Radio Astronomy: 101

Danae PolychroniUniversity of Athens – Department of Physics- 12/10/2015

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro

Radio Astronomy: intro The oldest of the 'new' astronomies since it is now 83 years since Karl Jansky first observed radio waves from a galaxy.

It has passed beyond the development stage (unlike IR, UV X-rays etc)

It was developed so early mainly due to the radio radiation being able to penetrate to ground level. From 10mm to 10m the atmosphere is practically transparent & between 0.5mm and 100mm there are some transparent windows.

Basic Definitions: Intensity Iν

dW=Iνcos d d d θ Ω σ ν W: power; d : the source sold angle; d : detector areaΩ σ

d : frequency bandiwdth; : the angle between the normals ν θTo source and detector

Flux Density Fν Fν=∫Ω Iν cos dθ Ω Normal units are Jansky: 1Jy = 10-26 Wm-2s

Detectable sources can be anywhere between 0.001 to 106 Jy

Radio Astronomy: Telescopes Usual radio telescopes are simple reflecting telescopes made of wire mesh or metal plates.

A paraboloidal reflector brings all rays to focus at the same point on the optical axis and eliminates spherical aberrations.

Advantages: High gain: The larger the 'dish' in termsof wavelengths, the higher the gain. High directivity: The higher the gain, the narrower the beamwidth. This can be a significant advantage in applications where the power is only required to be directed over a small area.

Disadvantages:

Requires reflector and drive element: the parabolic reflector itself is only part of the antenna. It requires a feed system to be placed at the focus of the parabolic reflector. Cost : The antenna needs to be manufactured with care. Size: The antenna is not as small as some types of antenna, although many used for satellite television reception are quite compact.

Radio Astronomy: Telescopes All large disk radio telescopes are mounted in an altitude-azumuth configuration (perpendicular and parallel to the horizon), where the axes are vertical and horizontal (cheaper and simpler to construct)

To track a sources across the sky both axes need to bedriven at variable angular velocity as alt-az does not followthe natural motion of the sky.

Optical Configurations:

>Prime Focus: detector or receiver at the focus of the paraboloidal reflector

>Cassegrain: a convex hyperboloidal secondary mirror or reflector is used to bring the rays to a focus below a hole in the primary

>Gregorian: the secondary is now a concave elliptical mirror or reflector above the primary focus

>Nasmyth: like Cassegrain, except a folding flat that is added above or below the primary

Radio Astronomy: Telescopes Homologous Deformation

Any telescope designed at should have surface tolerance of /20 to λ λmaintain efficiency. This accuracy must remain as it orientation changes i.e. as it tracks a source across the sky.

This is pretty difficult for dishes bigger than 10m in diameter.

What is really required is not that is figure does not change with orientation but that its figure must remain a paraboloid of rotation.

Large radio telescopes are designed to deform but remain paraboloids.

Radio telescope performance in terms ofmaintaining the figure is usually limited bythermal expansion efects and wind many→ modern high-frequency radio telescopes are made from carbon fibre or other com-posite material instead of steel or aluminum.

Radio Astronomy: Antennae & Receivers Basic Definitions:

Brightness Temperature:Tb = Iνc2/(2ν2k)

Radiative transfer equation:Tb=Tb(0)e- ( )τ ν + T(1-e- ( )τ ν )

Noise Temperature: is the thermodynamic temperature of a resistor which would have the same detected power as that in the system noise.

TN = W/k W is noise power per unit frequency K is the Boltzmann's constant

Where W is the power in a frequency band d :ν

W = Ed = h d /(exp(h /kT) -1)ν ν ν ν

In the Rayeigh-Jeans limit:W=kTdν

Radio Astronomy: Antennae & Receivers Antenna Temperature

Power observed from a source is given by:

W = ½ A∫ Bν

Where is the efficiency, A is the collecting area, and the factor is because only one polarisation is detected

Antenna temperature is then defines as:

TA = W/k = ½ (A/k)∫Bν

In the R-J limit:

TA = A∫(Tb/λ2)

If the source is of uniform brightness:

TA = ATb source/λ2

Beam solid angle: beam= πθ2/4 = πλ2(1.22)2/4d2 where d is the telescope diametre

Telescope area: A=πd2/4 => beam=(1.22π/4)2λ2/A

=> TA=Tbsource/beam

Radio Astronomy: Detectors Radio detection Is a two stage process, involving an antenna that produces an electrical signal from the electric field of the radio frequency radiation, and a receiver which processes it.

