Interstellar

The mass of gas in interstellar space is a hundred times larger than that of dust. Although there is more gas, it is less easily observed, since the gas does not cause a general extinction of light. In the optical region it can only be observed on the basis of a small number of spectral lines.

The existence of interstellar gas began to be suspected in the first decade of the 20th century, when in 1904 Johannes Hartmann observed that some absorption lines in the spectra of certain binary stars were not Doppler shifted by the motions of the stars like the other lines. It was concluded that these absorption lines were formed in gas clouds in the space between the Earth and the stars. In some stars there were several lines, apparently formed in clouds moving with different velocities. The strongest lines in the visible region are those of neu tral sodium and singly ionized calcium (Fig. 15.14). In the ultraviolet region, the lines are more numerous. The strongest one is the hydrogen Lyman a line (121.6 nm).

On the basis of the optical and ultraviolet lines, it has been found that many atoms are ionized in interstellar space. This ionization is mainly due to ultraviolet radiation from stars and, to some extent, to ionization by cosmic rays. Since the density of interstellar matter is very low, the free electrons only rarely encounter ions, and the gas remains ionized.

About thirty elements have been discovered by absorption line observations in the visible and ultraviolet region. With a few exceptions, all elements from hydrogen to zinc (atomic number 30) and a few additional heavier elements have been detected (Table 15.3). Like in the stars, most of the mass is hydrogen (about 70%) and helium (almost 30%). On the other hand, heavy elements are significantly less abundant than in the Sun and other population I stars. It is thought that they have been incorporated into dust grains, where they do not produce any absorption lines. The element abundances in the interstellar medium (gas + dust) would then be normal, although the interstellar gas is depleted in heavy

Observable phenomenon

Cause

Interstellar extinction and polarization

Non-spherical dust grains aligned by magnetic field

Dark nebulae, uneven distribution of stars and galaxies

Dust clouds

Interstellar absorption lines in stellar spectra

Atoms and molecules in the interstellar gas

Reflection nebulae

Interstellar dust clouds illuminated by nearby stars

Emission nebulae or HII regions (optical, infrared and radio emission)

Interstellar gas and dust cloud, where a nearby hot star ionizes the gas and heats the dust to 50-100 K

Optical galactic background (diffuse galactic light)

Interstellar dust illuminated by the integrated light of all stars

Galactic background radiation:

b) long wavelength (> 1 m)

Free-free emission from hot interstellar gas Synchrotron radiation from cosmic ray electrons in the magnetic field

Galactic 21 cm emission

Cold (100 K) interstellar neutral hydrogen clouds (HI regions)

Molecular line emission (extended)

Giant molecular clouds (masses even

105-106 Mg), dark nebulae

Point-like OH, H2O and SiO sources

Maser sources near protostars and long-period variables

Table 15.2. Phenomena caused by the interstellar medium

Table 15.2. Phenomena caused by the interstellar medium

Fig. 15.14. (a) The D lines Di and D2 of interstellar sodium (rest wavelengths 589.89 and 589.00 nm) in the spectrum of the star HD 14134. Both lines consist of two components formed in the gas clouds of two spiral arms. The radial velocity difference of the arms is about 30km/s. (Mt. Wilson Observatory). (b) The interstellar absorption lines of ionized calcium Ca II and ionized methyli-dyne CH+ in the spectra of several stars. The emission spectrum of iron is shown for comparison in (a) and (b). (Lick Observatory)

Table 15.3. Element abundances in the interstellar medium towards Z Ophiuchi and in the Sun. The abundances are given relative to that of hydrogen, which has been defined to be

Table 15.3. Element abundances in the interstellar medium towards Z Ophiuchi and in the Sun. The abundances are given relative to that of hydrogen, which has been defined to be

1,000,000. An asterisk (*) means that the abundance has been determined from meteorites. The last column gives the ratio of the abundances in the interstellar medium and in the Sun

