CHAPTER 3 Telescopes of Other Wavelengths 0. Nonoptical Telescopes

1
CHAPTER 3 Telescopes of Other Wavelengths 0.
Non-optical Telescopes

• The Earths atmosphere limits detection of light
  from the universe. The only two clean observing
  windows are in optical and radio wavelengths.
• Partially open windows include infrared and
  microwave.
• Totally blocked (opaque) windows are the high
  energy radiation (g-ray, X-ray, and UV), and long
  wavelength radio waves.
• We can access the partially blocked and opaque
  windows by going above the Earths atmosphere
  (airplane 15 km, balloon 30 km, orbiting
  altitude)

2
I. Infrared Observations

• Convention 1-5 mm (near IR), 5-200 mm (far-IR)
  Infrared radiation from space absorbed by gas
  molecules (H2O, O2, CO2) in the atmosphere
  (absorption bands)
• Ground-based infrared observations are done at
  optical telescopes, with detectors sensitive to
  IR wavelengths. The state of the art detector is
  the Mercury-Cadmium-Tellurium (HgCdTe)
  two-dimensional arrays (20482, compared with 2562
  being used in space right now!).
• Major source of background for IR observations is
  thermal radiation from the environment, i.e.
  blackbody radiation from the telescopes and
  detectors
• Full IR wavelength range observations can be
  achieved in space. Background can be reduced by
  active cooling of the detectors by dewers.

3
Blackbody Radiation

• All objects (including astronomical objects) emit
  radiation all the time which is independent of
  size, shape, and composition, and is dependent
  only of temperature, called the blackbody
  radiation
• The spectrum (intensity versus frequency
  distribution) is a function
• where k Boltzmanns constant
• 1.38 x 10-23 JK-1,
• and T temperature (unit Kelvin)
• The unit of intensity is in J s-1 m-2 Hz-1
  ster-1
• (ster steradian unit of solid angle)

Adopted from Astronomy Today by Chaisson
McMillan
4
Blackbody Radiation

• The temperature of an object is related to the
  wavelength lmax at which it emits the most
  radiation by the Wiens law
• e.g. (i) Surface temperature of the sun 6000 K,
• lmax at 0.0029/6000 480 nm (yellow-green)
• (ii) Room temperature 20oC 293 K,
• lmax at 0.0029/293 10 mm (IR!)
• (To be used later) Total energy emitted by a
  blackbody per unit area per second, F, (energy
  flux)
• F s T4, where s Stefan-Boltzmann constant
• 5.67x10-8 J s-1 m-2 K-4
• Stefans Law

Adopted from Astronomy Today by Chaisson
McMillan
5
Infrared Telesopes

• To improve IR observing, we need to observe from
  above the ground to be above as much of the
  atmospheric gas as possible (especially water).
  The observing window opens up as we observe at
  higher altitude.
• The Kuiper Airoborne Observatory is a 0.9m
  telescope flied inside a C-141 transport plane.

6
Kuiper Airborne Observatory

• Cruising altitude 9-13 km
• Flight duration 7-8 hours (11 hours max)
• Main Telescope 0.95m Cassegrain configuration
  reflector
• Detection Range 1 500 mm (very large
  background thermal radiation in the long
  wavelength limit though)
• Operating temperature -50 -70 oC (no active
  cooling)

7
SOFIA

• The Stratospheric Observatory for Infrared
  Astronomy (SOFIA), to be launched in October 2004
  2009(?)
• Successor to KAO
• A 2.7m f/19.6 Cassegrain (with Nasmyth focus)
  mounted on a Boeing 747 SP jet

8
SOFIA

• Wavelength range 0.3 1,600 mm
• Temperature of telescope 240 K
• Operating altitude 12 14 km
• Image stability 0.2? (rms)
• Diffraction-limited wavelengths gt 15mm