The Antenna may involve a reflecting dish or other structure, or may be a basic aerial

The receiver will involve amplification, probably frequency change, integration and converting the signal to digital form.

Most modern antennae incorporate a corrugated horn waveguide feed. This is a horn designed to concentrate the radiation on a smaller detector, without loosing energy, and also designed to reduce sidebands.

Radio Astronomy: Amplifiers Signal from the detector is weak and must be amplified before it can be used

Amplifiers are designed to add as little noise to the system as possible.

Total gain required is in the range of 109 to 1014. This cannot be achieved by on amplification stage so cascades of amplifiers are used.

If there are n stages of amplification the total gain is:

G=Σi=1 Gi

If the antenna temperature TA and stage I contribute a system temperatre Tsi then we can work out the total system noise Ts

Input power of stage 1 is P0=kTA Stage i amplifies the signal by a gain Gi and introduces noise Tsi:

Pi=(Pi-1 + kTsi)Gi & Pn=(TA+Ts)Σi=1Gi

=> Ts=Ts1+(Ts2/G1)+(Ts3/G1G2) +. .+ (Tsn/G1G2.Gn-1)

n

n

Radio Astronomy: Amplifiers

Note that the stage gains are all at least of order of 10.

Transitor amplifiers, used in radio recievers, have been replaced by more exotic devices which operate at cooler temperatures of 5K or less, so that the front end (i.e. first stage amplifier) contribute really low noise to the system.

Such devices are: Cooled Field Efect Transistors (FETs) Maser Amplifiers High Electron Mobility Transistor (HEMT)

Radio Astronomy: Heterodyne Receivers In a heterodyne receiver the signal from the sky is mixed with a stable Local Oscillator signal to produce a lower beat frequency, the Intermediate Frequency (IF). The reasons for doing this are:

> at high radio frequencies (microwave) it is difficult to construct high frequency amplifiers. IT is much easier to amplify the IF

>cascade amplifier chains are unstable a small amount of power leaking from →the output port back to the input port causes instabilities. So the output frequency is made to difer from the input frequency.

Radio Astronomy: Heterodyne Receivers The mixer is the device which performs frequency conversion. Any device with nonlinear relation between input voltage and out current can be used as a mixer, but the most popular and easiest to analyse are square law devices:

I=aU2

Where U is the input voltage and I the output current

U is the sum of the signal:Esin(2πνst+δs)

And the local oscillator is: Vsin(2πνLt+δL)

So the output is given by:I=a[Esin(2πνst+δs)+(2πνLt+δL)]2

Mixers produce a range of frequencies but the most important is νIF=νS-νLWhich is the intermediate frequency (IF). The others are filtered out while this one is amplified and processed.

Mixers nowadays tend to be Superconductor-Insulator-Superconductor (SIS) devices which operate at low temperatures and deliver less noise.

Radio Astronomy: JCMT -HARP

Radio Astronomy: JCMT -HARP James Clerk Maxwell Telescope is a 15m diameter; the largest telescope in the world currently designed to observe in the submillimetre. It is situated close to the summit of Mauna Kea, Hawaii, at an altitude of 4092m.

Telescope

The 15-metre diameter primary reflector of the JCMT is made up of 276 individual lightweight panels.

Each panel consists of a thin aluminium skin bonded to an aluminium honeycomb and is attached at three points to the backing structure of the antenna.

The backing structure is designed to maintain a parabolic figure as gravity distorts the antenna as it tips to diferent elevations.

The surface accuracy is routinely measured and adjustments required to each panel are calculated by making observations of a coherent millimetre source located on top of the UKIRT building.

