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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
3.7 Isotope Effect
In this section we will study a first astrophysical application of molecular spectroscopy,
namely the determination of abundance ratios of different isotopes. Molecules serve as a
much more ideal tool to distinguish different isotopes than atoms. This is because the
presence of more or less neutrons in the nucleus of a specific chemical element does not
strongly modify the electric field and thus has only a small influence on the electron
configuration. The same is true for the electron configuration of molecules of course.
However, the energy of nuclear motion in molecules due to vibration and rotation is
influenced already to first order if the number of neutrons and thus the mass of the nucleus is
changed. Thus, isotopic molecules have different frequencies of vibrations and rotations.
First we will consider the modification of molecular spectra due to the presence of
different isotopes. Then we will look at several examples that illustrate how knowledge of
isotope ratios allows us to gain insight into astrophysical objects and processes. Later, in
Chapter 6 (Astrobiology), we will discuss an additional important example where knowledge
of isotope ratios is crucial for determining the origin of water on Earth.
3.7.1
Vibration
For isotopic molecules, i.e. molecules that differ only by the mass of one or both of the nuclei
but not by their atomic number (for example 1H35Cl and 1H37Cl), the vibrational frequencies
are obviously different. Assuming harmonic vibrations the (classical) vibrational frequency is
given by (see Eq. (3.33))
ν osc =
1
2π
k
µ
,
(3.88)
where the force constant k, since it is determined by the electronic motion only, is exactly the
same for different isotopic molecules, whereas the reduced mass µ is different. Therefore, if
we let the superscript “i” distinguish an isotopic molecule from the “ordinary” molecule we
have
i
ν osc
µ
=
=ρ .
ν osc
µi
(3.89)
The heavier isotope has the smaller frequency. If the superscript “i” refers to the heavier
isotope the constant ρ will be smaller than 1. For example, the values of ρ for the pairs 1H35Cl
and 1H37Cl, 1H35Cl and 2H35Cl, 10BO and 11BO, and 16O16O and 16O18O are 0.99924, 0.71720,
0.97177, and 0.97176, respectively.
By substituting Eq. (3.89) into Eq. (3.38) we find for the vibrational levels of two
isotopic molecules (still assuming harmonic oscillations)
1
1


G i (υ ) = ωei υ +  =ρωe υ +  = ρ G (υ ) .
2
2


(3.90)
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Therefore, the separation of corresponding vibrational levels in two isotopic molecules are
somewhat shifted (and as a result spectral lines are also shifted). The levels of the lighter
isotope always lie higher than those of the heavier isotope.
If anharmonicity is taken into account, the calculations become rather more involved and
will not be reproduced here. The formulae for the energy levels are found to be in a very good
approximation
2
3
1
1
1



G (υ ) = ωe υ +  − ωe xe υ +  + ωe ye υ +  + … ,
2
2
2



2
(3.91)
3
1
1
1



G i (υ ) = ρωe υ +  − ρ 2ωe xe υ +  + ρ 3ωe ye υ +  + … .
2
2
2



(3.92)
In other words, the vibrational constants are modified as
ωei = ρωe ,
ωei xei = ρ 2ωe xe ,
(3.93)
ωei yei = ρ 3ωe ye .
3.7.2
Rotation
Since the reduced mass is inversely proportional to the rotational constant B, molecules
containing heavy isotopes have rotational lines corresponding to lower quantum energies and
smaller line spacing. Specifically, we find for the
rotational constant B of the two isotopic molecules
ℏ2
B =
=ρ 2 Be .
i 2
2 µ req hc
i
e
(3.94)
Note that the internuclear distances in diatomic
(and polyatomic) molecules are entirely determined
by the electronic structure. They are therefore
exactly equal in isotopic molecules as long as no
vibration occurs. The rotational energies of the two
isotopic molecules are thus connected by
F i = Bei J ( J + 1) = ρ 2 Be J ( J + 1) = ρ 2 F .
(3.95)
Rotational levels of the heavier molecule have
smaller energies. Furthermore, the separation of
neighboring lines in the rotational spectrum (which
is 2B in first approximation) differs for isotopic
molecules. For example, for the 12CO molecule, 2B
is found to be 3.842 cm–1, and for the 13CO
molecule containing the heavier isotope of carbon,
Figure 3.43: The isotope effect on the rotational energy levels and the corresponding
rotational spectrum of the CO molecule. The
line shifts are exaggerated in this drawing.