9
Spitzer Space Telescope (SST)

• Formerly known as Space Infrared Telescope
  Facility (SIRTF)
• Launched in August 2003
• A 0.85m f/12 Ritchey-Chrétien telescope made with
  Beryllium
• Operating wavelength 1-300 200 mm
• Lifetime (limited by cryostat) 2.5 5 years

10
SST

• Earth trailing heliocentric orbit (enables lower
  temperature observations than low-earth orbits)
• Scientific instruments are passively cooled with
  liquid helium to 5.5K to minimize the IR thermal
  radiation.
• Other advantages of orbit no need to avoid earth
  and moon during observations, no need to avoid
  Earths radiation belt (van Allen belt)

11
SST Orbit in Space
12
Observing with SST
13
II. Ultraviolet Observations

• Cutoffs of UV radiation 3300 Å (by O3
  absorption), 912 Å (Hydrogen Lyman continuum)
• Convention 2000 3300 Å (near UV), 912 2000 Å
  (far UV), 100 912 Å (extreme UV)
• UV can only be observed above the Earths
  atmosphere (need astronomical satellites)
• Detectors SiC and LiF-coated arrays
• Optical components of ultraviolet telescopes are
  similar to those used for the optical telescopes
• Hubble Space Telescope (mainly an optical
  telescope) is equipped with UV detectors.

14
International Ultraviolet Explorer (IUE)

• Launched on Jan 1978
• Geosynchronous orbit (revolve around the earth in
  24 hour orbit), allowing observers to make real
  time observations)
• 0.45m f/15 Ritchey-Cretien Cassegrain
  configuration
• Wavelength 1150 3300 Å, two spectroscopic
  channels, resolution (0.1 Å and 6 Å)
• Image quality 2?
• Detector Reticon tube (as used in a television
  set!)
• Satellite expected lifetime 3-5 yrs
• Actual lifetime 18.7 yrs

15
Far Ultraviolet Spectroscopic Explorer (FUSE)

• Launched June 1999
• FUSE operates in the far UV range 912 1150 Å
  not covered by IUE
• Data collected from four mirrors, each 35 cm2
• Specializes in high spectral resolution (echelle)
  observations
• Can now operate in zero-gyro mode

16
III. X-Ray and g-Ray Telescopes

• The combination of X-ray and g-Ray astronomy is
  also called High Energy Astronomy
• Convention Separation of X-ray and g-ray at 0.5
  MeV. Therefore, annihilation of a positron and an
  electron (Rest Mass 0.511 MeV/c2) will yield
  two g-ray photons.
• Convention Soft X-ray (0.1 1 keV), Hard X-ray
  (1 keV 0.5 MeV), g-ray (gt 0.5 MeV)
• All X-ray and g-ray are absorbed by the Earths
  atmosphere (so that we can live safely!) Space
  observatories are needed.

17
Chandra X-Ray Observatories
Jul 1999
XMM- Newton
Suzaku XIS 0.4 10 keV CCD M
Japan 2005 HXD 10- 700 keV
Scintillator
Adopted from Observational Astrophysics, by P.
Lena
18
Compton Gamma-Ray Observatory (CGRO)
Oct 2002

• Adopted from Observational Astrophysics, by P.
  Lena

19
III. X-Ray and g-Ray Telescopes

• Theoretical diffraction limit of a X-ray
  telescope
• e.g. For a 1 keV photon, l 12 Å,
• ? q 1.22 l/D 0.003?/D (very small!!)
• However, two problems
• First, the telescope surface needs to be
  manufactured to an accuracy at least 10 of the
  wavelength of the light.
• ?Accuracy for an X-ray mirror needed 1 Å,
  which is the size of an atom!
• Second, most X-ray photons simply penetrate
  through the surface of a material and are too
  energetic to be reflected at normal incidence.
• Therefore, while IR and UV telescopes have
  similar optics setup as the visible telescopes,
  high energy telescopes have very different
  configurations because X-ray and g-ray cannot be
  focused by a spherical mirror.

