Most of the gas in the Galactic disk is in an unionised, or neutral state, and is termed H I. Under the right conditions, such as near a hot star or after a supernova explosion, the gas can become ionised, and these places are called H II regions. The gas in both states is more generally called the interstellar medium or ISM.
H II regions are hot ( ~ 104 K) and sit in a background of ``cold'' ( ~ 102 K) H I. H II regions stand out very clearly on direct photographs (or CCD images) of external galaxies. This is because they are dominated by emission lines of various species, such as hydrogen itself, but also helium, oxygen, carbon, calcium, silicon, iron etc. Images taken in the light of Ha at l = 6563 Å (i.e. the a-line of the Balmer series) show H II regions dramatically, because much of the light in the visible region is emitted in this line. As a consequence H II regions typically appear pinkish on colour images.
H II regions are typical in spiral galaxies, in particular in the spiral arms, where young stars provide copious ultra-violet photons which are needed to ionise the gas. The most prominent are star forming regions, in which there are plenty of freshly born, hot stars. Some of these regions can contain so many energetic stars that they can ionise a significant fraction of the gas in a galaxy.
Star forming regions are prominantly associated with spiral arms in disk galaxies, and because of their clear emission lines, are an excellent way to trace the kinematics and rotational properties of the galaxies.
Elliptical and S0 galaxies rarely show prominant H II regions. These galaxies are mostly old, and have usually already used up most of their gas. They do show planetary nebulae however, and although they are rather faint, they are an important distance indicator.
H II regions have a certain size around an ionising source which is a balance between the flux of ionising photons on one hand, and the rate at which the plasma can cool and the ionised electrons can re-combine with the protons. If the ionisation rate is greater than the recombination rate, then the plasma will become almost fully ionised. Studies indicate that less than a per cent of the hydrogen remains unionised in typical regions.
Ultra-violet photons below a wavelength of l = 912 Å, have sufficient energy to ionise hydrogen, removing the electron from the ground shell. The electron is eventually recaptured, but because interstellar gas is of such low density, this can take some time.
Hot stars are the best sources of the UV photons needed to ionise the gas. In practice this means very early type stars, such as O and B stars, and white dwarfs. The surfaces of the O and B stars range in temperature from 10,000 to 60,000 K, while white dwarfs can reach surface temperature of up to 200,000 K.
Recombination leads to emission lines, as shown in figure 9.2. Important ones are Ha at 6563 Å, [N II] at 6583 Å, [O II] at 3726 and 3729 Å and [O III] at 4959 and 5007 Å.
A special notation has been used here, the square brackets, ``[ ]''. This indicates that the line is normally ``forbidden'', meaning it is only seen in very low density conditions. In the laboratory, plasmas are generally of such high density that these line are not seen. The lines originate from energy states just above the ground state, and are meta-stable with long lifetime.
The emission lines form very useful diagnostics of the physical conditions in the the gas. For example, the [O III] and [N III] lines are temperature sensitive, while [O II] and [S II] (at 6716 and 6731 Å) are more sensitive to the electron density. This is because the latter lines are emitted at different levels but with nearly the same excitation energy, so that theie relative ratio is a diagnostic of the collisional de-excitation, or the density, of the gas. The plasmas are found to be very thin by terrestrial standards, with particle densities of order 10 to 106 cm-3.
Observations at radio wavelengths are very interesting, because these regions typically emit bremsstrahlung radiation at these frequencies, giving an excellent independent probe of the temperatures and electron densities in the clouds.
Measurements of the line strengths also give information about the ratios of elemental abundances, so that the composition of the ISM can be studied.
The shapes of the lines also give information about the kinematic conditions in the nebula, since they are sensitive to the velocity of the ions, either through thermal broadening or bulk motions.
A higher resolution spectrum of H II regions in a spiral galaxy is shown in figure 9.4. The H II regions on either side of the center are redshifted and blueshifted relative to the center of the galaxy.
Complexes of young stars emerging from freshly collapsed gas characterise star forming regions. A magnificent example is the Orion nebula, which is visible as a fuzzy patch to the naked eye. Figure 9.5 shows the central region of this nebula, a region which in the optical is almost deviod of stars, but which when observed in the infra-red, has recently been shown to contain many sources behind a shielding screen of dust. Dust is very common in star forming regions.
Fresh stars are born with a range of masses, with low mass stars ( ~ 0.5 MO) dominating. The small number of very massive stars (M > MO) have a great effect on the development of the region, since they provide the bulk of the ionising UV photons, and can heat the surrounding gas to a few ×104 K. Just a few massive stars can clear an impressive region around themselves, exposing the rest of the newly formed, and suppressing further star formation. The massive stars are also short-lived, burning through their fuel at a fierce rate and eventually ending their lives as supernovae, which injects enormous amounts of energy into the surrounding gas and heating it to much greater temperatures (106) K.
Star forming regions range in size from order 10-3 pc to several 100 pc. Particle density ranges from > 106 cm-3 in the compact regions down to 10 cm-3 in a giant H II clouds. Three spiral galaxies are shown in Fig 9.1 in broadband light and the light of Ha, showing how closely spiral arms and other kinematical features in the disk are associated with star forming regions.
Planetary nebulae (PN) are small regions of ionised gas around a very hot central white dwarf, and are formed during the final stages of stellar evolution. They have a very wide range of shapes and properties, as can be seem from the splendid Space Telescope images in figure 9.7. This type of nebula was discovered by Herschell in 1785, who referred to them as ``planetary'' because they appear disk-like in telescopes (like planets). About 2000 PN are known in the galaxy, and they are also rather easily recognised in the closer external galaxies.
Temperatures, densities and compositions for PN can be derived fairly easily. The central star can always be seen, and its temperature measured, allowing the amount of UV flux it produces to be estimated. This UV flux then heats the nebula, and various emission lines used to derive electron densities and temperatures, and finally the abundances of the elements.
Most PN have abundances like stars in the Galactic disk, as expected if they have evolved from disk stars. PN of the Galactic halo are much more metal poor, also as one would expect. There can be interesting differences however. Helium, carbon and oxygen can be strongly enhanced relative to PN parent stars; this is because these elements, which are created in the parent population during nuclear burning, but which are hidden below the surface, can get mixed into the outer layers of the star when it is on the giant branch and be expelled into the surrounding gas, which is later lit up by the central star when it becomes a white dwarf.
Ionised masses in PN are typically a few tenths of a solar mass.
Lots more images can be found at http://www.noao.edu/image_gallery/planetary_nebulae.html
The name given to an ionised region is Strömgren Sphere, after the Danish astronomer Bengt Strömgren. Ionising photons moving outward from a central source can ionise a finite region around it, set by the ionisation rate and recombination rate. For a UV flux S*, the radius R of the region is given by
where np is the number density of protons in the gas surrounding the source, and b ~ 2.6 ×10-13 cm3 s-1 is the Hydrogen recombination coefficient. There is a thin layer around the ionised region, called an ionisation front, which forms the transition zone from hot ionised gas to the cold surrounding gas. The thickness of this zone is approximately equal to the mean free path of photons in the front.