From Haken & Wolf (2006).
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
2B is found to be 3.673 cm–1. Figure Figure 3.43 shows the resulting differences in the
rotational spectra of CO containing the isotopes 12C and 13C.
For simultaneous vibration and rotation, in a first approximation, we simply have to add
the vibrational and rotational isotope effects. As a result, the lines of a rotation-vibration band
of an isotopic molecule do not have exactly the same separations as the lines of the “normal”
molecule. In other words, the isotope displacement between corresponding lines of the two
bands is dependent on J.
Since the rotational energies are smaller than the vibrational energies, the rotational
isotope shift is, in general, smaller than the vibrational one, despite the fact that the former
scales with ρ2 while the latter scales with ρ.
For more precise calculation, it is necessary to take into account the interaction of
vibration and rotation, i.e. the rotational constants α, β, D, and possibly higher order ones. To
a very good approximation we may use
Dei = ρ 4 De ,
α i = ρ 3α ,
(3.96)
β i = ρ 5β .
3.7.3
Astrophysical Examples
Nucleosynthesis: Overview
The detection of isotopic molecules enables measurements of isotope ratios, which provide
information on details of the nucleosynthesis:
•
•
•
•
•
•
7
Li/6Li
12 13
C/ C
16 17
O/ O/18O
24
Mg/25Mg/26Mg
32 34
S/ S
56
Fe/57Fe/58Fe
⇒
⇒
⇒
⇒
⇒
⇒
lithium production rate
CNO cycle efficiency and mixing processes in stellar envelopes
3α-process efficiency
carbon burning efficiency
oxygen burning efficiency
supernova explosion details
In the following, as an example, we will have a closer look at the astrophysical diagnostics
with the 12C/13C ratio. Apart from studying nucleosynthesis at different stages of the stellar
evolution, measurements of isotope ratios serve as a test of primordial nucleosynthesis and
thus of cosmology. Isotope ratios also provide insight into the chemical evolution of our
galaxy (see below).
Solar Isotope Ratios
Figure 3.44 shows infrared bands with ∆υ = 1 of isotopic CO molecules in the solar spectrum.
Comparison of such observations with synthetic, i.e. computed, spectra allows us to determine
isotope ratios for carbon and oxygen atoms. The analysis yields solar isotope ratios 12C/13C =
80±1, 16O/17O = 1700±220, and 16O/18O = 440±6 (Ayres et al. 2006).
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Figure 3.44: Solar spectrum near 4.6 µm from the NSO (National Solar Observatory) solar atlas. The four
different spectra are shifted for easier comparison and represent the photosphere, hot umbra, medium umbra, and
cold umbra, respectively. Telluric lines, i.e. lines caused by Earth’s atmosphere, have been removed or corrected
for. The spectrum shows an extract of the rotational-vibrational bands with ∆υ = 1 in the electronic ground state
(X1Σ+) of different CO isotopic molecules.
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
Since the solar atmosphere is not simply an optically thin gas, the calculation of the solar
spectrum involves radiative transfer calculations based on solar model atmospheres. The
model atmosphere provides the height stratification of the thermodynamic properties such as
temperature, pressure, and density within the solar atmosphere.
Stellar 12C/13C ratios
Analogously to the Sun we can determine isotope ratios of other stars. The 12C/13C isotope
ratio is interesting because it is expected to change during the evolution of a star and, as a
result, also during the evolution of a galaxy.
Initially, the chemical composition of a star corresponds to the local condition of the
interstellar medium (ISM) and the cloud from which it was formed. Fusion in the stellar
interior modifies the chemical compositions of a star during its lifetime. In the stellar
spectrum this chemical evolution is only apparent if a mixing of the surface layers takes place
with the center or shells where fusion occurs. Only stars of spectral class F and later have an
outer convection zone, which extends all the way to the center only for late M stars. For
example, the Sun, a G2 V star, possesses an outer convection zone that extends to about 0.7
R (as measured from the center). Therefore, in the case of the Sun, the 12C/13C ratio at the
surface is still the same as for the young Sun despite modifications resulting from the CNOcycle that runs at the solar center. The CNO-cycle reduces 12C, increases 13C, and increases
14
N abundances, so that the 12C/13C ratio actually reduces with time. Only in some phases of
the late stages of stellar evolution the outer convection zone of a star like the Sun reaches
down to layers where hydrogen burning has taken place so that products of the CNO-cycle are
Figure 3.45: Comparison of the spectrum of the giant IV-101 (solid) with the average spectrum of three other
giants (dashed), all belonging to the globular cluster M3. The spectrum contains the infrared (3,1) band head of
13
CO at 2374 Å (lines of this band are labeled “13”) and lines belonging to the (2,0) band of 12CO (labeled “12”).
