Chapter 3 THE INTERSTELLAR MEDIUM

Transcription

1 Chapter 3 THE INTERSTELLAR MEDIUM Introduction The interstellar medium (ISM): the gas and dust distributed between stars in a galaxy In the Milky Way: mass of gas mass of dust : M dust 0.1M gas ISM is generally a small fraction of a galaxy s luminous mass: 0 % for an elliptical 1 25 % for a spiral (increases from Sa to Sd) % for an irregular Very diffuse: in the plane of the Galaxy, particle number density 10 3 to 10 9 atomic nuclei m 3 Mixture of: gas remaining from the formation of the galaxy gas ejected by stars gas accreted from outside (such as infalling diffuse gas or the ISM of accreted galaxies) 1

2 Important (1) in galaxy evolution - gas promotes star formation in denser regions absorption by dust allows molecular clouds to cool Important (2) observationally emission lines from gas are prominent and can be used to observe dynamics Chemical composition is about 90 % H, 9% He, plus a trace of heavy elements (expressed by numbers of nuclei) Heavy elements in the gas can be depleted into dust grains 2

3 Spectroscopy of Interstellar Gas Gas in the ISM readily emits detectable radiation The very low density of gas allows detection of forbidden lines : spectral lines not normally observed in the lab low transition probabilities in the lab the excited states get collisionally de-excited before they can radiate but in the ISM, although collisional times are the lifetimes of the forbidden states, the huge number of atoms in the ISM means that these are commonly observed In astronomy we use notation such as HI, HII, HeI HeII and HeIII where: I neutral atom II singly ionised positively charged ion III doubly ionised positive ion etc. So, HI is H 0, HII is H +, HeI is He 0, HeII is He +, HeIII is He 2+, LiI is Li 0, etc. A negatively charged ion, such as H, is indicated only as H, although few of these are encountered in astrophysics Square brackets indicate a forbidden line e.g. [OII] 3

4 Figure 1: The Milky Way with +/- 10 deg of Galactic plane (360 deg in longitude) in various wavebands [image credit: Jodrell Bank, Leiden Dwingeloo, Max Plank Institute, IRAS, COBE, A. Mellinger, J. Friedlander, S. Digel, ROSAT, NASA Goddard Flight Center] 4

5 Cold/Warm Gas: the 21 cm Line of Neutral Hydrogen Atomic hydrogen, HI, emits at 21 cm wavelength (radio) from hyperfine splitting of ground state cool/warm ISM T 10 to 100K in high density regions, 10 3 to 10 4 K in lower density regions spin-flip transition coupling of nuclear and electron spin forbidden line Energy difference between the two spin states: E = J = ev producing emission with a rest wavelength: λ 0 = hc/ E = cm and rest frequency: ν 0 = E/h = MHz 5

6 Transition probability, A = s 1, so lifetime of an excited state is 1/A = 11 million years 21cm transition itself cannot be observed in a laboratory But hyperfine splitting of ground state observed as the Lyα lines (UV) In the ISM, the 21cm line is observed primarily in emission, but can also be observed in absorption against a background radio continuum source HI observations have many uses: one critically important application is to measure the orbital motions of gas to determine rotation curves in our own Galaxy and in other galaxies HI observations can map the distribution of gas in and around galaxies 6

7 Figure 2: M83, The Southern Pinwheel Galaxy, type SABc, showing its extended disk at 21 cm radio wavelength in red on the left, with the UV image superimposed (near- UV in green, far-uv in blue). On the right is the UV image only showing near-uv in yellow and far-uv in blue. [Credit: NASA/JPL/Caltech/VLSA/MPIA]. 7

