Radio_Astronomy.ppt

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ASTR 3520ASTR 3520Observations & Instrumentation II:Observations & Instrumentation II: SpectroscopySpectroscopy

Intro to Radio Astronomy

• Concepts - Amplifiers - Mixers (down-conversion) - Principles of Radar - Radio Astronomy basics: System temperature, Receiver temperature Brightness temperature, The beam ( = / D) [ its usually BIG] Interferometry (c.f. the Very Large Array – VLA) Aperture synthesis

Detection of EM Radiation Across the Spectrum

-ray X-ray UV Visible NIR TIR sub-mm mm cm radio

>MeV 0.2 – 100 keV 4000 – 10,000A 3 – 200 m 1 – 10 mm >1 cm 100 – 4000 A 1 – 3 m 200 m – 1000m

Incoherent detection Coherent detection - Particle properties - Wave properties (photons) - Energy, arrival time - Amplitude, phase - quantum noise: Trec > h / 2

Transmission

History of Radio Astronomy (the second window on the Universe)

• 1929 - Karl Jansky (Bell Telephone Labs)• 1030s - Grote Reber• 1940s - WWII, radar - 21 cm (Jan Oort etc.)• 1950s - Early single dish & interferometry - `radio stars’, first map of Milky Way - Cambridge surveys (3C etc)• 1960s - quasars, pulsars, CMB, radar, VLBI aperture synthesis, molecules, masers (cm)• 1970s - CO, molecular clouds, astro-chemistry (mm)• 1980s /90s – CMB anisotropy, (sub-mm)

Karl Jansky, Holmdel NJ, 1929

300 m dish at Arecibo:300 m dish at Arecibo: / D ~ 48” at = 6 cm (5 GHz)

408 MHz radio sky408 MHz radio sky

Outline• A Simple Heterodyne Receiver System

– mixers and amplification• Observing in the Radio

– resolution– brightness temperature

• Radio Interferometry• Aperture synthesis

Mixers

signal inLO

local oscillator

signal out and

A mixer takes two inputs: the signal and a local oscillator (LO).

The mixer outputs the sum and difference frequencies.

In radio astronomy, we usually filter out the high frequency (sum) component.

Mixers

frequency

originalsignal

mixedsignal

Mixers

frequency

originalsignal

mixedsignal

The negative frequencies in the difference appear the same as a positive frequency.

To avoid this, we can use “Single Sideband Mixers” (SSBs) which eliminate the negative frequency components.

W-band (94 GHz,4 mm) amplifier

Local oscillatorDownconverted signal

Frequency

Single sideband mixer:

Band-pass of amplifier: Intermediate frequency = IF

Local oscillator

Amplifier passband

Frequency

Lower sideband Upper sideband

FLOFIF

Double sideband mixer:

Local oscillator

Amplifier passband

Frequency

Lower sideband Upper sideband

FLOFIF

ASTR 3520ASTR 3520Intro to Radio Aastronomy:Intro to Radio Aastronomy:

• Review heterodyne & mixing Radar examples: cars, Venus (review)• Single dish spectroscopy Orion nebula CO example• Radiometer equation (noise vs, exposure time)• Interferometer basics The U, V plane => Dirty beams, Fourier Inversion => Dirty maps De-convolve Dirty Beam => CLEAN maps • The VLA, VLBA, and ALMA

Freflect = f trans + / - f f = 2 f (V/c) f = 10 GHz V = 100 km/h = 27.8 m /s c = 3 x 1010 cm/s = 3 x 105 km/s f = 1850 Hz

Radar:

f = 1850 Hz

f trans

Freflect = f trans + / - f

Mixing: Adding waves together

Local oscillatorDownconverted signal

Frequency

Single sideband mixermixer:

f = 10 GHz

F + f = 10 GHz + 1850 Hz

1850 Hz f = fIF

Band-pass of amplifier: Intermediate frequency = IF

Planetary Radar imaging: Doppler shift + time delay = 2D map

Radar PulseEarly echo

Late echo

BlueshiftRedshift

VenusVenusUV

VenusVenusRadar

A Simple Heterodyne Receiver

low noiseamplifier

filter

receiver horn

tunablefilter

signal @ 1420 MHz

1570 MHz

1420 MHz

tunableLO

~150 MHz

Analog-to-DigitalConverter

Computer

outputs a power spectrum

150 MHz

Amplification• Why is having a low noise first amp so

important?– the noise in the first amp gets amplified by all

subsequent amps– you want to amplify the signal before subsequent

electronics add noise• Amplification is in units of deciBells (dB)

– logarithmic scale• 3 dB = x2• 5 dB = x3• 10 dB = x10• 20 dB = x100• 30 dB = x1000

)(log10dB2

110 VV

Observing in the Radio I

• We get frequency and phase information, but not position on the sky– 2D detector

• A CCD is also a 2D detector (we get x & y position)

Observing in the Radio II:Typical Beamsize (Resolution)

• i.e. The BURAO 21 cm horn (D ~ 1 m)

o1210021

Observing in the Radio II• i.e. The NRAO GBT (D ~ 100 m)

'2.710000

''2.6'10.010000

3.0≈

at 21cm = 1.420 GHz

at 0.3 cm = 100 GHz

Observing in the Radio II• i.e. The Arecibo Telescope (D ~ 300 m)

'4.230000

''1.230000

3.0≈

at 21cm = 1.420 GHz

at 0.3 cm = 100 GHz

300 m dish at Arecibo300 m dish at Arecibo

Transmission (and brightness of the Atmosphere) depends on H2O!