From Pilachowski et al. (2003).
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transported to the surface. This occurs a first time during the so called first dredge-up
dredge
while
the star moves up the red giant branch in the Hertzsprung-Russell
Hertzsprung Russell diagram. Thus, we expect
12 13
much lower C/ C ratios in the spectrum of red giants than in the solar spectrum.
spectrum
Figure 3.45 shows measured spectra containing
c
infrared 12CO and 13CO lines of red
giants in the globular cluster M3. The weakness of the 13CO band head in the M3 giant called
IV-101 compared with threee other giants is apparent even from visual inspection. The
strengths of the 12CO lines, which
which compose most of the remaining spectral features, are
similar in all four stars. Since the stars all have similar temperature, gravities, and
metallicities, including oxygen and carbon abundances, the apparent weakness of the 13CO
lines in IV-101 is a clear
ear indication of a high 12C/13C ratio in this star compared with the other
giants in the observed cluster. The 13CO strength in IV-101
101 is about a factor of 2 less than in
12 13
the other stars. A ratio C/ C = 11 ± 1.5 is found for IV-101, while 12C/13C = 6 ± 1.5 for the
three comparison stars (Pilachowski et al. 2003). The low value in the comparison giants
indicates a very strong mixing. On the other hand, IV-101 is known as Li--rich K giant (while
the other giants exhibit no Li enrichment). The higher Li abundance
abundance at the surface may be due
to the presence of a Li burning shell that influences the interior temperature structure and thus
the depth of the outer convection zone, apparently leading to less mixing. Such observations
allow us to test and improve stellar
ste
evolution models.
Figure 3.46 illustrates how the fitting of synthetic spectra to observations may be done, in
this case for an MgH band resulting from an electronic transition. From the best fit the isotope
ratio 24Mg:25Mg:26Mg = 78:13:9 was determined (Gay & Lambert 2000). Synthesis of
magnesium occurs primarily in the carbon
carb and neon-burning
burning shells of massive stars prior to
Figure 3.46: Spectrum of the star HD 23439A from 5134 Å to 5136 Å. The spectrum shows a portion of the
MgH A2Π–X2Σ+ ∆υ = 0 band. The observed spectrum (circles) is compared with synthetic
synth
spectra for the
isotopic ratios 24Mg:25Mg:26Mg = 100:0:0 (dashed), 78:13:9 (solid, best fit), and 72:16:12 and 83:10:6 (dotted).
From Gay & Lambert (2000).
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
their deaths as type II supernovae. Therefore, measurements of the magnesium isotope ratio in
stars provide information about the evolution of AGB stars, i.e. stars in the asymptotic giant
branch.
Galactic Chemical Evolution
The solar isotope ratios, e.g. 12C/13C = 80, corresponds to the chemical composition of our
galaxy 4.6 × 109 years ago at the birthplace of the Sun. A comparison with today’s isotope
ratios in the interstellar medium provides important clues on the production rate of heavier
elements in stars and allows us to determine empirically the chemical evolution of the
interstellar medium (ISM) within the galaxy during the last 4.6 × 109 years.
Conventionally, the local ISM has been used for such a comparison with the Sun, i.e. it
was implicitly assumed that the Sun had been formed at the same galactocentric distance R0 =
8.5 kpc which the Sun has today. However, the solar anomaly in metallicity (it is somewhat
higher) compared with the mean of nearby stars and with the galactic gradient of the iron abundance has given evidence that the Sun has formed at an initial galactocentric distance Ri, =
6.6 kpc (Wielen et al. 1996). Such a displacement is in good agreement with predictions on
the diffusion of stellar orbits within the galaxy.
