Download Lecture 10: Interstellar gas

Lecture 10: Interstellar gas
Interstellar Medium (ISM)
• In spite of the space between the stars though to be emptier than the best
vacuums created on the Earth, there is some material between the stars
composed of gas and dust. This material is called the interstellar
medium. The interstellar medium makes up between 10 to 15% of the
visible mass of the Milky Way. About 99% of the material is gas and the
rest is ``dust''. The interstellar medium affects starlight and stars (and
planets) form from clouds in the interstellar medium, so it is worthy of
study. Also, the structure of the Galaxy is mapped from measurements of
the gas.
• ISM is not uniform:
- Regions vary significantly in size, temperature, and density of matter.
- Highly rarified by Earth standards
- The densest portions of the ISM are the birth places of stars and
planetary systems in our galaxy.
- Most phases not seen in optical except for T = 104 K gas, heated and
ionized by OB stars
Interstellar gas
• The interstellar gas produces its own characteristics emission and absorption line spectra.
The temperature and density of the gas determine these characteristic spectral features.
• In general, the gas is transparent over wide range despite the fact that the total mass of the
gas in our galaxy is greater than the total mass of dust by a factor of about 100.
Interstellar optical absorption lines
• Some stars have in their spectra absorption lines that are quite out of character with the
spectral class.
• Many B stars shows sharp multiple Ca II lines.
• Some spectroscopic binaries show particular spectral lines that remain fixed in
wavelength while the rest of the spectral lines shift periodically (figure 15.7 A).
• Clearly, these lines originate in the interstellar medium. Multiple lines arise when there
are several absorbing clouds along the line-of-sight (Figure 15-7B). Optical absorption
lines, identified as interstellar in origin, include those from Ca I, Ca II, Ti I, Ti II, Na I
and the molecules CN and CH.
• The intensity of s line depends on the amount of gas lying between the star and the
observer. If the gas is distributed uniformly through space, the intensities of IS absorption
lines depend directly on the path length traversed by the starlight.
• Low gas density plays a role in preventing ions from recombining into neutral atoms
after photoionization. Sufficiently energetic photons and cosmic rays will occasionally
encounter and ionize the widespread gas atoms and molecules. In order to recombine,
an ion must capture an electron, but at typical IS densities the chance of such a capture
is very small.
Emission Nebulae: H II regions
I) Hydrogen line emission
• Emission nebula is a hot cloud of gas (mainly
hydrogen) whose visible spectrum dominated by
emission lines.
• Hot O & B stars emit a huge amount s of UV
radiation ; such energetic photons, with wavelengths
less than 91.2 nm, ionize any hydrogen atom they
encounter. If such a hot star is surrounded by a cloud
of gas, the hydrogen atoms close to the star will be
ionized and form an HII region.
• Away from the star, there are no sufficient photons
to ionize the hydrogen and H II regions sharply
terminates (neutral hydrogen H I prevails).
Emission lines in a typical gaseous nebula. The strongest
lines are Hα in red, [OIII] in green and [OII] in ultraviolet
Spectrum of NGC7009, a
planetary nebula, but similar
to a typical diffuse gaseous
nebula spectrum.
Diagram of spectrum
of the Orion nebula
Chemical composition of HII nebulae
element
H
He
C
N
O
log10N
12.0
11.0
8.5
8.0
8.8
All other elements have log10 N < 8.0
Physical processes in HII regions
H + hν (λ < 912 nm) → p + e (photoionization)
p + e → H* + hν
(recombination)
H* → H + hν
(cascading)
O++ +
(O++)*
e
→ (O++)* + e (collisional excitation)
→ O++ + hν (radiative deexcitation)
Typical radius and mass of HII regions
Spectral type of star
O5
B0
A0
radius of nebula (pc)
70–200
20
0.5
They can only readily be observed around stars of types
O to B0 (T* ~ 50 000 K to 25 000 K)
Mass: 0.1 to 103 M⊙
Evolution of HII regions
HII regions are surrounded by HI gas, but being much hotter,
they are high pressure regions which therefore expand.
The expansion is supersonic, and creates shock waves in the
surrounding HI gas. Usually hot stars disappear in a few ×
106 years, before pressure equilibrium can be achieved, and
so the HII region also dies out, reverting to HI condition.