20
High Resolution Imaging in X-Ray

• At very high incidence angle (gt 88o, grazing
  incidence), X-ray light reflect off metallic
  surfaces.
• X-ray can be focused by reflecting off a
  paraboloid and a hyperboloid (grazing-incidence
  telescope), with no spherical aberration!
• Several layers of mirrors are nested to increase
  the photon-collecting area.
• Collecting area decrease with increasing X-ray
  energy (hard X-ray need incidence angle closer to
  90o)

21
Focusing by grazing in X-ray telescopes
Chandra X-Ray Observatory as example
Chandra Science Center
22
X-ray Telescopes

• This technique can be applied to X-ray of energy
  up to 100 keV
• X-ray mirror now only need to be manufactured to
  the accuracy related to sinq, where q is the
  incidence angle (accuracy required several x 10
  Å)
• Also, because of the large incident angle,
  Grazing Incidence Telescopes are long compared to
  the entrance aperture (Question Is f/ number
  large or small?)

XMM- Newton
Adopted from Observational Astrophysics, by P.
Lena
Chandra X-Ray Observatories
23
X-ray Detectors

• Historically, the main type of detectors are the
  proportional counters for soft X-ray (lt20 keV)
  and scintillation detectors for hard X-ray
• A scintillation detector works as an incident
  X-ray photon ionizes an atom in the crystal
  lattice, producing a high-energy free electron.
  Some of these electrons are later recaptured by
  impurity atoms in the lattice and produces
  flashes of visible light. Then these photons are
  detected by a photomultiplier tube attached to
  crystal.
• These days, specially coated CCD (Charge- Coupled
  Device, to be discussed in next Chapter) are used
  for imaging.
• Reflection-grating arrays are used for
  spectroscopy

Principle of a scintillation detector
Adopted from Observational Astrophysics by Robert
C. Smith
24
Chandra X-Ray Observatories

• Launched on July 1999, with a expected lifetime
  of 5 years
• Diameter (of the largest pair of grazing mirror)
  1.2m, focal length 10m
• Four scientific instruments providing imaging and
  spectroscopic coverage for 0.2 10 keV X-ray
• Angular resolution 0.5? (best for X-ray)
• Highly elliptical and time-varying orbit perigee
  16,800 km, apogee 132,000 km, with an
  orbiting period of 65.5 hours

25
Chandra X-Ray Observatories
26
XMM-Newton

• Launched in Dec 1999
• Detecting 0.1 12 keV
• Aperture 0.7m, focal length 7.5m, Angular
  resolution 5?
• Three separate mirror modules providing the
  largest effective area, good for spectroscopy

27
Gamma-ray Telescopes

• The g-rays are too energetic for even grazing
  telescopes to be effective, leading to poor
  angular resolution.
• Imaging have historically been done with
  secondary particle track analysis
• Recent advances include developments of
  semiconductor detector arrays such as CdTe,
  (1282) and CsI (642 scintillation crystal).
• Three methods, sometimes used in combination, for
  focusing g-ray (1) partial or total absorption
  of the g-ray's energy within a high-density
  medium, such as a large crystal of sodium iodide,
  (2) collimation using heavy absorbing material,
  to block out most of the sky and realize a small
  field of view, and (3) at sufficiently high
  energies, utilization of the conversion process
  from g-rays to electron-positron pairs in a spark
  chamber, which leaves a directional signature of
  the incoming photon.

28
Compton Gamma-ray Observatory (CGRO)

• Launched by space shuttle in 1991, burnt during
  re-entry to Earth in 2000
• Four instruments to detect hard X-ray and g-ray
  photons from 20 keV to 10 GeV

29
CGRO Instrument Capabilities
Burst and Transient Source Experiment
Oriented Scintillation Spectrometer Experiment
Imaging Compton Telescope
Realized upper limit 10 GeV
Energetic Gamma Ray Experiment Telescope
30
Fermi Large Area Space Telescope

• Launched in 11 June 2008, joint US-Europe-Japan
  mission
• Successor to CGRO, will be first telescope to
  survey E gt 10 GeV
• Two instruments in one telescope
• LAT Large Area Telescope (g-ray 20 MeV 300
  GeV)
• GBM GLAST Burst Moniter (X-ray to g-ray 10 keV -
  25 MeV)
• Angular resolution for imaging (LAT) 1?

