In 1835 the French positivist philosopher Auguste Comte cited the chemical constitution of the stars as an example of knowledge that might be forever hidden. However, unknown to Comte, the development of spectroscopy was already revealing the composition of the stars and permitting the emergence of a true astrophysics. In 1802 English physician William Hyde Wollaston saw several dark gaps or lines in the Sun’s spectrum and conjectured that these might be the natural boundaries between colours. The dark lines in the solar spectrum were rediscovered around 1814 in Munich by Fraunhofer, who cataloged some 500 of them. Fraunhofer noted that his dark D line in the yellow part of the solar spectrum matched up with the well-known bright line in the spectrum of a candle flame. Fraunhofer also showed that light from Venus shows the same structure as sunlight, and he observed dark lines in the spectra of a number of bright stars.

A key step was taken in 1849 by French physicist Jean Foucault, who showed that the bright orange lines seen in the light emitted by a carbon arc could also be observed as dark absorption lines in sunlight that was passed through the gas around the arc. Thus, a gas that can be stimulated to emit a particular colour will also preferentially absorb that same colour. Around 1859 German chemist Robert Wilhelm Bunsen and physicist Gustav Robert Kirchhoff showed how to associate spectral lines with particular chemical elements. From an analysis of the dark lines in the solar spectrum, Kirchhoff concluded that iron, calcium, magnesium, sodium, nickel, and chromium were present in the Sun. In 1868 English astronomer Joseph Norman Lockyer identified an orange line in a solar-prominence spectrum that had no counterpart in that of any known element, so he ascribed it to a new element, which he called helium (after helios, the Greek name for the Sun and the Sun god). Helium was not isolated on Earth until 1895 by Scottish chemist William Ramsay.

In the 1860s Italian astrophysicist Angelo Secchi described the spectra of some 4,000 stars and classified them into four groups. A star’s spectrum is continuous, with all the colours present, though it may be brighter in one or another part of the spectrum according to the temperature of the star. (Cooler stars are redder.) Typically, the continuous spectrum is also overlaid with a number of dark absorption lines. Secchi’s classification scheme was based on the overall colour of the star, the number and kind of absorption lines, and other features of the spectrum. This work, performed before the application of photography to spectroscopy, was slow and very tedious.

Also in the 1860s English astronomer William Huggins observed the spectrum of a bright nebula and found that it consisted only of bright emission lines. This was therefore a glowing gas—a case of true nebulosity. Huggins went on to observe about 70 nebulae. He found that the nebulae consisted of two major groups. About one-third were gaseous, and about two-thirds showed the continuous spectrum that would be expected of unresolved stars.

A major centre of spectroscopy in the next generation was the Harvard College Observatory, under the direction of American astronomer Edward Charles Pickering. By putting a prism in front of the object lens of a telescope, his team was able to photograph the spectra of many stars at once. The resulting Henry Draper Catalogue (named to recognize the financial support for the project provided by Draper’s widow) appeared in nine volumes between 1918 and 1924 and contained over 225,000 spectra. Key to this work was a new stellar-classification scheme (still in use today—for example, the Sun is a G-type star) refined by American astronomer Annie Jump Cannon, who had joined Pickering’s team in 1895.

In the mid-19th century there was considerable dispute about the reality and nature of the Doppler effect. A shift in the frequency of light received from a moving source had been proposed in 1842 by the Austrian physicist Christian Doppler, who (wrongly) thought that in this way he could explain the colours of binary stars. The Doppler effect was demonstrated for sound by the Dutch physicist Christophorus Henricus Didericus Buys-Ballot in 1845 by putting musicians on a moving train. In 1868 Huggins measured a small shift in the position of the F line in the hydrogen spectrum for Sirius, which was interpreted as being caused by the radial motion of the star with respect to Earth. Strong confirmation of the Doppler effect for light was obtained in the 1870s by German astronomer Hermann Karl Vogel, who measured the spectral shift between the east and west edges of the rotating Sun. In the 1880s Vogel and German astronomer Julius Scheiner began to measure the radial velocities of stars by using photographic spectra. The tabulation of spectral types and radial velocities soon became a standard part of star cataloging.

The cataloging of stellar spectra opened the way for new discoveries, for it soon became clear that the spectral type of a star has a relation to the star’s intrinsic brightness. However, since a star will look dimmer the farther away it is, the intrinsic brightness (or absolute magnitude) of a star cannot be known unless one first has a way to determine the distance. American astronomer Henry Norris Russell in 1913 published a scatter plot correlating absolute magnitude with spectral type, using only stars for which he judged that the distances had been well determined. Slightly earlier, German astronomer Hans Rosenberg and Danish astronomer Ejnar Hertzsprung had plotted similar diagrams, using only stars from a single cluster, either the Pleiades or the Hyades. (Stars in a single cluster are all at roughly the same distance from Earth, so their apparent magnitudes can be used as replacements for their absolute magnitudes.) The resulting scatter plots are called Hertzsprung-Russell (H-R) diagrams. The H-R diagram revealed that most stars lie on a “main sequence,” in which absolute magnitude is positively correlated with temperature. Bluer main-sequence stars (spectral type O or B) are much brighter than main-sequence red stars (spectral type K or M). The H-R diagram also showed a second branch, in which there are reddish stars that are much brighter than those on the main sequence. If these bright red stars have the same surface temperature (because they are of the same spectral type) as a main-sequence star but are much brighter, they must be physically larger, and they soon came to be called “red giants.” White dwarfs were soon discovered as yet another branch. The H-R diagram became crucial for guiding speculations about the evolution of stars.

The source of the energy that drives the stars had been a great mystery. In the 19th century, chemical combustion and heating due to gravitational contraction were the only possibilities, but Scottish physicist William Thomson (Lord Kelvin) pointed out that a chemical process could hardly last more than 3,000 years. In various versions of heating by release of gravitational energy, the Sun was supposed to be contracting slowly (by about 75 metres [246 feet] per year) or else be heated by the continual infall of meteoric matter. After the discovery of radioactivity in the 1890s and the realization that Earth’s interior was warmed by this mechanism, various schemes were proposed for explaining stellar energy in terms of radioactive decay. The true explanation came only after German American physicist Albert Einstein’s 1905 publication of the mass-energy relation (E = mc², a consequence of special relativity). In the 1920s English astrophysicist Arthur Eddington proposed the proton-proton reaction, in which four atoms of hydrogen are combined to produce one atom of helium, with the mass difference released in the form of energy. Because of the primitive state of nuclear physics at the time, he could not say in detail how this might occur, but he pointed to the mere existence of helium in the stars as the surest proof that such a process must exist. Nuclear physics gained a firm foundation in the early 1930s with the discovery of the neutron and of deuterium (a heavy isotope of hydrogen with a proton and a neutron in its nucleus). From then on, progress was rapid. In 1937 German physicist Carl Friedrich von Weizsäcker discovered the CNO cycle, in which carbon, nitrogen, and oxygen act as catalysts in a sequence of nuclear reactions that leads to the conversion of hydrogen into helium. In 1939 German American physicist Hans Bethe published a more detailed and quantitative study of the CNO cycle that finally put stellar astrophysics on a secure footing. Bethe also treated in detail the proton-proton reaction that Eddington had only guessed at. In a collision at high temperature, two protons may stay close enough together for the brief time required for one of them to be converted into a neutron by emission of a positron; thus, deuterium is formed. From deuterium, helium may then be built up in several different ways. Bethe also showed that the CNO cycle is more important in high-temperature stars and the proton-proton reaction more important in cooler stars. Nuclear physics was successfully integrated with what was known about the conditions of temperature and density in the interiors of stars.