interstellar gas - university of rochester

Interstellar GasInterstellar Atoms and Molecules

Star FormationMolecular Clouds

March 18, 2021

University of Rochester

Interstellar dust & gas

I Refresher on spectral linesI Interstellar atoms and moleculesI Molecular clouds and gravitational

stabilityI Cloud collapse and star formation

Reading: Kutner Ch. 3.1–3.4 & 14.4–14.5,Ryden Ch. 5.1–5.3, 16.2-16.3, and Shu Ch. 11

HII region NGC 604 in Triangulum (NASA/HST)March 18, 2021 (UR) Astronomy 142 | Spring 2020 2 / 29

Nomenclature of spectral lines

Much of the evidence for the presence of interstellar gas comes from the detection ofatomic and molecular emission and absorption lines.

In astronomy, the ionization states of atoms are indicated with Roman numerals.I Neutral hydrogen (H) is denoted HI (“H-one”).

I Ionized hydrogen (a free proton) is HII (“H-two”); compare to the usual chemicalnotation H+.

I Do not confuse HII with H2, which is molecular hydrogen.

Other examples: Doubly-ionized oxygen (O++) is OIII. In the 106 K solar corona, it iscommon to observe FeIX and FeX (iron atoms that have lost 8 or 9 electrons).

March 18, 2021 (UR) Astronomy 142 | Spring 2020 3 / 29

Spectrum of Hydrogen

Recall that electron transitions in H give riseto radiation of discrete wavelengths.

Lyman: transitions to n = 1Balmer: transitions to n = 2And so on.

Lines within a series are sometimes labeledby Greek letters. E.g., the n = 2 ! 1transition is called Lyman-a, n = 3 ! 1 isLyman-b, etc.


Astronomically significantHydrogen lines

The Lyman-a line (121.567 nm) is commonly

used in cosmology since it is redshifted intothe visible spectrum.

Ha is the 656.28 nm line in the n = 3 ! 2Balmer transition. It shows up a lot in HIIregions due to atomic recombination.

The 21 cm HI line is due to an electrontransition in the hyperfine levels of the 1sground state. It is a tracer of neutral

hydrogen.


Fine and hyperfine structure in the spectrum

When you look with increasing precision at theH spectrum, you find that the energy levels aresplit into sub-levels.

In QM, we approximate the interactions insidethe atom to varying degrees of accuracy.Lowest order approximation Coulomb force“Fine structure” corrections for relativistic

motion of the electron.“Hyperfine structure” interactions between

electron orbital angular momentumand nuclear spin angularmomentum (spin-orbit coupling)

The dominant interaction in the hydrogen atom(Coulomb attraction of e and p) gives rise to theenergy levels. Sub-dominant interactions causesplitting of the levels, observed as “fine” and“hyperfine” structure.


Atomic excitationAtoms can be excited via one of two methods:

n = 2

hν∆Ee–

n = 1(a) (b)

Copyright © 2010 Pearson Education, Inc.Photoexcitation (or absorption) occurs whenthe atom absorbs a photon with an energyexactly equal to the difference betweenenergy levels.

Collisional excitation occurs when the atomcollides with another particle (normally afree electron). Some of the free particle’skinetic energy will transfer to the internalenergy of the atom.


Atomic de-excitationAtoms can drop down to lower energy levels via one of three methods:

2

1(a) (b) (c)

hν e–hνhν

hν

Copyright © 2010 Pearson Education, Inc.Spontaneous emission

occurs when the atomreleases a photon withenergy equal to the differencebetween energy levels (theinverse of photoexcitation).

Stimulated emission occurswhen the atom encounters aphoton with energy equal tothe difference between theenergy levels.

Collisional de-excitation

occurs when the atom’sde-excitation is triggered bya collision with another freeparticle.


Atomic de-excitation

I Spontaneous emission is directly related to the instability of the excited states.

I In stimulated emission, the emitted photon has the same phase and direction as thestimulating photon, thereby amplifying the original photon signal. The rate ofstimulated emission is proportional to the intensity of the radiation field at therelevant frequency.

I No photon is emitted in collisional de-excitation, as the free particle carries away theenergy difference as additional kinetic energy.


Ionizing gas in the interstellar medium

When an atom absorbs a photon with energy greater than the ionization potential for itselectron(s), the electron will be stripped from the atom. Photoionization results in thenow free electron having a kinetic energy equal to the energy difference of the incomingphoton and the ionization potential.

Collisional ionization can also occur, when a free electron with total kinetic energy greaterthan the ionization potential collides with the atom. The final speeds of the two freeelectrons is based on conservation of energy and momentum.

Free electrons can be captured by the ions, with a photon carrying away the excessenergy. Recombination can result in the electron being located in any of the availableenergy states (ground or excited).