Atomic

Name

Chemical

Interstellar

Solar

Abundance

number

symbol

abundance

abundance

ratio

1

Hydrogen

H

1,000,000

1,000,000

1 .00

2

Helium

He

85,000

85,000

ra 1

3

Lithium

Li

0.000051

0.00158*

0.034

4

Beryllium

Be

< 0.000070

0.000012

< 5.8

5

Boron

B

0.000074

0.0046*

0.016

6

Carbon

C

74

370

0.20

7

Nitrogen

N

21

110

0.19

8

Oxygen

O

172

660

0.26

9

Fluorine

F

-

0.040

-

10

Neon

Ne

-

83

-

11

Sodium

Na

0.22

1.7

0.13

12

Magnesium

Mg

1.05

35

0.030

13

Aluminium

Al

0.0013

2.5

0.00052

14

Silicon

Si

0.81

35

0.023

15

Phosphorus

P

0.021

0.27

0.079

16

Sulfur

S

8.2

16

0.51

17

Chlorine

Cl

0.099

0.45

0.22

18

Argon

Ar

0.86

4.5

0.19

19

Potassium

K

0.010

0.11

0.094

20

Calcium

Ca

0.00046

2.1

0.00022

21

Scandium

Sc

-

0.0017

-

22

Titanium

Ti

0.00018

0.055

0.0032

23

Vanadium

V

< 0.0032

0.013

< 0.25

24

Chromium

Cr

< 0.002

0.50

< 0.004

25

Manganese

Mn

0.014

0.26

0.055

26

Iron

Fe

0.28

25

0.011

27

Cobalt

Co

< 0.19

0.032

< 5.8

28

Nickel

Ni

0.0065

1.3

0.0050

29

Copper

Cu

0.00064

0.028

0.023

30

Zinc

Zn

0.014

0.026

0.53

elements. This interpretation is supported by the observation that in regions where the amount of dust is smaller than usual, the element abundances in the gas are closer to normal.

Atomic Hydrogen. Ultraviolet observations have provided an excellent way of studying interstellar neutral hydrogen. The strongest interstellar absorption line, as has already been mentioned, is the hydrogen Lyman a line (Fig. 15.15). This line corresponds to the transition of the electron in the hydrogen atom from a state with principal quantum number n = 1 to one with n = 2. The conditions in interstellar space are such that almost all hydrogen atoms are in the ground state with n = 1. Therefore the Lyman a line is a strong absorption line, whereas the Balmer absorption lines, which arise from the excited initial state n = 2, are unobservable. (The

1,000,000. An asterisk (*) means that the abundance has been determined from meteorites. The last column gives the ratio of the abundances in the interstellar medium and in the Sun

Balmer lines are strong in stellar atmospheres with temperatures of about 10,000 K, where a large number of atoms are in the first excited state.)

The first observations of the interstellar Lyman a line were made from a rocket already in 1967. More extensive observations comprising 95 stars were obtained by the OAO 2 satellite. The distances of the observed stars are between 100 and 1000 parsecs.

Comparison of the Lyman a observations with observations of the 21 cm neutral hydrogen line have been especially useful. The distribution of neutral hydrogen over the whole sky has been mapped by means of the 21 cm line. However, the distances to nearby hydrogen clouds are difficult to determine from these observations. In the Lyman a observations one usually knows the distance to the star in front of which the absorbing clouds must lie.