8 Cold Gas: Molecules Molecular hydrogen and other molecules T 10 to 100K, relatively high density (cold dusty molecular clouds) Molecular hydrogen, H 2, is rare and very difficult to detect directly abundance controlled by sticking of HI atoms to dust grains no radio lines, so no direct way of tracing H 2 in cold dense gas clouds using radio observations but some H 2 band absorption in the far UV can be detected H 2 can be photo-dissociated by UV radiation into atomic hydrogen, HI Other molecules do emit in radio/microwave they act as an indirect tracer of cold dense H 2 gas CO is particularly useful - has strong lines at 1.3 mm and 2.6 mm from rotational transitions CO and H 2 densities are similar so we can use CO as a tracer for H 2 8

9 Hot Gas: HII Regions Hot gas is readily observed in the optical via emission lines from largely ionised gas HII regions are regions of partially ionised hydrogen around hot young stars (O or B type), with T 10 4 K These stars emit strongly in UV Any UV (Lyman continuum) photons with wavelengths λ < 912 Å) will photoionise hydrogen producing H +, i.e. HII ions The ions and electrons recombine to produce excited hydrogen atoms The electrons then cascade down energy levels, emitting photons (radiative decay) Photons emitted in UV, optical, infrared and radio free/bound transitions continuum radiation bound/bound transitions emission lines Prominent optical lines from transitions down to first excited level (n = 2) give the Balmer series Transitions down to ground state (n = 1) give Lyman series (in UV) 9

10 Each series is labelled α, β, γ, δ,..., in order of increasing energy Transitions from n to n 1 levels are the strongest i.e. α lines are the strongest Lyman series lines of hydrogen are: Lyα λ = 1216 Å (in ultraviolet) Lyβ 1026 Å ( ) Lyγ 973 Å ( ) Balmer series lines of hydrogen are: Hα λ = 6563 Å (in optical) Hβ 4861 Å ( ) Hγ 4340 Å ( ) Hδ 4102 Å ( ) Hɛ 3970 Å ( ) 10

11 Collisional excitation can also occur not for H - no levels accessible at collision energies characteristic of HII regions (T 10 4 K) but possible for NII, OII, SII, OIII, NeIII [OIII] lines at 4959Å and 5007Å are particularly prominent Some of the most prominent optical lines of HII regions are: [OII] 3727 Å [OIII] 4959 Å [NeIII] 3869 Å [OIII] 5007 Å Hɛ 3970 Å HeI 5876 Å Hδ 4102 Å [NII] 6548 Å Hγ 4340 Å Hα 6563 Å Hβ 4861 Å [NII] 6584 Å 11

12 In HII regions one Balmer photon is produced for each Lyman continuum photon from the hot star so observations of Balmer lines of nebula gives number of UV photons from star This happens because most H atoms are in the ground state, and are therefore opaque to Lyman photons but transparent to others A Lyman continuum photon from star will be absorbed by a neutral H atom ionises H atom to produce a free electron The free electron is then recaptured (free-bound transition), emitting a continuum photon depending on which state it is captured into : If captured into the ground state (n = 1) emits another Lyman continuum photon back to where we started If captured into n = 2 emits Balmer continuum photon in going to n=2 one UV photon produces one Balmer photon then decays to n = 1 emitting a Lyα line photon which will almost certainly be absorbed again 12

13 If captured to n > 2 it can then decay to n=2 emitting a Balmer line photon, or directly to n = 1 but a transition to n = 1 emits a Lyman line photon that can excite another ground-state H atom, so the process repeats, eventually producing a Balmer line photon So each ionising UV photon from star (λ < 912 Å) will produce on Balmer (line or continuum) photon HII regions and planetary nebulae also produce thermal continuum radiation free-free emission: the free electrons in the HII can interact with protons without recombination electrons are accelerated, producing radiation The resulting spectrum is not blackbody because the gas is transparent to free-free photons: there is no redistribution of the energy of the free-free photons. In fact the spectrum is quite flat at radio frequencies Strengths of the emission lines from HII regions can provide information on temperature, density and chemical composition of the interstellar gas 13