Notice Different shape of a g-ray telescope (No
mirror or tubes)
31
Adopted from Observational Astrophysics, by P.
Lena
32
IV. Radio Telescopes

• While radio and optical radiation are the only
  two clear windows for observations through the
  atmosphere, there are significant differences
  between optical and radio telescopes.
• The major reason for difference lies in the big
  difference in wavelength lradio/loptical 105
  106, which also means optical photons have much
  higher energies than radio photons.
• The large wavelength ratios has two major
  consequences the design of the instruments, and
  the format of the resulting data obtained.
• The radio window, spanning from mm to tens of
  metres, is much wider than the optical window
  (300 1000 nm).

33
Language of Radio Astronomy

• Unit of flux density (Power per unit area per
  unit frequency) 1 Jansky (Jy) 10-26 Wm-2Hz-1
• Strongest sources 104 Jy detection limit mJy
• At radio wavelengths (hn ltlt kT), the blackbody
  spectrum can be expressed in terms of the
  Rayleigh-Jeans approximation Bn(bb) 2kT/l2
  (Unit Wm-2Hz-1ster-1)
• In radio astronomy, we frequently look at the
  brightness temperature TB of source defined as
  fn(source) 2kTB/l2, where fn is the specific
  flux density of the source measured.
• Each component of the radio telescopes can
  therefore be expressed in terms of temperature
  (e.g. background radio signal detected at
  telescopes system temperature Tsys)

34
Radio Background

• While radio signals are not absorbed by the
  atmosphere, observations in radio are affected by
  other background sources, both natural and
  manmade.
• The extreme weakness of power from astronomical
  source means all the transmission and
  amplification of signals must be kept with low
  power.

Courtesy Undergraduate Research Educational
Initiative
35
Design of Radio Telescopes

• The angular resolution of a telescope of aperture
  diameter D is q 1.22 l/D.
• To obtain the same resolution for optical
  telescope (500 nm) as radio telescope (1 cm), we
  need a diameter that is 1 cm/500 nm 20000 times
  bigger. (Exercise calculate the equivalent size
  for a 10m wave radio telescope to have the same
  diffraction-limited resolution as a 1m optical
  one)
• Therefore, usually, radio telescopes are
  diffraction-limited, i.e. the resolution q is
  given by 1.22 l/D.
• Radio astronomers use the technique of
  interferometry to achieve better resolution than
  optical telescopes.

36
Detection of Radio signals

• Optical detectors such as CCDs detect the
  individual photons that strike the surface,
  creating a response that is proportional to the
  number of photons striking it. Counting these
  photons will result in a measure of the intensity
  of the source (coherent detection).
• Radio photons carry energy that is too low to
  cause a reaction in such detectors.
• Radio receivers will instead detect the wave
  nature of the radio waves rather than the photon
  nature (incoherent detection).
• Radio waves are detected by the alternating
  electro-magnetic field generated in the detector,
  which is then detected as an AC voltage.
• Therefore, while we cannot count radio photons,
  we can measure amplitude, phase, and polarization
  of the wave.

37
Structure of Radio Telescopes

• Radio telescopes are mainly reflecting telescopes
  of either prime focus or Cassegrain types with a
  parabolic reflector (antenna)
• Radio telescopes are usually much bigger than
  optical telescopes to achieve better resolution.
• The biggest single dish radio telescope
  (filled-aperture telescope) is 305m diameter
  spherical dish at Arecibo (fixed) and 100m
  diameter parabolic dishes at Effelsberg, Germany
  and Greenbank, USA (fully steerable).
• The focal ratio f/D for radio telescope is
  usually smaller than 1 ( 0.3 - 0.4) why?
• The large research steerable radio telescopes
  usually have alt-az mounting, similar to that of
  optical telescopes.