Star formation
and glowing HII
regions in the
Great Orion
Nebula
An OB association is where O and B class stars are producing
ionizing radiation which makes an HII nebula glow.
Some famous HII nebulae
Orion nebula
M42
η Carinae nebula
30 Doradus (in LMC)
Lagoon nebula
M8
Rosette nebula
Trifid nebula
M20
NGC1976
NGC3372
NGC2070
NGC6523
NGC2237
NGC6514
Below: Lagoon
nebula M8 in Sagittarius
Above: Trifid
nebula, M20, in
Sagittarius
Right: Rosette nebula in Monoceros
Below, Tarantula nebula, 30 Doradus
in the Large Magellanic Cloud
Above: Orion
nebula, M42
Right: η Carinae nebula, in
southern Milky Way
Hubble Space Telescope
images of the Orion nebula
Right: detail of centre
Eagle nebula M17
• The hydrogen gas in IS space is extremely dilute and cold. Half of the gas is H I (neutral
hydrogen) in the ground state because collisional excitation is rare.
• Imagine a hot star with temperature greater or equal to 20,000 K, which produce ample
UV radiation. If the gas density is uniform , the UV radiation from the central star
ionizes all the hydrogen in a roughly spherical volume of space (Stromgren sphere).
Equilibrium is established when the rate of recombination (H II + e → H I) equals the
rate of photoionization. The H II region is maintained by the continual re-ionization of
recombined H I atoms due to the flux of UV photons from the central star.
• In an idealized case, recombinations will balance ionizations, the total number of
ionizing photons per second Nuv will equal the total recombination per second:
Nuv = (4π/3) Rs3 ne nH α(2)
where α(2) is the recombination coefficient (m3/s) of H excluding the n=1 state. Such
captures produce another ionizing photon; captures to n =2 or higher produce photons
longward of the Lyman limit. These quickly escape the H II region. So the Stromgren
radius is given by
Rs3 =
[Nuv /(4π/3) ne nH α(2)]
• At greater distance from the star, the inverse-square law diminishes the flux of UV
photons, and ionization of the recombined H I atoms is no longer possible. So the ratio of
H I to HII rises sharply with increasing distance from the star. In addition, most of the H
II recombines to an excited state of the neutral hydrogen ; the atom then quickly cascades
to the ground state (Balmer lines), emitting several low energy photons that escape from
H II region. So H II region fluoresces by converting the stellar UV radiation to lower
energy photons (visible light).
• Optical fluorescence lines of helium are also strong in the spectra of emission nebulae;
together with the radio recombination lines of helium (arising from transitions between
high excitation levels), these lines permit us to:
1- Study the excitation mechanisms operating in H II regions.
2- Investigate the elemental abundances (especially He/H) of the ISM.
3- Probe the spiral structure of our Galaxy.
• Radio line emission at cm wavelengths has been observed from very low energy
electronic transitions between very high excitation levels of H I, such as from level n =
110 to n = 109 and from n = 105 to n = 104.
II) Continuous radio emission
• The electrons in an H II region move freely through the gas, some recombining with ions
and sometimes, by collisions, exciting atoms or ions (leading to the emission of forbidden
lines), but more often interacting with ions in a free-free transition. When an assembly of
electrons and ions (a plasma) is involved, the individual free-free emissions add up to a
continuum, this continuum radiation occurs predominantly at IR and Radio wavelengths.
In short, an H II region is a source of radio emission characterized by the mean energy of
the electrons-the temperature of the gas. To distinguish this emission from synchrotron
radiation, we use the term thermal Bremsstrahlung.
Supernova remnants
• Ejected material from supernovae becomes part of the ISM. Moreover, the ejected matter
sweeps up any surrounding gas and dust as it expands; this produces a shock wave that
excites and ionizes the gas, which then becomes visible as an emission nebula. X-rays
emitted by supernovae are ionizing nearby gas. Supernova remnants are radio emitters
because of their synchrotron radiation.
• The huge shock waves plow through the IS gas heat it to at least few million Kelvins in
the zone just behind the wave. This gas emits X-rays by free-free emission because it has
such a high temperature. See the picture of Tycho’s SNR in radio and X-ray (Figure 1510).