Kirchoff’s Laws of Spectroscopy


Evidence for interstellar gas

Apparent through spectral lines seen in absorption against stars or in emission elsewhere.

“Nebulae:” hydrogen emission lines, plus a variety of other bright lines not readilyidentifiable in the laboratory.I Bowen 1928: extra lines are mostly forbidden lines in the spectrum of the ions,

mostly neutral, singly or doubly ionized, of the more abundant elements:I Oxygen (O/H ⇡ 7 ⇥ 10�4)I Nitrogen (N/H ⇡ 1 ⇥ 10�4)I Carbon (C/H ⇡ 3 ⇥ 10�4)

It was obvious right away that the lines must originate in very low density diffusematerial, which is difficult to reproduce in labs on Earth. Best vacuum on Earth is now⇠ 103 cm�3; in space ⇠ 0.1 � 1 cm�3.


http://adsabs.harvard.edu/abs/1928ApJ....67....1B

Forbidden lines

What are forbidden lines? And how can we see them if they are forbidden? In short, theterm is a bit of a misnomer.I Atomic and molecular transitions are governed by selection rules that tell us what

changes in quantum states are allowed during the transition. The selection rules aredetermined by the physical process governing the transition.

I But we approximate these interactions to varying degrees of accuracy, with thelowest-accuracy (“first-order”) effects dominating and then sub-dominant“higher-order” effects added as corrections.

I Selection rules that apply to first-order effects may forbid certain transitions, but thehigher-order interactions may allow the transitions (albeit at low rates). These are the“forbidden transitions” which produce forbidden lines.


https://en.wikipedia.org/wiki/Selection_rule

Planetary Nebula NGC 7662 & its visible spectrumNGC 7662: the Snowball Nebula (or “Blue Snowball Nebula”)

Almost all of the emission is in the form of spectral lines: H and He recombination linesand forbidden lines of heavier elements.

Color image taken at Mees Observatory in AST 111. Visible spectrum from Pic du Midi Observatory.


http://www.astrosurf.com/buil/us/mission2/mission2.htm

Role of interstellar gas in the galaxy

Interstellar gas is the reservoir of material for star formation and star death.I Stars form by gravitational collapse of interstellar clouds.

I Dying stars return fusion-processed material to the ISM, enriching it in heavierelements and providing material for new stars.

Many properties of interstellar gas clouds are measurable with high precision: densitytemperature, pressure, elemental/molecular abundance, etc.I For the purpose of studying galactic structure, dynamics, and evolution, the gas in

the ISM is a useful complement to the information available from stars.


The Stellar Life Cycle


Milky Way stars, dust, and interstellar gasStarlight, extinction

Starlight

Dust (emission)

Molecular gas

Neutral atomic gas


The components of the ISM

Diffuse ISM neutral atomic clouds embedded inionized medium

Dense ISM neutral molecular clouds

Ionized nebulae HII regions, planetary nebulae,SN remnants

H II region NGC 604 in Triangulum (NASA/HST).


The components of the ISMDiffuse ISM

Clouds of mass 103 � 106M�, 10 � 100 pc in size, 10 � 104 K in temperature.I Mostly neutral atomic material embedded in a much less dense and very hot

(> 105 K) ionized medium.

I Mostly in the form of dark clouds with nH = 0.1 � 10 cm�3 and T = 10 � 100 K.

I Dark because the extinction by dust they contain can be as high as AV ⇠ 3.

I The neutral diffuse ISM fills 40%–80% of the volume of the Milky Way.

I There is about 109 � 1010M� of neutral diffuse ISM in the Galaxy!

I Spectral line tracers: HI 21 cm line; C+ line at 157.7 µm.


The components of the ISMDense ISM

H — neutral and mostly molecular instead of atomic — in the form of clouds withdensities nH2 = 10 � 106 cm�3. Temperature is 10 � 100 K, mass is 103 � 106M�.I As much mass (total in Galaxy) as the diffuse ISM, 109 � 1010M� but volume is small

in comparison.

I Molecular cloud complexes are usually physically connected to complexes of diffuseatomic clouds.

I The visual extinction through a molecular cloud is � 1.

I Best spectral line tracers are rotational lines of CO. H2 radiates too poorly and isexcited too inefficiently to be an effective tracer.

I 180 molecular species have been detected so far in interstellar clouds. Smallestmolecule: H2; largest: C70 fullerene.

I See Al Wootten’s list of reported ISM molecules at NRAO


http://www.cv.nrao.edu/~awootten/allmols.html

The components of the ISMIonized nebulae

Include HII regions, planetary nebulae, and supernova remnants.

