Interstellar Medium (ISM)
Transcript of Interstellar Medium (ISM)
Interstellar Medium (ISM)
Hiroyuki Hirashita (平下 博之) ASIAA
1. ISM Composition/Phases 2. Interaction with Radiation 3. Multi-Phase ISM 4. Dynamical Processes 5. Summary
Topics
An ISM lecture needs one semester. We only briefly touch main topics and show some examples of numerical research (to motivate your further studies).
1. ISM Composition/Phases
Elemental Composition (Solar)Lodders (2010)
Metals/heavy elements
Metallicity
Hydrogenatomic hydrogen H I ionized hydrogen H II
ionization potential: E = 13.6 eV E = h(c/λ) → λ = 912 Å
H+e-H+
H+ H+e-
e-
molecular hydrogen H2
e-
Atomic Hydrogen (21 cm)
Ionized Gas (Hα; recombination lines)
H II regions ← mainly ionized by UV photons from OB stars ~ Recent star forming regions (star-formation tracer)
Finkbeiner (2003)
Hot Gas (X-Ray)
Diffuse hot gas (> 106 K) heated by supernovae, etc.
Molecular Gas (CO)
Dame et al. (2001)
Useful tracer of cold H2 gas: CO-to-H2 conversion factor XCO = NH2/WCO = 1.8 × 1020 cm-2 K-1 km-1 s (in the MW)
Dust
Lund Observatory
Optical (λ ~ 0.5 µm) Image of the Milky Way
Dark clouds in the Milky Way ← Dust extinction (= absorption + scattering)
Dust: Small particles most likely silicate and graphite, or something similar
Dust Emission in the Far InfraredAKARI 140 µm
Dust emission
2. Interaction with Radiation
Radiation Transfer
Definition of Intensity (energy flow on a line) dE = Iν dA dt dΩ dν Iν [erg/cm2/s/sr/Hz]
dA
dΩ
n
dI⌫ds
= �⌫⇢I⌫ + j⌫
Radiation Transfer Equation
s
absorbing and emitting medium
Iν(s): intensity
Emission jν [erg/cm3/s/Hz]Absorption
κν [cm2/g]: mass absorption coefficient ρ [g/cm3]: density
Optical Depth
Optical depth: dτν = ρκνds ∝ “number of absorptions” τν >> 1: optically thick τν << 1: optically thin
dI⌫ds
= �⌫⇢I⌫ + j⌫
Sν = jν/(ρκν): source function Sν = Bν(T) (Planck function) for LTE
dI⌫d⌧⌫
= �I⌫ + S⌫
A Simple Solution
For a constant source function:
τν
absorbing and emitting medium
Iν(τν): intensityIν(0)
dI⌫d⌧⌫
= �I⌫ + S⌫
I⌫(⌧⌫) = I⌫(0)e�⌧⌫ + S⌫(1� e�⌧⌫ )
Extinction of bkg light Emission
Dust Extinction
Magnitude: mλ = –2.5 log Iλ + const. Extinction (= absorption + scattering): Aλ = mλ(s) – mλ(0) = 2.5[log Iλ(0)e –τλ –log Iλ(0)] = (2.5 log e) τλ
Iλ(0) Iλ(0)e –τλ
Extinction Curves
Pei (1992)
Grain size distribution n(a) ∝ a – 3.5
(Mathis et al. 1977)
amin = 0.005 µmamax = 0.25 µm
A λ/A
0.44
µm
shorter wavelength ↓ larger extinction
“Reddening”
Scicluna & Siebenmorgen (2015)
The “effective” extinction curve (attenuation curve) depends on the clumpiness.
Effects of Complicated Dust Distribution
Infrared Emission of Dust
Li & Draine (2001)
Large grains (LGs) Silicate + carbonaceous material
Excess by very small grains (VSGs)
Wavelength (µm)
Inte
nsity
Effects of UV Radiation
Orion Nebula (HST)
H II Region
NLyc =4⇡
3R3
Sn2�B
Number of Lyman continuum (> 13.6 eV) photons emitted per time.
Case B: Lyman series/continuum (transitions down to n = 1) are optically thick. Recombination rate to n >= 2 levels contributes to the loss of ionizing photons.RS: Strömgren radiusRS ' 1.2
⇣ n
103 cm�3
⌘�2/3
NLyc
5⇥ 1049 s�1
!1/3
pc
Photodissociation Region (PDR)
Tielens & Hollenbach (1985)
Field et al. (1966)
Lyman-Werner bandsDissociation
Abundant Line Emissions from PDRs
Bisbas et al. (2021)
3. Multi-Phase ISMMcKee & Ostriker (1977)
HeatingFormation of free electrons by radiation causes heating (energy transfer from radiation to gas).