14 Colour optical images of HII regions show strong red/pink and green colours: the red and pink is produced mainly by the Hα line the green is produced by [OIII] and Hβ HII regions are seen prominently in images of spiral and irregular galaxies their emission lines dominate the spectra of late-type galaxies and are valuable for use in measuring redshifts 14

15 Figure 3: The Orion Nebula, M42. The most famous example of an HII region. The gas fluoresces because of the UV radiation from the hot young stars, recently formed in a dense region of gas [Hubble Space Telescope: NASA, ESA, M. Robberto (Space Telescope Science Institute/ESA), Hubble Space Telescope Orion Treasury Project Team.] Figure 4: The optical spectrum of the Orion Nebula, showing very strong emission lines from Hα (red/pink), [OII] (blue), and [OIII] and Hβ (both green). 15

16 Hot Gas: Planetary Nebulae Planetary nebula: compact HII regions around hot evolved stars gas is ejected by star through mass loss UV photons from star ionise gas gas emits photons like HII regions (similar emission process) relatively luminous, with prominent emission lines also observed in other galaxies useful for tracing distribution & kinematics of stars Figure 5: Examples of planetary nebulae: the Ring Nebula (M57), left, and the Helix Nebula, right. Gas has been ejected from a hot, evolved star and the ultraviolet radiation from the star ionises the gas. [Images from the Hubble Space Telescope. Ring Nebula: produced by the Hubble Heritage Team (AURA/STScI/NASA). Helix Nebula: produced by NASA, NOAO, ESA, the Hubble Helix Nebula Team, M. Meixner (STScI), and T.A. Rector (NRAO).] 16

17 Hot Gas: Supernova Remnants (SNRs) Supernovae eject material at very high velocities into the interstellar medium this gas shocks, heats and disrupts the ISM low density components of the ISM can be significantly affected dense molecular clouds are less affected hot gas from supernovae can even be ejected out of the Galactic disc into the halo Supernova remnants have strong line emission. They expand into and mix with the ISM Figure 6: Examples of supernova remnants: the Crab Nebula (M1), left, and part of the Veil Nebula, right. The Crab Nebula is a very young supernova remnant, produced by a supernova observed in the year The Veil Nebula is an older example. [Crab Nebula image from the Hubble Space Telescope: NASA, ESA, J. Hester and A. Loll (Arizona State University). Veil Nebula image from the 0.9m Burrell Schmidt Telescope at Kitt Peak National Observatory, Arizona: NOAO/AURA/NSF.] 17

18 Hot Gas: Masers Some very high density HII regions around young stars or old evolved stars can show maser emission density m 3 population inversion between certain energy states of molecules due to radiative excitation transitions are in the radio the overpopulated state decays by stimulated emission maser emission coherent emission - polarised very narrow emission lines of high intensity OH and H 2 0 masers are observed (e.g. in Orion) useful kinematic tracers 18

19 Hot Gas: Synchrotron Radiation Broad-band non-thermal radiation emitted by electrons moving relativistically in a magnetic field observed in both optical and radio photons are emitted in the instantaneous direction of electron motion polarised perpendicular to the magnetic field Spectacular sources of synchrotron emission are systems with jets young stellar objects with bipolar outflows, or active galactic nuclei, lobes of radio galaxies. 19

20 Absorption Line Spectra If interstellar gas is seen in front of a continuum source, light from the source is absorbed at certain wavelengths A number of interstellar lines and molecular bands are seen in absorption Some absorption features are not well-understood Particularly problematic are diffuse interstellar bands in the IR probably caused by carbon molecules, possibly polycyclic aromatic hydrocarbons (PAHs) Cold interstellar CN molecules: CN has rotational modes which produce radio lines, like most heteronuclear molecules radio lines can be observed directly, but more interesting are the optical lines that are split because of the rotational modes Optical observations of absorption by cold CN in continuum spectra of background stars show relative populations of the rotational modes (from line strengths) and hence the temperature of the CN temperature 2.7 K, i.e., these cold clouds are in thermal equilibrium with the CMB 20