Fig. 15.15. Interstellar absorption lines in the ultraviolet spectrum of Z Ophiuchi. The strongest line is the hydrogen Lyman a line (equivalent width, more than 1 nm). The observa-

tions were made with the Copernicus satellite. (Morton, D.C. (1975): Astrophys. J. 197, 85)

Fig. 15.15. Interstellar absorption lines in the ultraviolet spectrum of Z Ophiuchi. The strongest line is the hydrogen Lyman a line (equivalent width, more than 1 nm). The observa-

tions were made with the Copernicus satellite. (Morton, D.C. (1975): Astrophys. J. 197, 85)

The average gas density within about 1 kpc of the Sun derived from the Lyman a observations is 0.7 atoms/cm3. Because the interstellar Lyman a line is so strong, it can be observed even in the spectra of very nearby stars. For example, it has been detected by the Copernicus satellite in the spectrum of Arcturus, whose distance is only 11 parsecs. The deduced density of neutral hydrogen between the Sun and Arcturus is 0.02-0.1 atoms/cm3. Thus the Sun is situated in a clearing in the interstellar medium, where the density is less than one tenth of the average density.

If a hydrogen atom in its ground state absorbs radiation with a wavelength smaller than 91.2 nm, it will be ionized. Knowing the density of neutral hydrogen, one can calculate the expected distance a 91.2 nm photon can propagate before being absorbed in the ionization of a hydrogen atom. Even in the close neighbourhood of the Sun, where the density is exceptionally low, the mean free path of a 91 . 2 nm photon is only about a par-sec and that of a 10 nm photon a few hundred parsecs. Thus only the closest neighbourhood of the Sun can be studied in the extreme ultraviolet (XUV) spectral region.

The Hydrogen 21 cm Line. The spins of the electron and proton in the neutral hydrogen atom in the ground state may be either parallel or opposite. The energy difference between these two states corresponds to the frequency of 1420.4 MHz. Thus transitions between these two hyperfine structure energy levels will give rise to a spectral line at the wavelength of 21.049 cm (Fig. 5.8). The existence of the line was theoretically predicted by Hendrick van de Hulst in 1944, and was first observed by Harold Ewen and Edward Purcell in 1951. Studies of this line have revealed more about the properties of the interstellar medium than any other

V

I.I.I

. I.I.

- ^

1 \ 1 1 1

1 . j . 1 . 1 . 1 . 1

F t

1 1 1 ' 1 1 1 1 1 ' ' i= 1°

^ _^

, 1 ! .

-120 -100 -80 -60 -40 -20 0 +20 +4(3 Radial velocity v [km/s]

-120 -100 -80 -60 -40 -20 0 +20 +4(3 Radial velocity v [km/s]

Fig. 15.16. Hydrogen 21 cm emission line profiles in the galactic plane at longitude 180°, 90° and 1° (in the direction l = 0° there is strong absorption). The horizontal axis gives the radial velocity according to the Doppler formula, the vertical axis gives the brightness temperature. (Burton, W. B. (1974): "The Large Scale Distribution of Neutral Hydrogen in the Galaxy", in Galactic and Extra-Galactic Radio Astronomy, ed. by Ver-schuur, G.L., Kellermann, K.I. (Springer, Berlin, Heidelberg, New York) p. 91)

method - one might even speak of a special branch of 21 cm astronomy. The spiral structure and rotation of the Milky Way and other galaxies can also be studied by means of the 21 cm line.

Usually the hydrogen 21 cm line occurs in emission. Because of the large abundance of hydrogen, it can be observed in all directions in the sky. Some observed 21 cm line profiles are shown in Fig. 15.16. Rather than frequency or wavelength, the radial velocity calculated from the Doppler formula is plotted on the horizontal axis. This is because the broadening of the 21 cm spectral line is always due to gas motions either within the cloud (turbulence) or of the cloud as a whole. The vertical axis is mostly plotted in terms of the antenna temperature TA (see Chap. 5), the usual radio astronomical measure of intensity. The brightness temperature of an extended source is then Tb = TA/nB, where nB is the beam efficiency of the antenna.

For the 21cm line hv/k = 0.07 K, and thus hv/kT ^ 1 for all relevant temperatures. One may therefore use the Rayleigh-Jeans approximation (5.24)

2v2kT

In the solution of the equation of radiative transfer (5.42) the intensity can thus be directly related to a corresponding temperature. By definition, Iv is related to the brightness temperature Tb, and the source function Sv is related to the excitation temperature Texc, i. e.

