Download Molecular spectroscopy in Astrophysics

Molecular spectroscopy in Astrophysics Serena Vi4 University College London Universitat de Barcelona, December 2010 Outline 1.  Basic principles of molecular spectroscopy and an introduc4on on molecules in space 2.  Interpreta4on of emission and absorp4on molecular spectra from space 3.  Applica4ons of molecular spectroscopy –  Low temperature regime (e.g. starless cores) –  Warm regimes (e.g. star forming regions) –  Hot regime (e.g. atmospheres of cool stars) 4.  Recent Highlights from Herschel 1. Basic principles of molecular spectroscopy •  For atoms, the energy levels corresponding to
bound electronic states are quantized
•  In the case of molecules we have an extra
degree of freedom: vibration and rotation
•  We will limit our discussion to diatomics but
polyatomics are similar
Example: interaction of 2 H atoms
+
(free)
Nuclei close (no
shielding) repulsion
Einteraction
separation
Long range interaction (Van
der Waals) between atoms
(bound)
When bonding occurs
energy gets lower
Shielding of p by e-
•  Bottom of curve is approximated by
harmonic oscillator
Evib = ωe (v + 1 / 2), v = 0,1,2.....
i.e quantized vibration
•  ωe=angular frequency of oscillations
•  ν = vibrational quantum number
•  n.b. zero point energy is 1/2 ω
•  In this model molecules have vibrational
ladder of equally spaced levels: ΔEvib=
(Δv=±1) (can also have Δv=±2)
Harmonic Oscillator
•  Molecules can also rotate:
–  Erot(classical) = ½ Iωrot2 where I is the moment
of inertia about the centre of mass
–  Erot(quantum) = BJ(J+1)
–  B =! 2 / 2I is the rotational constant
–  So B tends to decrease as mass increases
•  Since E rot = BJ(J+1) then for a 1-0
transition à ΔErot(1-0)=2B
•  In general: heavy species à bigger I à
smaller B à levels closer, λ larger
•  Total energy of a molecule depends on:
electronic + vibrational + rotational state
(few eV) (~0.1eV)
(~10-3eV)
Molecules in space Over 120 molecular
species have been
discovered in
space (gas and dust) (this is
a very outdated list!)
Where are molecules found?
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Interstellar clouds:
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Diffuse clouds
Dark clouds
Star-forming regions
Circum-stellar discs around young stars
Cool envelopes around giant stars
Stellar atmospheres of relatively cool stars
Sunspots
Planetary atmospheres
Comets and asteroids
Ejecta of novae and supernovae
Dark cloud: B68 Star forming regions Cygnus Wall Planetary Nebula: NGC 6369 Sunspots •  Stars are born from clouds
in the interstellar medium.
•  These clouds are dense
(104-105 cm-3) and cold
(10K) and hydrogen atoms
are therefore bound as
molecules.
•  Molecular hydrogen is the
raw material for building
new stars.
•  However, the cloud also
contains microscopic dust
grains.
•  Typically, dust to gas ratio
is 1/100.
•  Due to instabilities (slow
MHD, turbulence)some
cores within such clouds
become unstable à
collapse
•  After a star is born
the material (gas
and dust)
surrounding the
protostar will heat
up
•  Atoms and
molecules locked in
icy mantles on dust
grains will sublimate
•  Molecular material
around new born
stars is transient
(e.g. stellar winds
and outflows will
disperse it and
destroy it)
The multi-million year cycle
H2O, NH3 etc back in gas HCOOCH3, CH3CN, e
Collapse
Large molecules,
Aggregation (icy mantles)
C, O, N, etc à grains Cooling
Radiation
O + grains à H2O N + grains à NH3 etc Winds, shocks,
dust evaporation
Large hydrocarbons
e.g. C6H6 Roles of molecules: Observa4ons of molecules à local condi4ons in interstellar clouds 1.  They trace interstellar gas: a.  rela4vely cool (T ≤103 K) b.  rela4vely dense (nH > 100 cm-­‐3) (note: star forming clouds have a density of 105 – 107 cm-­‐3) c.  This gas contains nearly all non-­‐stellar baryonic mass in the Galaxy à molecules are tools to understand interstellar condi4ons 2.  main coolants in denser interstellar gas: 1. 
2. 
3. 
radiate at long wavelengths (typically mm); transi4ons corresponding to low temperatures molecular radia4on allows collapse under gravity of gas clouds in forma4on of galaxies, globular clusters and star forma4on (poten4al energy radiated away) 3.  Control level of ioniza4on: e.g. HCO+ + e-­‐ à H + CO fast , Mg+ + e-­‐ à Mg + hν slow Link magne4c field to gas: –  High ioniza4on à ions 4ed to field lines, many ion-­‐
neutral collisions so neutral and field fixed together so magne4c pressure. –  Low ioniza4on à neutral gas weakly 4ed to B-­‐field à no magne4c pressure 2. Interpreta4on of emission and absorp4on molecular spectra Detection of molecules
•  Diffuse clouds (nH ~ 100 cm-3, T~100K):
–  In general molecules are detected in absorption
against a hot star (effectively a continuum source)
Credit: Adapted from a diagram by James B. Kaler, in "Stars and their Spectra," Cambridge University Press, 1989 •  So, for example:
–  H2(v =0,J=0,1) + photon à H2 (v =1, J=0,1…)
•  Dark and dense clouds (nH ≥ 104 cm-3; T ~ 10K):
–  Molecules seen in emission
–  A molecule can only emit radiation from an upper level if it is
excited to that level. Usually this excitation occurs via collisions
with H2.