A series of different types of fusion reactions in highmass stars lead to luminous supergiants
• When helium fusion ceases in the core, gravitational compression
increases the core’s temperature above 600 million K at which
carbon can fuse into neon and magnesium.
• When the core reaches 1.5 billion K, oxygen begins fusing into
silicon, phosphorous, sulfur, and others
• At 2.7 billion K, silicon begins fusing into iron
• This process immediately stops with the creation of iron which
can not fuse into larger elements and a catastrophic implosion of
the entire star initiates.
High-mass stars die violently by blowing
themselves apart in supernova explosions
Remnants of supernova explosions can be
detected for millennia afterward
The most famous “before and
after” picture Supernova
1987 A
Planetary nebulae
• Planetary nebulae differ from H II regions in that they are more compact and of higher
surface brightness and have a different exciting source. Closer examination reveals that
the nebula is excited by very hot central Gas densities in the nebulae surrounding these
stars are higher than in H II regions; hence, collisions between electrons, atoms and ions
occur more frequently. Collisional excitation and de-excitation are therefore significant,
so the spectra of planetaries differ in important ways from those of H II regions, mainly
by emission from forbidden lines.
Low-mass stars expand into the giant
phase twice before becoming planetary
nebulae
Stages in the evolution of low-mass stars
beyond the helium flash:
•
•
•
•
Movement to horizontal branch
Core helium fusion
Asymptotic GIANT branch (AGB)
Planetary nebula formation
Interstellar radio lines
The neutral-hydrogen line at 21 cm
• Where IS atomic gas is cold, hydrogen is neutral and in its ground state. This ground state
has two levels separated by a very small energy difference. The reason for this
phenomenon lies in the fact that both the proton and the electron have an intrinsic spin.
Both the proton and the electron have angular momentum.
• According to the rules of quantum physics, the electron and the proton can be oriented in
the atom so that the two spins either align or oppose each other. If the spins oppose, the
total energy of the atom is just a bit less than if the spins align. As usual, the atom prefers
to be in the lower energy state. Suppose that the spins are aligned. Eventually the electron
flips over and emits a low energy photon energy that corresponds to a wavelength of
21.11 cm (1420 MHz).
• How are the protons and electrons in hydrogen atoms aligned in the first place? By
collisions with electrons and other atoms. The gas in IS space is very sparse, and
collisions between two atoms occur only once every few million years.
• On the other hand, once the spins in a hydrogen atom have become aligned, about 10
million years, on the average, pass before the proton flips and the atom drops to its lowest
energy state. It emits a 21 cm photon. This is a rare event for any one atom. But because
so many hydrogen atoms exist in IS space, enough are emitting 21 cm radiation at any
given time that the IS gas radiates strongly at this wavelength and detected with radio
telescopes.
Molecular Lines (IS molecules)
• IS molecules range from simple molecules like CO, CN, and OH to such complex organic
molecules as formaldehyde (H2CO) and methanol (CH3OH), all found by searching for
spectral lines at radio wavelengths. The study of these molecules will eventually lead to a
better understanding of the chemistry of the interstellar medium (see Table 15-1).
•
•
•
•
•
Although grains make up a very small fraction of the total IS medium, they influence the
form of the gas. Grains are probably the sites of molecule formation for some of the
simpler molecules (at least H2). Their surfaces act as catalysts by allowing atoms (or
simple molecules) to stick to them so that there is time for a second atom to land, interact,
and form a molecule that then evaporates back into the gas. Dust grains also shield
molecules from dissociation by UV radiation, thus letting the molecule population build
up within a cloud.
Optical astronomers in the 1930s made the first discoveries of some molecules (such as
CH and CN). The radio search for molecules made of many atoms began in earnest in the
1960s. The first detected molecule was OH by radio in 1963.
Molecules appears to connected with dust because OH and CO lines are fairly widespread
and found in large dust clouds. Many of these dense clouds lie in the direction of and are
connected to H II regions; (e.g. Orion nebula). The number densities in such clouds are
estimated as 109 to 1012 H2 molecules/m3, other molecules are, far less abundant although
more readily observed. The cloud temperature are low, usually 10 to 30 K and sometimes
as high as 100 K.
