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An introduction to the physics of the interstellar medium I. Overview II. Thermal processes in the ISM III. Hydrodynamics in the ISM IV. Gravity in the.

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Presentation on theme: "An introduction to the physics of the interstellar medium I. Overview II. Thermal processes in the ISM III. Hydrodynamics in the ISM IV. Gravity in the."— Presentation transcript:

1 An introduction to the physics of the interstellar medium I. Overview II. Thermal processes in the ISM III. Hydrodynamics in the ISM IV. Gravity in the ISM V. Magnetohydrodynamics in the ISM

2 An introduction to the physics of the interstellar medium I. Overview Patrick Hennebelle

3 The Milky Way about stars mass of gas : 5% mass of stars

4 STARS Hot Ionised Gas Warm Ionised Gas Warm Neutral Gas Cold Neutral Gas Heavy Elements Kinetic energy Radiation Cosmic Rays Molecular Gas Dense Cores Cooling, mhd turbulence, mhd, gravity Gravity, mhd Accretion discs Large scale structures The Interstellar Cycle Planets

5 What do we find in the Interstellar Medium ? -photons -gas -magnetic field -dust -cosmic rays As we will see, they all interact with each other, they carry comparable energies and none of them can be ignored… Welcome in the ISM!

6 The Radiation Due to the contribution of the stars, the dust, the CMB, the hot gas. It heats and ionizes the gas. It fluctuates very strongly (neighbourhood of a O or B stars). Radiation energy is about: Urad = erg/cm 3 Energy flux as a function of frequency (Black 1987) submillimeter CMB Far Infrared Dust Visible starlight UV OB star H,He absorption X-rays Extragalactic HIM

7 The Gas in the ISM The ISM is very inhomogeneous in density and temperature. HIM: Hot Ionised Medium ionised, 10 6 K, 0.01 cm -3, 10 8 solar masses produced by supernovae explosions WIM: Warm Ionised Medium ionised, 8000 K, 0.5 cm -3, 10 9 solar masses WNM: Warm Neutral Medium atomic neutral, 8000 K, 0.5 cm -3, solar masses CNM: Cold Neutral Medium atomic neutral, 70 K, 50 cm -3, solar masses Molecular Hydrogen: neutral, 10 K, cm -3, 10 9 solar masses Comparable Pressure Higher pressure

8 Abundances of heavy elements He/H~10 -1 D/H~ C/H~ N/H~ O/H~ Metals: Although the abundances are relatively small, the heavy elements are playing a very important role in the physics of the ISM in particular for the thermodynamics.

9 Galactic south pole (IRAS) Emission of the dust associated to the HI. 21 cm line, emission and absorption Kulkarni & Heiles (1987) Parkes Survey Atomic hydrogen

10 Atomic hydrogen: a thermally bistable Medium Warm Neutral Medium (0.5 cm -3, 10 4 K) Cold Neutral Medium (50 cm -3, 10 2 K) HI Spectra (21cm) Absorption: Emission: Emission Heiles 2001 Absorption

11 Molecular Clouds of the Galaxy Extinction Map in Near Infrared of the Orion Molecular Cloud (Bontemps et al., 2MASS data) CO Map of our Galaxy (Dame et al. 1987) 50 pc

12 Padoan, Cambresy et al. 04 Ward-Thompson et al. 01 Oph molecular cloud seen at 1.3 mm with the Iram 30m telescope Several dense cores can be seen Taurus molecular cloud seen in infrared extinction (total mass about 10 4 Ms) L1544 dense core (belonging to Taurus) Seen at 1.3 mm (dust emission) total mass is about 2 solar mass Molecular clouds and Dense cores molecular clouds ( Ms) contain dense core ( Ms) roughly times denser Molecular clouds are often (not always) filamentary.

13 The Mechanical Energies Non thermal line width implies that significant motions are observed in the gas. Thermal Energy: P/k=4000 K/cm 3, Utherm= erg/cm 3 Kinetic Energy: sonic to supersonic velocity dispersion (up to Mach 5) => Eturb / Etherm = 1-20 PDF of the Mach number in HI clouds (Crovisier 91 )

14 Big powerlaws in the sky….. Turbulence ? Density of electrons within WIM (Rickett et al. 1995) Intensity of HI and dust emission Gibson 2007

15 The Larson Laws (1981) A Turbulent Cascade ? In the molecular gas, Power laws are observed over several order of magnitude. Universal Mass Spectrum dN/dM M -1.6 (Heithausen et al.98) L (pc) M (solar mass) Velocity dispersion (km/s) Mass versus size of CO clumps Velocity dispersion versus size of CO clumps Falgarone 2000

16 The structure of the gas appears to be fractal The molecular gas presents the following power laws over 4 orders of magnitude: This suggests that there is a turbulent energy cascade, presumably from the large towards the small scales In any case, there is a continuous energy injection in the ISM. What are the possible sources ? Falgarone et al. 91

17 Sources of Energy Injection (Mac Low and Klessen 2004) How to produce and sustain turbulence in the ISM ? Energy Dissipation: turbulent energy is dissipated in about a crossing time: For typical numbers, we get: Sources of Energy must compensate this dissipation. What are the possible sources ?