21 Components of the Gaseous ISM It s convenient to divide diffuse gas in the ISM into distinct components also called phases: cold neutral medium neutral hydrogen (HI) and molecules at temperatures T K and relatively high densities warm neutral medium neutral hydrogen (HI) but at temperatures T K and lower densities warm ionised medium ionised gas (HII) at temperatures T 10 4 K and lower densities hot ionised medium ionised gas (HII) at very high temperatures T K but very low densities The phases are pressure-confined and are stable in the long term Ionisation by supernova remnants is an important mechanism in producing the hot ionised medium Cold neutral medium makes up 50% of the ISM s mass, but very small fraction by volume Supernova remnants, planetary nebulae, and giant molecular clouds not normally included in these phases because they are not in pressure equilibrium with the other components 21

22 Interstellar Dust Consists of particles of silicates or carbon compounds relatively small, but broad range in size largest 0.5 µm in size with 10 4 atoms some have 10 2 atoms like large molecules Dust has a profound observational effect absorbs and scatters light extinction diminishes light of background sources e.g.dark nebulae, zone of avoidance for galaxies at low galactic latitudes Dust in galaxies is important because: it allows dense molecular clouds to cool - facilitates star formation it catalyses the formation of molecules e.g. molecular hydrogen 22

23 Interstellar Dust - extinction The absorption and scattering of light by dust is called extinction Light of wavelength λ, specific intensity I λ (i.e. flux) passing through an element of interstellar space will experience a change di λ in intensity I λ due to extinction This is related to the change dτ λ in the optical depth τ λ at the wavelength λ that the light experiences along its journey by di λ I λ = dτ λ Integrating over the line of sight from a light source to an observer, the observed intensity is I λ = I λ 0 e τ λ where I λ 0 is the light intensity at the source and τ λ is the total optical depth along the line of sight What is the loss of light in magnitudes? Magnitude m in some photometric band is related to the flux F (i.e. intensity) in that band by: m = C 2.5 log 10 ( F ) 23

24 where C is a calibration constant (see Appendix A of the Course Notes for more information about the magnitude system) In the presence of extinction, for a particular wavelength λ we have (substituting for the expression for I λ above) : m λ = C 2.5 log 10 ( I λ 0 e τ λ ) = C 2.5 log 10 I λ log 10 ( e τ λ ) = C 2.5 log 10 I λ ln( e τ λ) ln(10) = C 2.5 log 10 I λ ln(10) τ λ = C 2.5 log 10 I λ τ So the observed magnitude m is related to intrinsic magnitude m 0 by: m λ = m 0 + A λ where the intrinsic magnitude, m 0, is the magnitude that the star would have in the absence of interstellar extinction, and A is the extinction in magnitudes, given by: A λ = 2.5 log 10 (e τ ) = τ Note that A λ depends on the photometric band 24

25 For example, for the V (visual) band (yellow-green, centred at 5500 Å) : V = V 0 + A V For the B (blue) band (centred at 4400 Å): B = B 0 + A B A λ is a strong function of wavelength and scales as A λ 1/λ (not as strong as Rayleigh Law 1/λ 4 ) There is much stronger absorption in the blue than in the red reddening by interstellar dust 25

26 The interstellar extinction law. The extinction caused by dust is plotted against wavelength and extends from the ultraviolet through to the near-infrared. [Based on data from Savage & Mathis, Ann. Rev. Astron. Astrophys., 1979.] Colour indices, e.g. B V, are reddened so that the observed value is: (B V ) = ( B 0 + A B ) ( V 0 + A V ) = ( B 0 V 0 ) + ( A B A V ) (B V ) 0 + E B V where (B V ) 0 is the intrinsic value (no extinction) and E B V = A B A V is the colour excess or reddening, which tells us how reddened a source is, based on the extinctions in the two magnitudes For the V photometric band, A V 3 E B V (as shown in the plot above) If the intrinsic colour, (B V ) 0, can be predicted from spectrum, then E B V can be calculated using E B V = (B V ) (B V ) 0 E B V data can then be used to map the dust distribution in space 26