Example: CO
•  CO(J=0,1) + H2 à CO(J=2) + H2
•  CO(J=2) à CO(J=1) + photon
•  ΔE(1-0)=2B=5.6K
•  ΔE(2-1)=4B=11.2K
•  ΔE(2-0)=6B=16.8K
•  ΔJ=J (upper)-J (lower)
ΔJ:
-2 -1
0
+1
Branch: O P
Q
R
+2
S
Energy level diagram for the water molecule P
ee
(0,4)
O
ee
(0,3)
P
ee
(0,2)
O
ee
(0,1)
P
ee
(0,0)
O
eo
(1,4)
P
eo
(1,3)
O
eo
(1,2)
P
eo
(1,1)
P
fe
(1,3)
O
fe
(1,2)
P
fe
(1,1)
O
O
fo
(2,3)
P
fo
(2,2)
O
fo
(2,1)
P
ee
(2,2)
O
ee
(2,1)
P
O
eo
(3,2)
P
eo
(3,1)
P
fe
(3,1)
O
O
fo
(4,1)
P
ee
(4,0)
J
4
3
fe
(3,0)
2
ee
(2,0)
1
fe
(1,0)
0
Critical density
•  CO(J=2) à CO(J=1) + photon
This Emission is spontaneous (occurs at rate of Einstein
coeff) but collisional de-excitation can in fact occur:
A-
CO(J=2) + H2 ßàCO (J=1) + H2 à no radiation
•  For any molecular transition there is a critical density:
Aul
n* =
!"
at which the radia4ng molecule suffers collisions at the rate n(H2)σv equals A •  So, the critical density is defined as the density at
which the rate of upward transitions through
collisions is equal to the rate of downward
transitions through spontaneous decay
•  Unless density is a significant fraction of
critical density emission line is invisible
•  Each transition of any molecule will have a
different critical density à each transition is
potentially tracing a different gas component
For CO (J=1-­‐0) ncrit~103cm-­‐3 While CS has ncrit~105cm-­‐3
à can use different molecules/transi4ons to map different parts of clouds: iden4fy clumps where stars may form. •  Emission is usually op4cally thin, so more intensity means more molecules. •  Line profile can be broadened by bulk
motions:
Δv / c = Δν / ν
in some molecular clouds (e.g. B5) , Δν ~ 10 kms-1 cf. with Δν(thermal) ~ 0.3 kms-1
Radia4on can be shined to a higher or lower frequency: Doppler shin Velocity of the object rela4ve to the observer (+ when moving towards us, -­‐ve when moving away from us) 1
Observed frequency υ = υ+
Rest frequency υv
c
Table 1: Cosmic (i.e. solar) rela4ve abundances H
He
O
1
0.075-0.14
2-4(-4)
C
N
Si, Mg
Fe
S
Na, Al, Ca, Ni
1-2(-4)
7(-5)
3(-5)
2(-5)
1(-5)
2(-6)
How do molecules form and get destroyed in
space?
•  To account for the molecular richness that
we see in dense clouds we need an
efficient formation mechanism i.e it occurs
at nearly every collisions
•  Ion-molecule reactions probably the
most important
•  Much of the chemistry starts with the
formation of H3+ via H2:
–  H+3 created by c.r. ionization (n.b. UV cannot
create H3+ - c.r. are fast, MeV, protons):
H 2 + c.r. → H + e
+
2
−
H + H2 → H + H
+
2
+
3
•  Abundant species such as oxygen
preferentially react with H3+ because they
do not react with H2 at low temperatures
(≤ 100K) – the energy barrier is ~ 4000 K
O + H 3+ → OH + + H 2
OH + + H 2 → H 2O + + H
H 2 O + + H 2 → H 3O + + H
⎧OH + 2 H
H 3O + e → ⎨
⎩ H 2O + H
+
−
All very efficient
•  These reactions supply the OH and H2O
•  Other species, such as some ions, react
directly with H2:
e.g:
H2 + O+ à OH+ + H
Then chemistry goes on as in previous slide
•  Another important example is the
formation of hydrocarbons:
C+ +H2
CH+ + H
Because it is too slow so:
+
3
+
C + H ! CH + H 2
+
+
2
+
2
+
3
CH + H 2 ! CH + H
CH + H 2 ! CH + H...etc
#CH + 2H
+
"
CH 3 + e ! $
%CH 2 + H
•  So from one molecule, H2, a range of
molecules can be formed. But these are all
hydrides. To make molecules with 2 heavy
atoms:
•  Example:
– C+OH à CO + H (neutral exchange)
•  Or:
– C+ + OH àCO+ + H
ion-mol
– CO+ + H2 à HCO+ + H ion-mol
– HCO+ + e- à CO + H
diss rec.
•  Schemes containing thousands of reactions
are now routinely studied.
So, in summary:
•  Ion-molecule:
A+ + BC à AB+ + C
•  Diss.rec. and rad. Reac: AB+ + e- à A + B
AB + hv
•  Neutral exchange:
A + BC à AB + C
So, at least some of the dark cloud chemistry is
driven by cosmic rays: it creates ions, drives ionmol reactions, produce molecular ions which
recombine (dissociatively) to form new neutral
species.
More detailed summary of chemical
reactions occurring in the ISM
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In addition molecules can be formed on the surface of the grains (this is
important in dense environments such as in hot cores – dense
cores, remnant of high mass star formation):
•  However remember that everything starts
with H2 e.g recall that:
H 2 + c.r. → H + e
+
2
−
H + H2 → H + H
+
2
+
3
But how easy is it to make H2?