18 Supernovae Explosions: probably the most significant source of energy injection in the ISM : mean mechanical energy of a supernovae remnant : galactic supernovae rate (1/50 year -1 ) : efficiency of energy injection

19 and also: Galactic Differential Rotation (e.g. through the magneto-rotational instabilities) Gravitational Instabilities (spiral arms of the galaxies) Massive Stars (ionisation pressure, winds) Jets and Outflows

20 The Magnetic Field Magnetic field is observed through the Galaxy. Its origin is still debated, likely to be some sort of dynamo. Magnetic Energy: magnetic intensity: 5 G => Emag / Etherm = 1-5 Magnetic Intensity as a function of density (Troland & Heiles 86) Han et al. 06

21 The Dust solid phase of the ISM : 1% of the gas phase complex chemical and physical structure (amorphous carbon, graphite, silicate…). Size a, from a=0.01 to a=1 m, dn=A a -3.5 da (Mathis et al. 1977). Play an important rôle for the thermal (heating and cooling of the gas), for the chemistry, for the absorption and diffusion of the interstellar radiation. Emission of the dust (Desert et al. 1990)Absorption by the dust (Draine & Lee 1984)

22 Energy equipartition in the ISM U rad ~U therm ~U turb ~U mag ~U cosmic ~ erg cm -3 Physical reasons not fully understood yet but suggest that the various processes are coupled to each other and exchange energy. Some examples will be given in the lectures.

23 An introduction to the physics of the interstellar medium II. Thermal processes Patrick Hennebelle

24 Thermal processes in the ISM -Heating -Cooling -Cooling time -Thermal balance and equilibrium

25 Heating Processes Several heating mechanisms. Most of them are based on the ionisation of an ISM components by an energetic radiations. Then the electrons quickly (~1 year) interact with the ISM gas and thermalise. Cosmic rays ( the first proposed source of heating ) Low energy protons (few MeV) ionise the gas during collisions. The heating is related to the ionisation rate, p, induced by cosmic rays ( Black 1987, Lequeux 2002 ). =>relatively low value but may be important in well shielded clouds as the dense cores.

26 Photoelectric effect on small dust grains and PaH UV radiation leads to the ionisation of interstellar grains (particularly the small one). The electrons undergo collisions in the ISM and heat it ( Watson 1972, de Jong 1977, Draine 1978, dHendecourt & Léger 1987, Bakes & Tielens 1994 ). Full calculations not straighforward. Seems to be enough to explain the temperature of the diffuse gas (WNM, CNM)

27 Other heating processes: -ionisation of atoms and molecules (like C) by UV -X-rays (ionisation of ions and atoms, efficient at low column density) -chemistry (formation of H 2 onto grains) -exchange between gas and grains (probably important deep inside the dense cores, FIR heats the dust) -mechanical heating (discussed in next lectures)

28 Cooling Processes Let us consider an atom with 2 energy levels l and u which is excited from l to u by collisions and desexcited by spontaneous radiative emission or collisions (radiatively induced excitation of desexcitation can be neglected) with some impactor of density n i. Stationary state: If the gas density is smaller than some critical density:

29 The most important lines for the cooling are the fine structure lines which result of the L-S coupling. Those lines are forbidden. Lequeux 2002 Energy of the transitions are: -158 m (or 92 K) for CII -63 m (or 230 K) and 145 m for O -370 m and 610 m for C Expression of the cooling by CII due to collisions with electrons and H:

30 Other cooling mechanisms -atomic cooling in other type of transitions Ly of atomic hydrogen. Level n=2. -electron recombination onto positively charged grains as it removes an energy 3/2kT from the gas phase. Reverse of the photo-electric effects. Can be important at high temperature or if the grains charge is significant. -molecular cooling (e.g. Goldsmith & Langer 1978 ) for example at low temperatures: CO molecule rotation lines At high temperature H 2 lines can be very efficient

31 The various heating and cooling contributions of the atomic phase Wolfire et al. 1995

32 Cooling Processes: summary -due to collisional excitation followed by a radiative deexcitation (photon energy is lost) H (Lyman alpha) for T> few 1000K, and C+, O for T < few 1000 K -mainly proportional to the square of density -strong dependence on the temperature ( depending on whether the dominant lines are saturated ) Cooling function of the ISM Dalgarno & Mac Cray 72

33 Cooling Times Cooling time for HIM: 100 Myr is very long (it is metastable) Cooling time for WNM: 1 Myr is comparable to its dynamical time Cooling time for CNM: 0.01 Myr is shorter than its dynamical time

34 Thermal balance and Thermal Instability Temperature is determined by the equilibrium between cooling and heating Heating is proportional to the density and depends smoothly on the temperature Cooling is proportional to the square of the density and depends stiffly on the temperature If (T) does not vary much with T, the gas is thermally unstable Cooling function of the ISM Dalgarno & Mac Cray 72 HIM (metastable) WNM CNM heating

35 Thermal equilibrium curve (Field et al. 69, Wolfire et al. 95) CNM WNM Unstable Field 65: performs linear stability analysis of the radiatively cooling fluid equations. Obtains the isobaric criteria for instability: Wolfire et al. 95

36 Influence of stronger radiation flux on the thermal equilibrium Wolfire et al. (1995)

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