27 Extinction gets less severe for λ 1 µm as the wavelength gets much longer than the grains For sight lines through the Galaxy at the Galactic poles: A V 0.00 to 0.05 mag At intermediate galactic latitudes: A V 0.05 to 0.2 mag In Galactic plane, extinction can be many magnitudes in V and UV (less in IR). Distribution can be patchy (e.g. Baade s Window, in the bulge) Towards the Galactic Centre: A V 20 mag X-rays can pass through dust grains (A K 3 mag) 27

28 Interstellar Dust: polarisation Dust grains are not spherical Spinning dust grains tend to align with their long axes perpendicular to the local magnetic field Preferentially block light perpendicular to the magnetic field: extinction produces polarised light Polarisation will tend to be parallel to the magnetic field polarisation measurements of starlight reveal the direction of the magnetic field Dust also reflects light, with some polarisation - observable as reflection nebulae, where faint diffuse starlight can be seen reflected by dust 28

29 Interstellar Dust: Radiation by Dust Dust absorbs light warms the dust re-emitted as black-body radiation (approximately) So dust has diffuse black-body emission superimposed on reflected starlight spectrum Wien s displacement law states that the maximum of the Planck function B λ of a black body at a temperature T is found at a wavelength λ max = T K m This predicts that the peak of the black-body spectrum for dust at a temperature of T = 10 K will be at a wavelength λ max = 290 µm for dust at T = 100 K will be at a wavelength λ max = 29 µm and for T = 1000 K will be at λ max = 2.9 µm So radiation emitted by dust will found in the infrared, given the expected temperatures of dust 29

30 Star Formation in the ISM Stars form by collapse of dense regions of the ISM under their own gravity i.e. in cores of molecular clouds, where gas is cold ( 10 K) and densities relatively high molecules m 3 Figure 7: A star-forming HII region within M16, The Eagle Nebula. The blue-green colour from the mostly ionised gas is caused by the light of [OIII] and Hβ emission lines from neutral hydrogen atoms. The gas is being ionised by ultraviolet radiation from hot, young stars off the top of the picture. The dark pillars, in contrast, are regions of cold, dense molecular hydrogen gas in which star formation is occurring. They are dark because the cold molecules emit virtually no light and because of the absorption of light by dust mixed with the gas. The ultraviolet radiation is burning away the surface of the cold gas by photoionisation. [Hubble Space Telescope image produced by NASA, ESA, STScI, J. Hester and P. Scowen (Arizona State University).] 30

31 A region of cold gas collapses when its gravitational self-attraction > hydrostatic pressure support For gas of uniform density ρ, the Jeans length λ J is the diameter of a region of gas just large enough for gravitational force to exceed pressure support: λ J = c s π G ρ where c s is the speed of sound in the gas The Jeans mass is the mass of a region that has a diameter equal to the Jeans length: ( ) 4 M J = 3 π(λ J/2) 3 ρ = π 6 ρ λ J 3 The free-fall collapse time T ff, is the time taken for a static cloud to collapse under its own gravity in the absence of gas pressure The free-fall collapse time for a spherically symmetric distribution of mass with a total mass M and initial radius R to collapse from rest is T ff = π R 3 8GM = 3π 32 1 G ρ where ρ is the mean density before the collapse starts Star formation can be self-propagating Stars form, heat up and ionise cold molecular gas outward flow of gas compresses gas ahead of it 31

32 Causes instabilities locally collapse to form new stars The enhanced density in the spiral arms of spiral galaxies means that star formation occurs preferentially in the arms 32