Figure 3.47 shows the comparison of the 12C/13C and 16O/18O isotope ratios of the Sun
with today’s ISM. As expected the 12C/13C ratio reduces significantly with time. According to
theoretical studies the 16O/18O ratio should also decrease with time. Assuming a birthplace of
the Sun closer to the galactic center is thus in better agreement with theoretical expectations
on the chemical evolution, so that the isotopic study is a further argument in favor of a shift of
the Sun in galactocentric distance.
Figure 3.47: The isotope ratios of 12C/13C (left) and 16O/18O (right) of the interstellar medium (open circles) as a
function of the galactocentric distance R. The time evolution of the isotope ratios is shown by a comparison with
the solar system value, indicated by . In the previous, conventional comparison, the Sun is compared with the
local interstellar medium at R0 = 8.5 kpc. In the improved analysis, the solar value of the isotope ratio is
compared with the interstellar medium at the birth-place of the Sun at Ri, = 6.6 kpc. From Wielen & Wilson
(1997).
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3.8 Observation of Molecular Spectra
This section is devoted to a brief introduction into infrared and radio astronomy, including
important current and future observatories and satellite missions. As we have learned previously, the infrared to radio spectral regions allow us to observe rotational and rotationalvibrational spectra of molecules. In addition, the same spectral domains cover continuum
spectra due to thermal and non-thermal processes. We will leave near-UV and optical observations aside here as they apply also to the detection of atomic spectra and are better known.
3.8.1
Atmospheric Transmission
Earth’s atmosphere is transparent only in selected spectral windows (Figure 3.48). It is
transparent only in the optical, in selected windows in the near infrared and in a broad radio
wavelength region. Most of the infrared light reaching us from space is absorbed by
molecular bands due to water vapor and carbon dioxide in the Earth's atmosphere. Radiation
with shorter wavelengths than optical light is blocked by ozone, molecular oxygen, and
molecular nitrogen. Ground-based observations are only possible within the transmission
windows while all remaining spectral regions must be observed from space.
3.8.2
Infrared Observations
The infrared is usually divided into three spectral regions: near-infrared, mid-infrared and farinfrared (Table 3.2). However, the boundaries between these regions are not strictly agreed
upon and can vary for different authors.
Figure 3.48: Opacity of Earth’s atmosphere as a function of wavelength.
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
Table 3.2: Infrared spectral regions and astrophysical objects that can be observed in these regions.
Spectral region
Wavelength range
[µm]
Temperature range
[K]
Objects
Near-infrared
(0.7–1) to 5
740 to
(3,000–5,200)
Cooler red stars, red giants;
dust is transparent
Mid-infrared
5 to (25–40)
(92.5–140) to 740
Planets, comets and asteroids,
protoplanetary disks,
dust warmed by starlight
Far-infrared
(25–40) to
(200–350)
(10.6–18.5) to
(92.5–140)
Emission from cold dust,
central regions of galaxies,
very cold molecular clouds
Table 3.3: Atmospheric infrared windows.
Wavelength
range [µm]
Band
Sky transparency
Sky brightness
1.1–1.4
J
high
low at night
1.5–1.8
H
high
very low
2.0–2.4
K
high
very low
3.0–4.0
L
3.0–3.5 µm: fair
3.5–4.0 µm: high
low
4.6–5.0
M
low
high
7.5–14.5
N
8–9 µm and 10–12 µm: fair
rest: low
very high
17–40
17–25: Q
28–40: Z
very low
very high
very low
low
330–370
The transmission windows of Earth’s atmosphere in the infrared are labeled to simplify a
reference to a certain wavelength region (Table 3.3). This is analogous to wavelength bands
characterized by standard filters employed in the optical (Johnson UBVRI filters).
Near-Infrared
Near-infrared observations have been made from ground-based observatories since the
1960’s. As we enter the near-infrared region, the hot blue stars seen clearly in visible light
fade out and cooler stars come into view. Red giants and low mass red dwarfs dominate in the
near-infrared. The near-infrared is also the region where interstellar dust is the most
transparent to infrared light.
Mid-Infrared
Mid and far-infrared observations can only be made by instruments brought above the
atmosphere. Entering the mid-infrared region of the spectrum from the small wavelength side,
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the cool stars begin to fade out and cooler objects such as planets,
comets and asteroids come into view.
