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Home , Emission nebula

Lecture 10: Interstellar gas (ISM) • In spite of the space between the 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 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) line emission • Emission 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 Spectrum of NGC7009, a , but similar Diagram of spectrum to a typical diffuse gaseous of the nebula nebula spectrum. Chemical composition of HII nebulae

element log10N H 12.0 He 11.0 C 8.5 N 8.0 O 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++ + e → (O++)* + e (collisional excitation) (O++)* → O++ + hν (radiative deexcitation) Typical radius and mass of HII regions

Spectral type of star radius of nebula (pc) O5 70–200 B0 20 A0 0.5

They can only readily be observed around stars of types

O to B0 (T* ~ 50 000 K to 25 000 K)

3 Mass: 0.1 to 10 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. and glowing HII regions in the Great 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 NGC1976 η Carinae nebula NGC3372 30 Doradus (in LMC) NGC2070 M8 NGC6523 NGC2237 M20 NGC6514 Below: Lagoon nebula M8 in Sagittarius

Above: Trifid nebula, M20, in Sagittarius

Right: Rosette nebula in Below, , 30 Doradus in the

Above: Orion nebula, M42 Right: η Carinae nebula, in southern Milky Way Hubble Space Telescope images of the Orion nebula Right: detail of centre 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: 3 Nuv = (4π/3) Rs 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

3 Rs =[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 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 15- 10). A series of different types of fusion reactions in high- mass 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, 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 9 12 3 estimated as 10 to 10 H2 molecules/m , 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 7 5 4- The total mass of the complexes range from 10 to 10 M๏; 10 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 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 44 10 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 9 3 of a few billion M๏; temperature as low as 10 K; density about 10 molecules/m 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%.