Spectral Classification of Stars
Spectral Classification of Stars
Although the composition of most stars is very similar, there are systematic variations in stellar spectra based on their temperatures. A typical star has a spectrum consisting of a continuous range of colors overlaid with dark lines. The positions, strengths, and shapes of these lines are determined by the temperature, density, gravitational fields, velocity, and other properties of the star. In order to be able to study stars systematically, it is useful to classify stars with others that have similar properties. This is the basis for the classification scheme used by astronomers. Stars are classified according to the patterns and relative strengths of their dark spectral lines, which are indicators of both their temperature and their intrinsic luminosity, or brightness. Although roughly 10% of stars do not fit into the classification scheme, it provides a convenient way to understand the systematics of stellar formation and evolution.
When the light from a star is divided into its component colors using a spectrograph, it appears as a continuous band of colors, broken up by dark, narrow lines. These lines are created by atoms and ions (atoms missing [positive ion] or acquiring [negative ion] one or more electrons) in the outer layers of a star’s atmosphere. These layers absorb light at specific wavelengths, which are unique for each type of atom or ion. Atomic physics predicts the positions and intensities of these lines, called absorption lines, based on the temperature and composition of the star. Thus, the number, strengths, and positions of these lines vary from star to star.
The first stellar spectra were observed in 1814, long before the atomic physics that creates them was understood. In an attempt to understand the processes that formed the spectra, similar stars with similar spectra were grouped together in the hopes that stars that were alike would produce similar spectra. In 1863, Italian astronomer Father Angelo Secchi (1818–1878) made one of the first attempts at trying to classify stars, when he divided stars into two groups based on their spectral lines. He eventually extended this categorization, dividing more than 4,000 stars into four classes.
The basis of the current system of classification of spectral types began in the late 1800s at the Harvard College Observatory, under the direction of Professor Edward C. Pickering. Scottish-American astronomer Williamina Paton Stevens Fleming (1857–1911) initially classified 10,000 stars using the letters of the alphabet to denote the strength of their hydrogen absorption lines, with A being the strongest, followed by B, C, etc. At the time, she did not know that these lines were due to hydrogen, but since they were visible in almost all stellar spectra, they provided a convenient means by which to organize her data.
Several years later, the classifications were reordered to be in what scientists now know to be the order of decreasing temperature: O, B, A, F, G, K, M, in order to have a smooth transition between the class boundaries. This reordering was done primarily by American astronomer Annie Jump Cannon (1863– 1941), also at the Harvard Observatory, in preparing the Henry Draper catalog of 225,000 stars. She also further subdivided each class into as many as ten subclasses, by adding the numbers 0 through 9 after the
letter, to account for changes within a class. This spectral classification scheme was formally adopted by the International Astronomical Union in 1922, and is still used today.
It was not until 1925 that the theoretical basis behind the ordering was discovered. At first, scientists believed that the strength of the lines directly determined the amount of each element found in the star, but the situation proved more complex than that. Most stars have very similar compositions, so the strength of hydrogen (and other) lines in the spectrum is not a measure of the makeup of the star. Instead, it is a measure of temperature, as a result of the atomic physics processes occurring in the star. At relatively low temperatures, the gas in a stellar atmosphere contains many atoms, (and even some molecules), and these produce the strongest absorption lines. At higher temperatures, molecules are destroyed and atoms begin losing electrons, and absorption lines of ions begin to appear. More and more ionization occurs as the temperature increases, further altering the pattern of absorption lines. Thus, the smooth sequence of line patterns is actually a temperature sequence.
Another of the early classification workers at Harvard, American astronomer Antonia Caetana De Paiva Pereira Maury (1866–1952), noted that certain dark lines (absorption lines) in stellar spectra varied in width. She attempted a classification based partially on line widths, but this was not adopted by Annie Cannon in her classification, and was not used in the Henry Draper catalog. However, Maury’s work laid the foundation for the subsequent discovery that the line widths were related to stellar size: very large stars, now called giants or supergiants, have thin lines due to their low atmospheric pressure. These stars are very luminous because they have large surface areas, and so line width was eventually recognized as an indicator of stellar luminosity.
In 1938, American astronomer William Wilson Morgan (1906–1994) at the Yerkes Observatory added a second dimension to the classification scheme, by using the luminosity of the star as an additional classifying feature. He used roman numerals to represent the various types of stars. In 1943, the MKK Atlas of Stellar Spectra (after Morgan, P.C. Keenan, and E. Kellman) was published, formalizing this system. Approximately 90% of stars can be classified using the MKK system.
The temperature range of each class, along with the most prominent spectral lines that form the basis of the spectral identification, are described below. A common mnemonic for remembering the order of the spectral classes is Oh Be A Fine Girl (or Guy) Kiss Me. In the original scheme, there were a few additional classes (S, R, C, N), which turned out to represent stars that actually do have abnormal compositions. Today, these stars are usually in transient evolutionary phases, and are not included in the standard spectral classifications.
O (30,000 to 60,000 Kelvin (K), blue-white)—At such high temperatures, most of the hydrogen is ionized, and thus the hydrogen lines are less prominent than in the B and A classes (ionized hydrogen with no remaining electron has no spectral lines). Much of the helium is also ionized. Lines from ionized carbon, nitrogen, oxygen, and silicon are also seen.
B (10,000 to 30,000 K, blue white)—In stars in this spectral class, the hydrogen lines are stronger than in O stars, while the lines of ionized helium are weaker. Ionized carbon, oxygen, and silicon are seen.
A (7500 to 10,000 K, blue white)—A stars have the strongest hydrogen lines (recall the ordering of the original Harvard classification). Other prominent lines are due to singly ionized magnesium, silicon, and calcium.
F (6000 to 7500 K, yellow-white)—Lines from ionized calcium are prominent features in F stars.
G (5000 to 6000 K, yellow)—The ionized calcium lines are strongest in G stars. The sun is a G2 star.
K (3500 to 5000 K, orange)—The spectra of K stars contain many lines from neutral elements.
M (less than 3500 K, red)—Molecular lines seen in the spectra of M stars mean that the temperature is low enough that molecules have not been broken up into
Ionized —Missing or acquiring one or more electrons, resulting in a charged atom.
Spectrograph —Instrument for dispersing light into its spectrum of wavelengths then photographing that spectrum.
Spectrum —A display of the intensity of radiation versus wavelength.
their constituent atoms. Titanium oxide (TiO) is particularly prominent.
The MKK luminosity classes are: I-Supergiants; II-Bright Giants; III-Normal Giants; IV-Subgiants; V-Main Sequence.
See also Spectroscopy.
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