Planets absorb light from the Sun and heat up. They then reradiate this heat as infrared light, in first approximation as a
black-body spectrum. This radiation has to be distinguished from
the visible light emitted by planets, which is reflected light from
the Sun (or from the central star in the case of an extrasolar
planet). The planets in our solar system have temperatures
Figure 3.49: Earth’s infraranging from about 50 to 570 K and emit most of their light in the
red emission.
mid-infrared. For example, the Earth itself radiates most strongly
at about 10 µm (Figure 3.49). Asteroids also emit most of
their light in the mid-infrared making this wavelength
band the most efficient for locating dark asteroids.
Dust warmed by starlight is also very prominent in
the mid-infrared. An example is the zodiacal dust which
lies in the plane of our solar system. This dust is made up
of silicates and ranges in size from a tenth of a µm up to
Figure 3.50: Mid-infrared image (at
the size of large rocks. Silicates emit most of their radia25 µm) of comet Iras-Araki-Alcock
tion at about 10 µm. The dust from comets also has strong
observed with IRAS.
emission in the mid-infrared (Figure 3.50). Protoplanetary disks also shine brightly in the mid-infrared.
Far-Infrared
In the far-infrared, the stars have all vanished. Instead we
now see very cold matter (140 K or less). Huge, cold
clouds of gas and dust in our own galaxy, as well as in
nearby galaxies, glow in far-infrared light (Figure 3.51).
In some of these clouds, new stars are just beginning to
form. Far-infrared observations can detect these protostars long before they are seen in the visible by sensing
the heat they radiate as they contract.
The central region of most galaxies shines brightly in
the far-infrared because of the thick concentration of stars
embedded in dense clouds of dust. These stars heat up the
dust and cause it to glow brightly in the infrared.
Some galaxies, called starburst galaxies, have an
extremely high number of newly forming stars
heating interstellar dust clouds. These galaxies by
far outshine all other galaxies in the far-infrared.
Figure 3.51: Composite image of the
giant molecular cloud in Orion observed with IRAS in the far-infrared
at 12 µm, 60 µm, and 100 µm. The
Orion star constellation is outlined.
Ground Based Infrared Observations
Infrared detectors attached to ground-based telescopes can detect the near-infrared wavelengths
Figure 3.52: Composite image of our
galaxy in the far-infrared at 60 µm, 100 µm,
and 240 µm observed with COBE.
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
which make it through our atmosphere. The best location for ground based infrared observatories is on a high, dry mountain, above much of the water vapor which absorbs infrared radiation. At these high altitudes, we can study infrared bands up to about 35 µm.
The first infrared sky survey (at 2.2 µm) was conducted at the Mount Wilson Observatory
(California, next to Los Angeles), where also data for the first infrared star catalogue was
obtained. Nowadays, the largest group of infrared telescopes can be found on top of Mauna
Kea (a dormant volcano) on Big Island, Hawaii. At an elevation of 4200 m, the Mauna Kea
Observatory, which was founded in 1967, is well above much of the infrared absorbing water
vapor.
High-Atmosphere Infrared Astronomy
It is best to get above as much of the atmosphere as possible to observe in the infrared. To do
so, infrared detectors have been placed on balloons, rockets and airplanes, allowing us to
study also mid- and far-infrared wavelengths. For example, infrared telescopes onboard
aircraft such as the Kuiper Airborne Observatory were used to discover the rings of Uranus in
1977.
NASA and the German Aerospace Center (DLR) are currently developing jointly a new
airborne observatory, the Stratospheric Observatory For Infrared Astronomy (SOFIA), an
optical/infrared/sub-millimeter 2.7 m telescope mounted in a Boeing 747. It will observe in
the range from 0.3 to 655 µm. In 2009 the aircraft is undergoing test flights at high altitude
(the intended observing altitude is 12 km). The start of science operation is scheduled to 2011,
though full capability will only be reached by 2014.
Infrared Astronomy from Space
Infrared space telescopes are not restricted by the transmission windows of Earth’s
atmosphere. Furthermore, they can view a large area of the sky and observe regions for longer
time periods than is possible with telescopes onboard balloons or aircrafts.
Infrared detectors have to be kept in an environment which is as cold as possible. The
colder the environment, the more sensitive the instruments are to infrared light. Any
surrounding heat will create infrared signals that can interfere with signals from space. This
includes heat radiation from the instruments and telescope, from the atmosphere (for groundbased observatories), and from warmer objects in space like the Sun. To keep infrared
detectors and instruments cold, cryogens such as liquid helium are used. The cryogen is
usually contained in a chamber called a cryostat. Solar shields are also placed on infrared
space telescopes to protect the instruments from the Sun’s heat.