38
Arecibo, Puerto Rico

• 305m spherical dish
• Prime focus feed

39
Greenbank, West Virginia, USA Prime focus
Effelsberg, Germany Cassegrain
40
Structure of Radio Telescopes

• Radio telescopes are made up of two main parts
  antenna and receiver
• The antenna collects and focus the plane waves
  from the distant astronomical source and into
  converging spherical waves and converts these
  spherical waves into an electrical AC signal.
• The receiver amplifies the input signal (with a
  low-noise amplifier LNA), chooses the signal
  frequency and bandwidth, processes the signal and
  records it. (This is very similar to the system
  found in conventional radio and television)

41
Antenna

• Most research radio telescopes for millimetre and
  centimetre waves have parabolic or spherical
  dishes
• Antennas, like primary mirrors of optical
  telescopes, need to be built accurately within at
  least 10 of the wavelength.
• Therefore, surfaces of radio antennas are simply
  metal sheets, and metal wire meshes (for longer
  radio waves)
• For radio waves of length gt 1m, Yagi-Uda type
  antenna (similar to the TV antennas seen on
  rooftop) can be used.

76 m Lovell Telescope, Jodrell Bank
42
Feed/Feed Horn
A dipole placed here

• Placed at the prime focus
• Convert the radio waves into an electric AC
  signal to a transmission line
• Example a half-wave dipole consists of a
  straight metal rod of length l/2, with a cable to
  the receiver attached at the centre.

Circular Waveguide
43
Beamwidth
Polar power pattern of normal illuminated source

• Power of radio waves received by a radio dish at
  various angular directions can be expressed in a
  power pattern diagram.
• The angular resolution of a radio antenna, also
  known as the half-power beamwidth (HPBW), or
  simply beamwidth, is determined by the width of
  the main peak
• Circular aperture q1.22 l/D
• Main beam can be thought of as the central Airy
  disk of the diffraction pattern of a point source
  for optical telescopes.

main peak
side lobes
Adopted from Robert Smith Observational
Astrophysics
44
Filled Aperture Telescope Array

• To improve resolution for radio data, we can
  construct an array of telescopes all feeding
  their signal to a single receiver.
• This is similar to the multi-mirror approach for
  optical telescopes.
• The signals from different antennas are at
  different distance from the receiver. We can
  introduce electrical phase delay to the system to
  ensure that all signals arrive at the receiver
  simultaneously.

Australia Telescope Compact Array (ATCA), 6x22m
45

• The combination of the electrical signals can be
  accomplished by electrical steering.

• Consider the three-element array above, the
  signal from zenith will arrive in phase (assuming
  the same cable length)
• For sources at angle a, we need to introduce
  phase shifts to the signals so that they will
  arrive at receiver in phase.
• For example, need to introduce phase shift of f,
  and 2f to signals at element 2 and 1 respectively
  to match with that at element 3, where
• Note usually phase change due to separation of
  elements will be many cycles, i.e. f 2np f0.,
  where ngtgt1 and 0lt f0 lt2p.

46
Radio Interferometry

• Resolution of an aperture-filled telescope
  actually depends only on the largest dimension of
  the antenna.
• Therefore, we can keep the same resolution even
  when large parts of the apertures are removed!
  (unfilled-aperture telescopes)
• We can have a series of antennas (size d)
  separated at distance D, then angular resolution
  achieved will be l/D, not l/d
• The combination of signals from these antennas is
  called interferometry. We need precise knowledge
  of the phase of the radio waves received.

