R

Wavelength

Fig 14.10.

On top is the arrangement for doing optical or uv absorption spectroscopy on an interstellar cloud. Below is a schematic absorption spectrum.

These simple molecules are not generally stable in the laboratory. CH+ is charged and would combine with a negative ion or electron under laboratory conditions. CH and CN have an outer electronic shell with only one electron (as does H), making them chemically reactive. The presence of these unstable molecules in the interstellar gas suggests densities much lower than in the typical laboratory.

In these early studies no hydrogen absorption lines were observed. This is not because there was no hydrogen present. The temperatures in interstellar clouds are generally low, and most of the hydrogen is in the ground state. Therefore, the only H absorption lines possible are the Lyman lines in the ultraviolet. Now that ultraviolet observations are possible from satellites, astronomers can study these absorption lines.

You might wonder how we know that the absorption lines are coming from the interstellar gas and not from the stars themselves. After all, we have already seen the large number of absorption lines present in stars. One distinguishing feature is that the interstellar lines are much narrower than the stellar absorption lines. By narrower we mean they cover a smaller range of frequency (or wavelength). Interstellar lines have Doppler broadenings that correspond to a few kilometers per second. If the Doppler broadening is produced by random thermal motions, this suggests a temperature of about 100 K. Also, systematic studies show that, on the average, the absorption lines are stronger when detected in the light of more distant stars. The more distant the star is, the more interstellar material there is between us and the star. The narrow interstellar lines do not appear in the spectra of all stars. This suggests that the interstellar gas is clumpy, just as the interstellar dust is clumpy.

14.4.2 Radio studies of atomic hydrogen

Much of what we know about the interstellar medium comes from radio observations. We have already seen that supernova remnants, planetary nebulae and pulsars are sources of radio emission. These are generally hot sources, or sources with high energy electrons that produce a high radio luminosity. However, most of the interstellar gas is cool and does not produce a strong radio continuum emission. The cool interstellar gas must be observed via radio spectral lines.

The first interstellar radio line to be observed was from atomic hydrogen, but it was not a transition in which an electron jumps from one orbit to another. As we have said, these transitions are in the visible and ultraviolet parts of the spectrum. For the radio transition, the hydrogen stays in the ground electronic state. This is illustrated in Fig. 14.11. Both the electron and proton have intrinsic angular momentum, called spin. We have already seen that this spin can have two

Spins Opposite

Spins Aligned

Spins Opposite

Spins Aligned

(Expanded by factor of 105 )

Fig 14.11.

(a) Origin of the 21 cm line.The splitting comes from a magnetic interaction which depends on the spin directions of the electron and the proton.The energy is higher when the spins are parallel and lower when they are antiparallel. (b) Energy level diagram showing the splitting of the hydrogen ground state (n = 1).The splitting is greatly exaggerated in this figure.

possible orientations. We refer to them as 'up' and 'down'. This means that the electron and proton spins can be either parallel or antiparallel. The relative orientation of the spins affects the magnetic force between the electron and the proton. The state with the spins parallel has slightly more energy than the state with the spins antiparallel. The atom can undergo transitions between these two states. The energy difference corresponds to a frequency of about 1400 MHz, or a wavelength of 21 cm. This is generally referred to as the 21 cm line.

If we take the energy of the transition hv and divide by Boltzmann's constant, the quantity hv/k gives the temperature necessary to see collisional excitation of the hydrogen upper state. This is about 0.07 K. This means that even at the low temperatures of interstellar space there will be sufficient energy to excite transitions between these two states in hydrogen. The 21 cm line is easily observable under interstellar conditions. The possibility of detecting this line was discussed in Leiden (Netherlands) in the early 1940s by Henk van de Hulst. After that there was a race among Australian, Dutch and American groups to detect the line. The first detection of the 21 cm line from interstellar hydrogen was in the early 1950s by a group at Harvard, led by Edwin Purcell, who won the Nobel Prize in physics for his work.

Since that time there have been extensive observations of the 21 cm line by radio astronomers all over the world. It is probably fair to say that it was the dominant tool for studying the interstellar medium and galactic structure through the 1960s, and continues to be very useful. In these studies, the line was observed in both emission and absorption. The conditions for emission or absorption lines are shown in Fig. 14.12. In order for the line to be in absorption, there must be a background continuum source whose brightness temperature at 21 cm is greater than the excitation temperature of the atoms in the particular cloud being observed. Under most conditions the excitation temperature of the 21 cm line is close to the kinetic temperature of the clouds.

By studying both absorption and emission lines in a given region it is possible to deduce

Continuum Source

Continuum Source

Cloud

Fig 14.12.

Conditions for radio absorption and emission lines. If a radio continuum source is viewed through an interstellar cloud, then radio absorption lines can be seen against the continuum source. (This also requires the continuum source to appear hotter than the cloud at the wavelength of observation.) If there is no background continuum source, then only emission lines can be seen.

both the excitation temperature and the optical depth. The excitation temperature enables us to calculate the kinetic temperature of the gas, an important quantity. The optical depth can be converted into a column density for atomic hydrogen. If we know the column density and size of a cloud, we can also find the average local density of hydrogen. So you can see that the 21 cm line observations provide astronomers with an important tool for studying the physical conditions in many interstellar clouds.

One important feature of the radio observations is that interstellar dust is transparent at radio wavelengths. Therefore we can use radio telescopes to detect objects across the galaxy, far beyond what we can see optically in the presence of dust. Since we can use it to observe clouds anywhere in the galaxy, the 21 cm line is a very useful tool for studying galactic structure. Also, since it is a spectral line, we can observe its Doppler shift and learn about motions throughout our galaxy. We will see how these studies are used in Chapter 16.

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