The lifetime of an infrared space mission depends on how long the instruments can
remain cold. Shortly after the cryogen (or coolant) runs out, the detectors will become useless
and the data gathering portion of the mission will end. The relatively short lifetime of a few
years is a major disadvantage of space missions.
In the following we list the most important infrared space telescopes, covering past and
current missions as well as currently planned, future telescopes.
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Infrared Astronomical Satellite (IRAS):
• United States (NASA), Netherlands (NIVR), and
Great Britain (SERC)
• Launched January 1983; duration: ten months
• Earth orbit (height 900 km)
• 0.6 m telescope
• Wavelengths: 12, 25, 60 and 100 µm
• Scanned more than 96% of the sky four times
• First high sensitivity all-sky maps
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Figure 3.53: IRAS.
Discoveries:
• Detected ~350,000 infrared sources
• Six new comets
• Disk of dust grains around Vega
• Warm dust (infrared cirrus) in almost every direction of space
• Revealed for the first time the core of our galaxy
• Very strong infrared emission from interacting galaxies
Infrared Space Observatory (ISO):
• ESA
• November 1995 to April 1998
• Highly elliptical Earth orbit
• 0.6 m telescope
• Wavelengths: 2.5–240 µm (imaging, photometry, and
spectroscopy)
• Much more sensitive than IRAS and higher resolution
• Detected dry ice in interstellar dust and hydrocarbons in
some nebulae
• Detected several protoplanetary disks
Spitzer Space Telescope:
• NASA
• Launched August 2003
• May 2009: liquid helium exhausted (temperature rose
from 2 K to 31 K) ⇒ only two channels of the camera
(3.6 µm and 4.5 µm) still operational (until 2014?)
• Earth-trailing solar orbit
• 0.85 m telescope
• Wavelengths: 3–180 µm (imaging, imaging photometry
and spectroscopy)
Figure 3.54: ISO.
Figure 3.55: Spitzer Space
Telescope.
Discoveries:
• Detection of protoplanetary disks
• Observation of very young stellar objects
• Direct (but unresolved) observation of extrasolar planets (day vs. night time
temperatures; temperature map of HD 189733b)
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
•
First detection of water vapor and methane in the atmosphere of an extrasolar planet
(HD 189733b)
Discovered light from the possibly very first generation of stars in the universe
(redshifted from the UV or visible to the infrared), consistent with clustered first
stars
October 2009: discovery of tenuous, huge ring around Saturn (extending from 128 to
207 times Saturn’s radius)
•
•
Herschel Space Observatory:
• ESA
• Launched May 2009, intended minimum duration of 3
years
• Positioned in an orbit about Lagrange point L2 of the
Sun-Earth system (1.5 Mio km from Earth, further away
from the Sun than Earth)
• Imaging photometer, high and low resolution spectrometry in the far-infrared to sub-mm (55-672 µm)
Figure 3.56: Herschel Space
• The Herschel Space Observatory with the 3.5 m
Observatory.
telescope will perform spectroscopy and photometry at
55-672 µm. It will be used to study galaxy formation,
the interstellar medium, star formation, and the atmospheres of comets and planets in
the solar system.
James Webb Space Telescope:
• NASA, ESA, Canadian Space Agency
(CSA)
• Scheduled launch 2014, planned duration
10 years
• Will be positioned in an orbit about
Lagrange point L2 of the Sun-Earth
system
• Imaging and spectroscopy
• The James Webb Space Telescope (6.5 m)
Figure 3.57: James Webb Space Telescope.
is an infrared space mission to be operated
between 0.6 and 28 µm. It will have
extremely good sensitivity and resolution, giving us the best views yet of the sky in
the near- to mid-infrared. It will be used to study the first stars and galaxies in the
Universe, the formation and evolution of galaxies, the formation of stars and
planetary systems, and to study planetary systems and the origin of life.
3.8.3
Radio Astronomy
Production of Radio Waves
There exist many processes that produce radio waves in astrophysical environments. These
include:
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•
•
•
•
3–58
Spectral line radiation from atomic or molecular transitions that occur in the
interstellar medium or in the gaseous envelopes around stars.
Bremsstrahlung or thermal radiation from hot gas in the interstellar medium.