47
Principle of Interferometry
48
Aperture Synthesis

• Signals from separate antennas combined. Phase
  shifts are introduced to signals for in-phase
  combination.
• Phase shift applied depends on (a) orientation
  and length of baseline (b) diurnal motion of
  Earth
• Correlator used to maximize combined signal from
  different antennas
• Detailed maps resulted from Fourier Transform of
  the data from differenet baselines (helped by
  diurnal motion!)
• To obtain wider coverage of the sky, we need to
  maximize the number of distinct baselines of
  antennas.

Adopted from Observational Astrophysics by P. Lena
49
Very Large Array

• Socorro, New Mexico, USA
• 27 antenna, 25m each
• Maximum separation 36 km
• Detection range 0.1 50 GHz
• Antennas in Y configuration to maximize the
  number of distinct baseline orientations and
  separations.

50
Very Long Baseline Interferometry (VLBI)

• To further improve angular resolution of the
  resulting radio maps, we need to increase
  separation between antennas.
• However, there is a practical limit on the
  separation of antennas linked by electrical
  cables to the receivers.
• One solution Antennas around the world observe
  an object simultaneously, record the data on
  storage device (magnetic tape), and send the data
  to a central location for correlation and
  combination.
• Biggest challenge To calculate the correct phase
  shift applied to all data, all observing stations
  need synchronous clocks (hydrogen maser
  oscillator locked to Global Positioning System
  time) for accurate timing (phase accuracy 10-14
  over a few hours.
• Resolution 10-3 10-4 arcsec in millimetre
  region

51
USA VLBI Network
52
European VLBI Network
Sheshan, Shanghai 25m
Urumqi, Xinjiang 25m
53
Small Radio Telescope (SRT)

• A small radio telescope developed by MIT/Haystack
  observatory for teaching purposes
• Dish 2.3m diameter
• Receiver 1370-1800 MHz
• Can detect emission from hydrogen atoms (through
  spin-flip emission line at 1420 MHz)
• Construction completed in Dec 2006, undergoing
  testing now
• Possible projects radio emission from the sun,
  our Galaxy, pulsars,
• Possible dish in the future for interferometry

54
Optical Interferometry

• Interferometry technique can also be applied in
  optical wavelength
• E.g. CHARA Six 1-m diameter telescopes,
  connected by vacuum pipes for light to be
  combined
• Maximum baseline 330 m
• Best resolution 200 mas
• Need very accurate correlator to combine lights
  from different telescopes

55
V. Neutrino Telescope

• Other than photons, we have only directly
  detected neutrinos from astronomical sources.
• Neutrinos are low mass, spinless, chargeless
  particles that hardly interact with other matter.
  The mean free path of neutrinos through water is
  0.1 light year.
• First neutrino telescope in 1964 is a tank of
  450m3 (100,000 gallons) of bleaching fluid
  (C2Cl4) 1.5km underground in a landmine.
  Chlorine37 capture neutrinos by reaction Cl37n
  ? Ar37 e-. Count the number of radioactive Ar37
  counted every few months when the tank is flushed
  out.
• Result dozens of n (from Sun) detected every few
  months!
• Neutrinos from Supernova 1987A (14 within two
  seconds of explosion in two observatories
  Kamiokande II and IMB) detected.

56
Sudbury Neutrino Observatory (SNO)
2.1 km (6800 feet) underground 1000 tonnes of
heavy water (D2O) Most expensive ground telescope
57
VI. Gravitational Wave Telescope

• According to General Relativity, rapid changing
  of gravitational field should emit gravitational
  waves, which are basically oscillations in
  space-time.
• Indirect evidence of existence by the spiralling
  down of binary pulsars
• We expect violent events such as the formation of
  black holes, and neutron star-black hole merger
  to be the source of gravitational waves.
• Detect by gravitational effects of these waves,
  which is extremely small (length variation of
  order 10-22) with laser interferometer.

58
Laser Interferometer Gravitational Wave
Observatory (LIGO)
Started operation in 11/2005
LIGO Hanford Observatory Redmond, WA, USA
LIGO Livingston Observatory Livingston, LA, USA
L shape design, 4km on a side