Observing in the Radio III:Brightness Temperature

Flux: erg s-1 sr-1 cm-2 Hz-1 (1023 Jy)B(T): erg s-1 sr-1 cm-2 Hz-1 (1023 Jy)

We can use temperature as a proxy for flux (Jy)

Conveniently, most radio signals have h/kT << 1, so we can use the Raleigh-Jeans approximation

B(T) = 2kT/2

Thus, flux is linear with temperature

Antenna TemperatureTB = F2/2k

• Brightness temperature (TB) gives the surface temperature of the source (if it’s a thermal spectrum)

• Antenna Temperature (TA):– if the antenna beam is larger than the source, it will

see the source and some sky background, in which case TA is less than TB

TA ~ TB s/b

System Temperature

• Noise in the system is characterized by the system temperature (Tsys)– i.e. you want your system temperature

(especially in the first amp) to be low

Radiometer Equation• Trms = Tsys / ( t) ½

– Trms = r.m.s. noise in observation– a ~ (2)1/2 since you have to switch off-source = position switch

off-frequency = frequency switch

– Tsys = System temperature– = bandwidth, frequency range observed– t = integration time (how long is the exposure?)

Spectral Resolution• The spectral resolution in a radio

telescope can be limited by several issues:– integration time (signal-to-noise)– filter bank resolution (if you’re using a filter

bank to generate a power spectrum in hardware)

•Single Dish line & continuum basics • = / D ~ arc minutes 10 m @ 1 cm => 250”

• Brightness temperature: B(T) => I T

• Radiometer equation: TRMS = Tsys / (

Ex: = 100 GHz, = 1 MHz (= 3 km/s), = 1 sec, Tsys = 1000 K TRMS = 1 (Kelvin)

Orion B

Orion A

Orion Nebula

Orion MolecularClouds

13CO 2.6 mm

13CO J=6-5 661 GHz

(Wilson et al. – in prep)

12CO J=9-8(discovery)

12CO J=9-8

OMC1OMC1-S(Kawamura et al. 2002)

THz spectra!

Milky Way all-sky: Visual wavelengths

Milky Way all-sky: Infrared wavelengths (COBE)

408 MHz radio sky408 MHz radio sky

21 cm HI all-sky map21 cm HI all-sky map

Dame et al. CO map of Milky WayDame et al. CO map of Milky Way

R Cor Aust.

-Ophiucus

Galactic Center

Multi-transitionAnalysis ofOutflow massspectra

Multi-transition analysis of outflow mass spectra

L1551 J=2-1 CO

[SII] with COL1551

Radio Interferometry

e dela

Two Dish Interferometry• The fringe pattern as a function of time

gives the East-West (RA) position of the object

• Also think of the interferometer as painting a fringe pattern on the sky– the source moves through this pattern,

changing the amplitude as it goes

Extended Sources

+ -+ - -+ + • Spacing of the fringes is a function of the projected baseline

• The area under the fringes determines the amplitude of the signal (positive fringes add, negative fringes subtract)

• The projected baseline changes as the source rises or sets

Extended Sources

The narrow fringes (not visible here), represent the positional delay.

The broad envelope is the self-interference of the extended source.

A four hour observation of the Sun at 12 GHz using a two dish E-W interferometer.

Extended Sources• In order to determine the extended object’s

shape, we must disentangle the fringes due to the projected baseline (which we’d get from a point source) from the interference of the different parts of the source– to do this, we use delay lines– we introduce a delay between the two antennas

to compensate for the the positional delay– this leaves only the fringes from the structure of

the source

Aperture Synthesis• A two dish interferometer only gives

information on the E-W (RA) structure of a source

• To get 2D information, we want to use several dishes spread out over two dimensions on the ground

Radio Telescope ArraysThe VLA:An array of 27 antennas with 25 meter apertures

maximum baseline: 36 km

75 Mhz to 43 GHz

Very Large Array radio telescope (near Socorro NM)Very Large Array radio telescope (near Socorro NM)

The VLAThe VLA

An amplifier

The U-V Plane• Think of an array as a partially filled aperture

– the point source function (PSF) will be have complicated structure (not an airy disk)

– the U-V plane shows what part of the aperture is filled by a telescope

– this changes with time as the object rises and sets

– a long exposure will have a better PSF because there is better U-V plane coverage (closer to a filled aperture)

The U-V plane

a snapshot of the U-V plane(VLBA)

U-V coverage in a horizon to horizon exposure

Point Spread FunctionThe dirty beam : the diffraction pattern of the array

Fourier plane samplingUV plane covered over 6 hrs Amplitude of fringes on a source

Examples of weightingDirty Beams:A snapshot (few min) Full 10 hrs VLA+VLBA+GBT

Image Deconvolution• Interferometers have nasty PSFs• To get a good image we “deconvolve” the

image with the PSF– we know the PSF from the UV plane coverage– computer programs take a PSF pattern in the

image and replace it with a point– the image becomes a collection of point

sources

UV Plane Coverage and PSF

images from a presentation by Tim Cornwell (given at NRAO SISS 2002)

UV Plane Coverage and PSF

Image Deconvolution

The Orion Nebula (at 1.4 GHz)

SNR Cassiopeia AContinuum in 3 colors:1.4 GHz (L band)5.0 GHz (C band)8.4 GHz (X band)

Jupiter andits magnetosphere

W50 SNR: SS 433 (accreting neutron star + 0.27c jets)

21 cm HI map21 cm HI mapof M33of M33

Centaurus A (the nearest radio galaxy)

PH 227PH 227

Radio galaxyand jets

Radio galaxyRadio galaxy3C2863C286

VLBAVLBA

Radio Telescope ArraysALMA:An array of 64 antennas with 12 meter apertures

maximum baseline: 10 km

35 GHz to 850 GHz

Radio_Astronomy.ppt

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