We can not observe H2 directly by radio because it emits no lines in that wavelength
range. Instead, we observe CO and assume that it acts as a good tracer of H2.
Most IS molecules are concentrated in dark, dense and cold clouds (called molecular
clouds). The dust in these clouds shields out the UV light that destroys molecules.
Giant Molecular clouds
• Observations show that the bulk of the material of the ISM is bound up in complexes of
giant molecular clouds. Typical properties are:
1- They consist mostly of molecular hydrogen; many other molecules are present but make
up only a small fraction of the mass.
2- The cloud complexes have average densities of a few hundred million molecules per cubic
meter; the individual clouds are slightly denser.
3- They have sizes of a few tens of parsecs.
4- The total mass of the complexes range from 104 to 107 M๏; 105 M๏ is the typical. Masses
of individual clouds are about 1000 M๏.
• The core of these clouds are unusual places compared with the average ISM. Here the
temperatures are a frigid 10 K and densities get as high as 1012 molecules /m3.
• Giant H II regions, which surround young, massive stars, are always found near
molecular cloud complexes. This proximity suggests that giant molecular clouds play the
essential role in the process of star formation. Chemically, clouds with active star
formation appear to be very different from those in which none in now taking place.
Stars form out of enormous volumes of
dust and gas
• Interstellar medium
• H2 (mostly), CO, H2O, NH3,
H2CO
• Most is concentrated in giant
molecular clouds
Intercloud Gas
• Optical observations of IS absorption lines and 21-cm data indicate that a large fraction of
the IS gas consists of cool clouds along with denser molecular clouds (H I regions).
Filling the space between these clouds is an ionized gas, some of which is very hot.
• Radio observations indicate one partially ionized component at 10,000 K with a mean
density of 104 ions/m3, roughly. UV and X-ray observations reveal a hotter component at
106 K (Called coronal interstellar gas). It appears to occupy the largest volume of local IS
space.
The evolution of the interstellar gas
• Driven by star birth and death, the IS gas evolves in various forms. Supernovae play an
important role in IS gas dynamics. A supernova blasts a huge amount of energy (about
1044 J) and material ( 1 to 50 M๏) into space. This material blown off by a supernova
expands as a shell into the ISM. The expanding shell compresses and heats up the IS gas;
behind the shell the gas is left hot and rarefied.
• The supernova shells are large structure, a few of them hundreds of parsecs (up to 3 kpc)
in diameter, these sweep through the IS gas, they heat it to about 50,000 K and thin it out
to a low density. In the process, they distort and destroy existing cool IS clouds. These
large expanding shells may power the evolution of a large fraction of the IS gas.
• Consider the ISM as containing several components, which may evolve from one to
another as the imbedded stars go through their various life stages:
1. H II regions. Zones of glowing, ionized hydrogen surrounding young, hot stars (O & B);
contain a minor amount of IS gas, perhaps 10 million M๏ total in the Galaxy; temperature
about 104 K; density about 106 H II/m3.
2. H I regions/diffuse neutral clouds. Clouds of cool neutral hydrogen roughly 5 pc in
diameter and each containing about 50 M๏ of material; total mass in Galaxy may be 3
billion solar masses; temperature about 100 K; density about 108 atoms/m3.
3. Molecular clouds. Small to huge, containing mostly molecular hydrogen (H2); total mass
of a few billion M๏; temperature as low as 10 K; density about 109 molecules/m3 or
greater. Although they occupy less than 1% of the IS space, they contain a substantial
portion of the matter that constitutes the ISM. Stars form out of the dense molecular
clouds, some fragments of which develop into H II regions.
4. Intercloud medium. A relatively hot gas composed largely of neutral hydrogen (and
therefore observable at 21 cm) plus about 20% ionized gas, including electrons
(observable in the radio continuum). This gas surrounds the cooler IS clouds and fills
about 20% of the volume; temperature from 5000 to 10,000 K; density about 3x105
atoms/m3 and 5x104 electrons/m3.
5. Coronal gas. A very hot (106 K), low-density (<104 particles/m3), ionized gas that
permeates the rest of IS space and occupies well over half of it, perhaps as much as 70%.