I Hydrogen is fully ionized in ionized nebulae; other elements may be multiplyionized.

I These objects have negligible mass on the galactic scale but they are very bright atvisible and radio wavelengths.

I They are the most easily noticed components of the ISM.

I Spectral line tracers: hydrogen recombination lines, “forbidden” lines of relativelyabundant ions and atoms (C, N, O, etc.).

I Electron densities are usually around ne = 10 � 104 cm�3. Temperatures are about104 K in HII regions and planetary nebulae.


The components of the ISMIonized nebulae

Planetary nebulae consist of gas ejected and ionized by stars with core masses below theSAC mass that are becoming white dwarfs.

HII regions are associated with young, massive O-type stars in star forming regions.I They always seem to occur on the edges of giant molecular cloud

complexes. The Orion clouds are the nearest and best example.

Supernova remnants are what their name implies. Their ionization traces the advance ofthe blast wave into the ISM rather than photoionization by UV starlight.I The matter in SNRs tends to be in lower ionization states than in HII

regions and planetary nebulae.


Heating and cooling of the ISMHII regions and planetary nebulae

Heating Photoionization by starlight; UV light with E > 13.6 eV ionizes hydrogen,imparting kinetic energy to the electrons produced.

Cooling Collisional excitation followed by emission of forbidden lines by ions of C,N, and O (mostly). The more metals, the more efficiently the cooling.

EAtomicenergylevels

e -atominelasticcollision

hn


Heating and cooling of the ISMNeutral, Diffuse ISM

Heating UV light (hn = 5 � 13 eV) in background starlight, known as the interstellarradiation field (ISRF).I Heating occurs via the photoelectric effect on dust grains.I Carbon can be ionized by E = hn = 11 eV photons and thus is usually

singly ionized (C II) in the diffuse ISM even when hydrogen is neutraland atomic.

Cooling Excitation of excited carbon (CII) by collisions with H atoms and electrons,followed by radiation in the forbidden 157.7 µm line.


Heating and cooling of the ISMMolecular Clouds

Heating Not well understood, but we know it cannot be starlight directly. Bestcandidates:Turbulence driven by stellar winds and outflows.Cosmic rays ions accelerated to 1015 eV in supernova remnants.

Cooling Collisional excitation followed by radiation in the rotational lines of themore abundant molecules like CO, OH, H2O.

I Molecular clouds are frequently dense and cold enough to be unstable and collapseunder their own weight.

I As we will see, the gravitational instability of molecular clouds is the principalmeans by which stars are formed in the galaxy.


Molecular clouds: Internal structure

I Molecular clouds are clumpy; they tend to consist of denser clumps in a range ofsizes and masses, blending into a less-dense background.

I Molecular clouds are cold; they tend to have T < 20 K.

I Molecular clouds are turbulent: the random internal velocities are typically⇠ 1 km/s, much larger than the average molecular speeds in a quiescent gas inequilibrium at the same temperature.

I Molecular clouds (and clumps) generally rotate slowly.

I Molecular clouds are magnetized: they are threaded by the same magnetic fluxpresent when the material was diffuse and atomic. Now it has been compressed to amuch smaller size, and the fields are correspondingly larger.


Molecular clouds & star formation

As a result:I Clumps get massive enough and cold enough that the gas pressure cannot hold up

their weight, and they collapse.I As clumps collapse they heat in their cores.I If they are sufficiently massive and collapse to a small enough scale, temperatures can

reach the fusion ignition point and a star is formed.

I The clumps are constantly being rearranged, compressed, or distended byturbulence.

I Collapse does not happen with spherical symmetry; often collapse is easier along theaxis of rotation because of centrifugal forces.


Collapse of a clump: the Jeans massConsider a constant-density clump and its central pressure:

Pc ⇡GM2

R4 from weight

=rkTµ

from ideal gas law

Balancing the pressure gives

GM2

R4 = rGMR

=rkTµ

=) MR

=kTµG

and assuming the mass is uniformly distributed gives

r =3M

4pR3 =) R =

✓3M4pr

◆1/3

) M =kTµG

✓3

4pr

◆1/3M1/3 =) M =

✓kTµG

◆3/2 ✓ 34pr

◆1/2


Collapse of a clump: the Jeans massDefine the Jeans mass MJ as the critical mass for collapse:

M =

✓kTµG

◆3/2 ✓ 34pr

◆1/2

If a clump’s mass exceeds the Jeans mass for its density, temperature, and composition, itwill collapse under its weight.

Example: A pure molecular Hydrogen cloudFor a cloud of H2:

µ = 3.3 ⇥ 10�24 gT = 20 K

n = 1.4 ⇥ 105 cm�3

) MJ = 1M�


interstellar gas - university of rochester

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