Synchrotron radiation from relativistic electrons in weak magnetic fields.
Pulsed radiation resulting from the rapid rotation of neutron stars surrounded by an
intense magnetic field and energetic electrons.
Interestingly, radio astronomy is not only concerned with cold environments. Synchrotron
radiation brings in an immediate cross-disciplinary contact between radio and high-energy
astrophysics, a connection that might have been thought unlikely because of the low energies
of radio photons. In fact, there is a high relevance of radio observations to high-energy
phenomena since the radio and X-ray observations of active galactic nuclei and quasars have
close relationships to one another.
Discoveries
Radio waves penetrate much of the gas and dust in space as well as the clouds of planetary
atmospheres and pass through the terrestrial atmosphere with little distortion. Apart from
different physical processes causing radio waves, this makes radio astronomy so valuable and
complementary to optical astronomy since is allows us to observe objects that remain hidden
to visible wavelengths. Nevertheless, optical observations are essential to understand what
types of objects are the sources of radio emission, and to put radio observations into a
physical and astrophysical context. Not surprisingly, this has led to important advances. We
just mention here the most famous discoveries of radio astronomy:
•
•
•
•
•
•
1933: Karl Jansky first detected cosmic radio noise from the center of our galaxy
(resulting from synchrotron radiation) while investigating radio disturbances
interfering with the transoceanic telephone service.
1940/50’s: Australian and British radio scientists located a number of discrete radio
sources associated with old supernovae and active galaxies, which later became
known as radio galaxies.
1963: Maarten Schmidt discovered quasars.
1965: Arno A. Penzias and Robert W. Wilson detected the cosmic microwave background radiation at a temperature of 3 K.
1967: Jocelyn Bell and Antony Hewish discovered pulsars, rapidly rotating magnetic
neutron stars.
More than 100 different molecules are detected in the space, including familiar
chemical compounds like water vapor, formaldehyde, ammonia, methanol, ethanol,
and carbon dioxide.
Radio Bands
Radio waves have wavelengths longer than about 1 mm (300 GHz). For wavelengths in the
range 1 cm – 20 cm Earth’s atmosphere, in particular the ionosphere, introduce only minor
distortion to incident radiation. At wavelengths below about 1 cm (30 GHz) absorption in the
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
atmosphere becomes increasingly critical, and observations from the ground are possible only
in a few specific wavelength bands that are relatively free of atmospheric absorption. At
wavelengths longer than 20 cm irregularities in the ionosphere distort incoming signals. This
causes a phenomenon known as scintillation, which is analogous to the twinkling of stars seen
at optical wavelengths. The absorption of cosmic radio waves by the ionosphere becomes
more important as the wavelength increases. At wavelengths longer than about 10 m (30
MHz), the ionosphere becomes opaque to incident signals. Therefore, radio observations of
cosmic sources at these wavelengths are difficult from ground-based radio telescopes.
Radio Telescopes
The instrumental methods of the radio astronomer often appear to be quite different from
those of the optical astronomer. The distinguishing feature of a radio telescope is that the
radiation energy gathered by the parabolic antenna is not measured immediately, a process
known as detection in radio terminology. Instead, the radiation is amplified and manipulated
coherently, preserving its wave-like character, before it is finally detected. It is crucial to cool
the amplifier cryogenically to reduce the internal noise as much as possible and to increase
the sensitivity of the instrument. The instrumental goals of the radio astronomer, i.e. obtaining
a larger collecting area, greater angular resolution, and more sensitive detectors, are otherwise
the same as they are for all astrophysical disciplines.
Radio telescopes are used to measure broad-bandwidth continuum radiation as well as
spectroscopic features due to atomic and molecular lines found in the radio spectrum of
astrophysical objects. In early radio telescopes, spectroscopic observations were made by
tuning a receiver across a sufficiently large frequency range to cover the various frequencies
of interest. This procedure, however, was extremely time-consuming and greatly restricted
observations. Modern radio telescopes observe simultaneously at a large number of
frequencies by dividing the signals up into as many as several thousand separate frequency
channels that may range over a total bandwidth of tens to hundreds of megahertz.
The most straightforward type of radio spectrometer employs a large number of filters,
each tuned to a separate frequency and followed by a separate detector to produce a multichannel, or multi-frequency, receiver. Alternatively, a single broad-bandwidth signal may be
converted into digital form and analyzed by the mathematical process of autocorrelation and
Fourier transformation. In order to detect faint signals, the receiver output is often averaged
over periods of up to several hours to reduce the effect of noise generated in the receiver.