In certain directions in the Milky Way there is so much hydrogen along the line of sight that the 21 cm line is optically thick, tv ^ 1. In that case

Fig. 15.17. The distribution of neutral hydrogen in the galaxy from the Leiden and Parkes surveys. The density is given in atoms/cm3. (Oort, J.H., Kerr, P.T., West-erhout, G.L. (1958): Mon. Not. R. Astron. Soc. 118, 379)

Fig. 15.17. The distribution of neutral hydrogen in the galaxy from the Leiden and Parkes surveys. The density is given in atoms/cm3. (Oort, J.H., Kerr, P.T., West-erhout, G.L. (1958): Mon. Not. R. Astron. Soc. 118, 379)

i. e. the brightness temperature immediately yields the excitation temperature of the cloud. This is often referred to as the spin temperature TS.

The excitation temperature need not always agree with the kinetic temperature of the gas. However, in the present case the population numbers of the hyperfine levels are determined by mutual collisions of hydrogen atoms: the time between collisions is 400 years on the average, whereas the time for spontaneous radiative transitions is 11 million years; thus the excitation temperature will be the same as the kinetic temperature. The observed temperature is T « 125 K.

The distance to a source cannot be obtained directly from the observed emission. Thus one can only study the number of hydrogen atoms in a cylinder with a 1 cm2 base area extending from the observer to outside the Milky Way along the line of sight. This is called the projected or column density and is denoted by N. One may also consider the column density N(v) dv of atoms with velocities in the interval [v,v + dv].

It can be shown that if the gas is optically thin, the brightness temperature in a spectral line is directly proportional to the column density N of atoms with the corresponding radial velocity. Hence, if the diameter L of a cloud along the line of sight is known, the gas density can be determined from the observed line profile:

The diameter L can be obtained from the apparent diameter, if the distance and shape of the cloud are assumed known.

The distances of clouds can be determined from their radial velocities by making use of the rotation of the Milky Way (Sect. 17.3). Thus if the observed peaks in the 21 cm line profiles (Fig. 15.16) are due to individual clouds, their distances and densities can be obtained. Since radio observations are not affected by extinction, it has been possible in this way to map the density distribution of neutral hydrogen in the whole galactic plane. The resulting distribution, based on observations at Leiden and Parkes, is shown in Fig. 15.17. It appears that the Milky Way is a spiral galaxy and that the interstellar hydrogen is concentrated in the spiral arms. The average density of interstellar hydrogen is 1 atom/cm3, but the distribution is very inhomogeneous. Typically the hydrogen forms denser regions, a few parsecs in

size, where the densities may be 10-100atoms/cm3. Regions where the hydrogen is predominantly neutral are known as H I regions (in contrast to H II regions of ionized hydrogen).

The hydrogen 21 cm line may also occur in absorption, when the light from a bright radio source, e. g. a quasar, passes through an intervening cloud. The same cloud may give rise to both an absorption and an emission spectrum. In that case the temperature, optical thickness and hydrogen content of the cloud can all be derived.

Like interstellar dust hydrogen is concentrated in a thin disc in the galactic plane. The thickness of the hydrogen layer is about twice that of the dust or about 200 pc.

HII Regions. In many parts of space hydrogen does not occur as neutral atoms, but is ionized. This is true in particular around hot O stars, which radiate strongly in the ultraviolet. If there is enough hydrogen around such a star, it will be visible as an emission nebula of ionized hydrogen. Such nebulae are known as HII region (Figs. 15.18 and 15.19).

A typical emission nebula is the great nebula in Orion, M42. It is visible even to the unaided eye, and is a beautiful sight when seen through a telescope. In the middle of the nebula there is a group of four hot stars known as the Trapezium, which can be distinguished inside the bright nebula, even with a small telescope. The Trapezium stars emit strong ultraviolet radiation, which keeps the gas nebula ionized.

