The light from an H+ region is of a quite different character from that of a star. Whereas a normal star emits a continuous spectrum with dark absorption lines, an H+ region emits almost all its light in a few comparatively narrow emission lines .The most important of these emission lines are those of hydrogen, helium and oxygen. The visible hydrogen lines form the so-called ‘Balmer Series’. and arise as electrons in the upper states of a hydrogen atom drop to the second lowest state. Helium lines are only seen from ‘high-excitation’ nebulae – in other words those with very hot exciting stars; the ionization potential for helium is much larger than that for hydrogen (24.6 electron volts instead of 13.6), so that only stars which have a large flux shortward of 50.4 nm can ionize helium in the nebula.

The hydrogen and helium lines are sometimes called RE¬COMBINATION LINES, since they occur as an indirect result of the recombination of an electron with a hydrogen or helium ion. When the recombination takes place the resultant neutral atom is prob¬ably in an excited rather than in the ground state; it then cascades down to lower and lower energy states, emitting a photon at each transition, until it reaches the ground state, where it will remain until it is ionized again by an ultraviolet photon from the exciting star. In a nebula with density 109 ions m-3 and radius 1 pc a hydrogen atom would typically remain ionized for some 100 years before recombining. It would then cascade to the ground state in a fraction of a second, and then remain there for something like two months before being ionized again.

In most nebulae the strongest spectral lines are not those of hydrogen or helium, but those of the oxygen ions O++at 495.9 and 500.7nm, and 0+ at 372.6 and 372.9nm. The prominence of these lines is remarkable not only because oxygen ions are a thousand times rarer than hydrogen ions, but also because these particular spectral lines are extremely difficult to produce in a terrestrial laboratory. The atomic transitions that give rise to these lines are called FORBIDDEN TRANSITIONS and are normally too weak to be seen. Under the extreme low-density conditions found in interstellar clouds, however, these lines come to dominate the spectrum. The oxygen ions, in fact, act as a thermostat that keeps all H+ regions at much the same temperature (about 10 000K) by absorbing kinetic energy from the fast-moving electrons in the nebula and radiating the energy away as emission lines. The physics of why these forbidden lines are so strong is complicated, and for many years after the discovery of these lines in nebular spectra, their origin was a mystery. For a while there was speculation that the lines were due to a hypothetical chemical element NEBULIUM, and it was only in 1927 that the lines were correctly identified. Since then, forbidden lines have been seen from other elements such as nitrogen, neon and sulphur and are much used to measure the temperature and density of H+ regions and of some other astronomical objects including planetary nebulae and quasars.

Emission lines are also seen at infrared and radio wavelengths, superimposed on the strong continuum. Spectra at these wave¬lengths are particularly useful for the study of distant and of obscured H+ regions because of the fact that infrared and radio radiation is much less absorbed by dust than is light. At infrared wavelengths, the brightest lines are certain forbidden lines, such as the 12.8 urn line of Ne+. At radio wavelengths, it is the hydrogen and helium recombination lines that are the most important; for example, the transition from the 110th to the 109th energy level of hydrogen occurs at a radio frequency of 5.009 GHz and has been much studied. These radio recombination lines are particularly useful in that they allow the radial velocities of H+ regions to be measured, even if they are invisible from the Earth. For many obscured H+ regions the radial velocity is the only clue we have to its distance, through a model of the rotation of the Galaxy.

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