Radio Interferometry
The angular resolution θ of a telescope depends on the wavelength of observations λ and the
size of the dish D and can be approximated by
θ=
λ
D
.
(3.97)
Because radio telescopes operate at much longer wavelengths than optical telescopes, radio
telescopes must be much larger than optical telescopes to achieve the same angular resolution.
Yet, even the largest antennas, when used at their shortest operating wavelength, have an
3 Molecular Spectroscopy
3–60
angular resolution only a little better than one arc minute, which is comparable to that of the
unaided human eye at optical wavelengths.
The high angular resolution of radio telescopes is achieved by using the principles of
interferometry to synthesize a very large effective aperture from a number of small elements.
In a simple two-element radio interferometer, the signals from an unresolved, or point, source
alternately arrive in phase and out of phase as the Earth rotates and causes a change in the
difference in path from the radio source to the two elements of the interferometer. This
produces interference fringes in a manner similar to that in an optical interferometer. In a
simple two-element radio interferometer, the angular resolution is defined by the distance
between the two telescopes, i.e. the baseline D. If the radio source has finite angular size, then
the difference in path length to the elements of the interferometer varies across the source.
The measured interference fringes from each interferometer pair thus depend on the detailed
nature of the radio brightness distribution in the sky.
Each interferometer pair measures one Fourier component of the brightness distribution
of the radio source. Movable antenna elements combined with the rotation of the Earth can
sample a sufficient number of Fourier components and thereby reconstruct high-resolution
images of the radio sky.
In the following we list the most famous radio interferometers, the first two being
operational while the last one is currently under construction:
Very Large Array (VLA): The VLA is located
west of Socorro, New Mexico, USA. It consists
of 27 antennas with 25 meter diameter each. The
antennas are distributed over 3 arms, each with a
length of 21 km and with 9 movable antennas
(Figure 3.58). The maximum baseline is 36 km,
and the corresponding smallest angular resolution is about 0.04 arcsec reached at the highest
frequency. The VLA operates in several different
frequency bands from 74 MHz to 43 GHz (400
cm to 0.7 cm).
Figure 3.58: Very Large Array (VLA).
Very Long Baseline Array (VLBA): The VLBA consists of 10 antennas with 25 m diameter
each, scattered around the US territory (Figure 3.59). The longest baseline in the array is 8611
km. The telescopes are capable of observing in 10 frequency bands ranging from 300 MHz to
86 GHz (100 cm to 0.4 cm). The maximum angular resolution is 0.13 milli-arcsec. For comparison, the diameter of the red supergiant Betelgeuse in Orion (as measured in the mid-infrared) is 55 milli-arcsec. However, only about a dozen stars can be resolved from Earth with
current instruments, and only two stars have a larger apparent angular diameter than
Betelgeuse, namely R Doradus and, of course, the Sun.
Atacama Large Millimeter Array (ALMA): ALMA is an interferometer currently under
construction, expected to start full-scale science operation in 2012 (Figure 3.60). It is an
imaging and spectroscopic instrument that will contain 50 to 64 antennas of 12 m diameter
(the number of antennas is a financial question currently under debate), plus an additional
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Molecular Universe, HS 2009, D. Fluri, ETH Zurich
Figure 3.59: The ten instruments belonging to the VLBA.
compact array of four 12 m antennas and twelve 7 m antennas. The array is located Llano de
Chajnantor Observatory in the Atacama Desert in northern Chile at an altitude of 5000 m. The
12 m antennas will have reconfigurable baselines ranging from 150 m to 16 km. The
instruments can observe wavelengths in the sub-millimeter to millimeter range and operate in
all atmospheric windows between 350 µm and 10 mm, and thus cover the short wavelength
end of radio waves. ALMA can reach a spatial resolution of 10 milli-arcsec and will become
the largest and most sensitive instrument in the sub-millimeter and millimeter range. The
science goals of ALMA include studying the first stars and galaxies in the universe by
imaging the redshifted dust continuum emission at epochs as early as z = 10, studying the
physical and chemical conditions of protoplanetary disks, young stars, and circumstellar
shells and envelopes around evolved stars, novae and supernovae, and probing the interstellar
medium in different galactic environments.
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