Unlike a star a cloud of ionized gas has a spectrum dominated by a few narrow emission lines. The continuous spectrum of H II regions is weak. In the visible region the hydrogen Balmer emission lines are particularly strong. These are formed when a hydrogen atom recombines into an excited state and subsequently returns to the ground state via a sequence of radiative transitions. Typically a hydrogen atom in a HII region remains ionized for several hundred years. Upon recombination it stays neutral for some months, before being ionized again by a photon from a nearby star.

The number of recombinations per unit time and volume is proportional to the product of the densities of electrons and ions, nrec a neni. (15.21)

< Fig. 15.18. The great nebula in Orion (M42, NGC 1976). The nebula gets its energy from newly formed hot stars. The dark regions are opaque dust clouds in front of the nebula. Radio and infrared observations have revealed a rich molecular cloud behind the nebula (Fig. 15.20). In the upper part of this picture is the gas nebula NGC 1977, in the lower part the bright star i Orionis. (Lick Observatory)

In completely ionized hydrogen, ne = ni, and hence nrec a n2. (15.22)

Most recombinations will include the transition n = 3 ^ 2, i. e. will lead to the emission of a Ha photon.

Thus the surface brightness of a nebula in the Ha line will be proportional to the emission measure,

where the integral is along the line of sight through the nebula.

The ionization of a helium atom requires more energy than that of a hydrogen atom, and thus regions of ionized helium are formed only around the hottest stars. In these cases, a large H II region will surround a smaller central

Fig. 15.19. The Lagoon nebula (M8, NGC 6523) in Sagittarius. This H II region contains many stars of early spectral types and stars that are still contracting towards the main sequence. Small, round dark nebulae, globules, are also vis-

ible against the bright background. These are presumably gas clouds in the process of condensation into stars. (National Optical Astronomy Observatories, Kitt Peak National Observatory)

He+ or He++ region. The helium lines will then be strong in the spectrum of the nebula.

Although hydrogen and helium are the main constituents of clouds, their emission lines are not always strongest in the spectrum. At the beginning of this century it was even suggested that some strong unidentified lines in the spectra of nebulae were due to the new element nebulium. However, in 1927 Ira S. Bowen showed that they were forbidden lines of ionized oxygen and nitrogen, O+, O++ and N+. Forbidden lines are extremely difficult to observe in the laboratory, because their transition probabilities are so small that at laboratory densities the ions are de-excited by collisions before they have had time to radiate. In the extremely diffuse interstellar gas, collisions are much less frequent, and thus there is a chance that an excited ion will make the transition to a lower state by emitting a photon.

Because of interstellar extinction, only the nearest HII regions can be studied in visible light. At infrared and radio wavelengths much more distant regions can be studied. The most important lines at radio wavelengths are recombination lines of hydrogen and helium; thus the hydrogen transition between energy levels 110 and 109 at 5.01 GHz has been much studied. These lines are also important because with their help radial velocities, and hence (using the galactic rotation law), distances of H II regions can be determined, just as for neutral hydrogen.

The physical properties of H II regions can also be studied by means of their continuum radio emission. The radiation is due to bremsstrahlung or free-free emission from the electrons. The intensity of the radiation is proportional to the emission measure EM defined in (15.23). HII regions also have a strong infrared continuum emission. This is thermal radiation from dust inside the nebula.

HII regions are formed when a hot O or B star begins to ionize its surrounding gas. The ionization steadily propagates away from the star. Because neutral hydrogen absorbs ultraviolet radiation so efficiently, the boundary between the H II region and the neutral gas is very sharp. In a homogeneous medium the H II region around a single star will be spherical, forming a Strom-gren sphere. For a B0 V star the radius of the Stromgren sphere is 50 pc and for an A0 V star only 1 pc.

The temperature of a H II region is higher than that of the surrounding gas, and it therefore tends to ex pand. After millions of years, it will have become extremely diffuse and will eventually merge with the general interstellar medium.

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