to be calibrated by comparison with known ex posures. In modern CCD spectrographs the intensity curve is determined directly without the intervening step of a photographic plate. For measurements of line strengths the spectrum is usually rectified by dividing by the continuum intensity.

Figure 8.2 shows a photograph of the spectrum of a star and the intensity curve obtained from a calibrated and rectified microdensitometer tracing. The second pair of pictures shows the intensity curve before and after the normalisation. The absorption lines appear as troughs of various sizes in the curve. In addition to the clear and deep lines, there are large numbers of weaker lines that can barely be discerned. The graini-ness of the photographic emulsion is a source of noise

b) 588 590 592 594 596 598 600 c) 588 590 592 594 596 598 600

588 590 592 594 596 598 600 588 590 592 594 596 598 600

Wavelength [nm] Wavelength |nmj

Fig. 8.2. (a) A section of a photograph of a stellar spectrum picture the intensity curve of the first picture has been rec-

and the corresponding rectified microdensitometer intensity tified by normalizing the value of the continuum intensity tracing. The original spectrum was taken at the Crimean Ob- to one. (Pictures by J. Kyrolainen and H. Virtanen, Helsinki servatory. (b) A more extensive part of the spectrum. (c) The Observatory)

588 590 592 594 596 598 600 588 590 592 594 596 598 600

Wavelength [nm] Wavelength |nmj

Fig. 8.2. (a) A section of a photograph of a stellar spectrum picture the intensity curve of the first picture has been rec-

and the corresponding rectified microdensitometer intensity tified by normalizing the value of the continuum intensity tracing. The original spectrum was taken at the Crimean Ob- to one. (Pictures by J. Kyrolainen and H. Virtanen, Helsinki servatory. (b) A more extensive part of the spectrum. (c) The Observatory)

which appears as irregular fluctuations of the intensity curve. Some lines are so close together that they appear blended at this dispersion.

The detailed shape of a spectral line is called the line profile (Sect. 5.3). The true shape of the line reflects the properties of the stellar atmosphere, but the observed profile is also spread out by the measuring instrument. However, the total absorption in the line, usually expressed in terms of the equivalent width, is less sensitive to observational effects (see Fig. 5.6).

The equivalent width of a spectral line depends on how many atoms in the atmosphere are in a state in which they can absorb the wavelength in question. The more atoms there are, the stronger and broader the spectral line is. For example, a typical equivalent width of a metal line (Fe) in the solar spectrum is about 10 pm. Line widths are often expressed in angstroms (1 A = 10-10m = 0.1 nm).

Only in weak lines the equivalent width depends linearly on the number of absorbing atoms. The equivalent width as a function of the amount of absorbing atoms is known as the curve of growth. It is, however, beyond the scope of this book.

Line profiles are also broadened by the Doppler effect. In stellar atmospheres there are motions of small and large scale, like thermal motion of the atoms and convective flows.

The chemical composition of the atmosphere can be determined from the strengths of the spectral lines. With the introduction of large computers it has become feasible to construct quite detailed models of the structure of stellar atmospheres, and to compute the emergent spectrum for a given model. The computed synthetic spectrum can be compared with the observations and the theoretical model modified until a good fit is obtained. The theoretical models then give the number of absorbing atoms, and hence the element abundances, in the atmosphere. The construction of model atmospheres will be discussed in Sect. 8.6.

Telescopes Mastery

Telescopes Mastery

Through this ebook, you are going to learn what you will need to know all about the telescopes that can provide a fun and rewarding hobby for you and your family!

Get My Free Ebook


Post a comment