Major heating mechanisms: Photoionization (H + γ → H+ + e-) Photoelectric emission (dust + γ → dust+ + e-)
e-
CoolingThermal energy is converted into radiation: ex.
nk
nj
nγjk nγkj Aul
n: number density of colliding species
Ejk
e- A A
A
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⇤ = nX
j<k
Ejk(nj�jk � nk�kj)
excited
Two-Phase Model of Neutral MediumWolfire et al. (1995)
Thermal equilibrium: Λ = Γ (cooling) = (heating)
Lyα line cooling ~ Photoelectric heating at low density (T ~ 104 K)
C II line cooling ~ Photoelectric heating at high density (T ~ 100 K)
Field et al. (1969)
ISM Structures in SimulationAoyama, Hirashita, et al. (2017)
Realistic structures are more dynamic, but we still see some characteristic features in the phase diagram.
Molecular Clouds
See Yueh-Ning’s talk for the processes leading to star formation.
Image courtesy of Nigel Sharp (NOAO), NSF
ESO
Various Tracers and Critical Densitynu
nl
nγlu nγul Aul
Statistical equilibrium: nlnγlu = nunγul + nuAul γlu = (gu/gl)γul exp(-Eul/kBT)
nu
nl=
gu/gl exp(�Eul/kBT )
1 + ncr/nncr ⌘
Aul
�ulCritical density
Critical density ~ Typical density which the line is suitable for tracing. Critical densities at 10 K: CO(1-0): 4×103 cm-3, CO(3-2): 7×104 cm-3 HCO+(1-0): 2×105 cm-3, HCO+(3-2) : 3×106 cm-3 HCN(1-0): 2×106 cm-3, HCN(3-2): 6×107 cm-3
n: number density of colliding species
4. Dynamical Processes
See Hung-Yi’s talk for hydrodynamic equations.
Characteristic ScalesFree-fall time: gravitational contraction tff ~ 1/(Gρ)1/2 Crossing time: hydrodynamic effects tcr ~ L/v v = cs = (γP/ρ)1/2 ∝ T1/2 (for sound wave) v = vA = (B2/4πρ)1/2 (for Alfvén wave) Cooling time: gas cooling tcool ~ kBT/nΛ Jeans length/mass (typical size/mass for gravitationally bound structures) λJ ~ cs tff ∝ (T/ρ)1/2 MJ ~ ρλJ3 ∝ T 3/2/ρ1/2 Estimate these quantities for the objects you like.
ShocksConservation laws in hydrodynamic equations:
ρ1 u1 → p1 T1
ρ2 u2 → p2 T2
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p1 + ⇢1u21 = p2 + ⇢2u
22
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u2
✓1
2⇢2u
22 + U2
◆� u1
✓1
2⇢1u
21 + U1
◆= u1p1 � u2p2
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U =1
1� �p
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p2p1
=2�
� + 1M2 � � � 1
� + 1
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u2
u1=
⇢1⇢2
=2
� + 1
1
M2+
� � 1
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M2 ⌘ ⇢1u21
�p1=
u21
C21
Mach number
Expansion of H II Region
ri
ionization frontShock
Conservation laws at the ionization front and shock front Constant ionization photon luminosity (ρII ∝ ri -3/2; see Strömgren radius)
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riri0
=
✓1 +
7
4
CIIt
ri0
◆4/7CII : Sound speed in the H II region ri0: initial radius of H II region (under ρII =ρI)
External Ionizing RadiationIliev, Hirashita, Ferrara (2006)
Ionizing radiation
Ionizing photons incident on a cloud could cause compression, inducing star formation, or evaporation, suppressing star formation.
Supernova Blast Wave
rs ~ (E/ρ1)1/5t2/5
Sedov solution
rs
ρ1Based on the pressure inside the SN remnant (described by E and rs), and assuming a strong shock, we obtain the relation Vs and rs (under a given ρ1.
See Yen-Chen’s talk for SN explosion.
Supernovae and Dust DestructionHu et al. (2019)
5. SummaryOverviewed the following processes in the ISM: (1) Radiation transfer. (2) Ionization and dissociation. (3) Heating and cooling. (4) Multi-phase ISM. (5) Dynamical processes (hydrodynamic treatment of the ISM).
Further Reading
• Spitzer, L. 1978, “Physical Processes in the Interstellar Medium” • Tielens, A. G. G. M. 2005, “The Physics and Chemistry of the Interstellar Medium”