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Lecture 10: Interstellar

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Lecture 10: Interstellar 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 Diagram of spectrum to a typical diffuse gaseous of the Orion 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. 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 NGC1976 η Carinae nebula NGC3372 30 Doradus (in LMC) NGC2070 Lagoon nebula M8 NGC6523 Rosette nebula NGC2237 Trifid nebula M20 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: 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 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 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, 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.
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