The sub-reflector or secondary mirror can be adjusted in three axes to compensate for changes in focus as well as changes in the figure of the primary. In addition, the secondary can be tilted or chopped in two axes in order to perform sky background cancellation.

Radio Astronomy: JCMT -HARP Carousel

The JCMT carousel co-rotates with the antenna and is designed to protect the telescope from the elements and to provide a safe and comfortable working environment for astronomers and engineers. An important feature of the carousel is the membrane which is deployed in front of the antenna at all times and is transparent at millimetre and submillimetre wavelengths. In addition to providing protection from the wind, the membrane performs the useful function of reflecting the visible and near-infrared radiation, providing protection from the solar heat which could damage the antenna, thereby allowing daytime astronomical observations including direct observations of the Sun itself.

Instruments

Receivers for the telescope can be located either in one of the bays of the Cassegrain cabin or on the two Nasmyth platforms located at the ends of the elevation bearing.

The heterodyne receivers, covering the atmospheric windows between 215 and 680 GHz, are mounted in the cabin while the continuum bolometer array instrument, SCUBA-2, and the heterodyne array receiver, HARP, occupy the Nasmyth platforms.

When observing together with the SMA and the CSO, the JCMT is part of the Extended Submillimeter Array (eSMA).

Radio Astronomy: JCMT -HARP Observing Techniques with Heterodyne Receivers:

Position-switching: the telescopes observes alternately the source and a nearby refe-rence position, from which pair of spectra the software obtains a calibrated result from which most instrumental and atmospheric efects have been eliminated. The integration time spent on a single signal or reference position is a compromise between observing efficiency and the rapidity with which conditions change, but is typically 30 seconds or less.

Beam-switching: In this technique, the secondary mirror is `chopped' (usually at 1 Hz, and in the azimuth direction, although chopping in any direction is possible) with the source in one beam. At a slower rate, typically every 10-30 seconds, the telescope is moved by an amount equal to the chop throw, to bring the source into the other beam. Thus a negative line signal is obtained when the source is in the reference beam, and a positive one when it is in the signal beam. Both signal and reference phases are neces- sary; if one attempts to omit the reference part of the cycle, standing waves are the likely result. Generally the technique is efective for sources of small angular size (i.e. smaller than the beam throw) , and should produce better results whenever sky noise is a problem (most of the time).

Frequency switching: In this case the frequency of the local oscillator is switched between two values of the LO frequency separated by, say, 16 MHz.

Point-by-point mapping: One can map an extended source by observing points defined via ofsets from a central position. The points may be observed using any of the standard single-point observing modes (see above).

Radio Astronomy: JCMT -HARP Observing Techniques with Heterodyne Receivers:

Raster mapping: For extended fields with relatively bright line emission, the raster mapping technique is a sensible choice. In this mode the telescope is scanned along a line of constant latitude while the ACSIS digital autocorrelation spectrometre continuously integrates the incoming signals, forming the average once every few seconds and saving the result. Raster mapping runs only in the position-switched mode, with an integration at a user-supplied reference position once at the end of every row.

Normally the sampling interval (`cell') should be less than half the HPBW to avoid smearing the map in the scan direction. One-third HPBW sampling is a good default choice.

The integration time per point ($ T$) should be 3-6 seconds per point; If you want to obtain more than 6 seconds/point, then repeat the map at the same position as often as necessary and average them together (this will give better baselines and calibration than a longer integration time per point).

Calibrations should be done at regular intervals, say once per 30, 60 or 90 minutes, depending on the sky conditions. Raster maps containing more than 1000 points are routine. Larger maps can be constructed by mosaicing sub-maps together in data reduction.

Radio Astronomy: JCMT -HARP Observing Techniques with Heterodyne Receivers:

HARP: the heterodyne receiver array program, it was commissioned and accepted by the JCMT in 2007. It has 16 receptors separated by 30 arcsec with a footprint of 2 arcminute. The beamsize of each receptor is approximately 14 arcseconds.

Receptor MapR13(H12) R14(H13) //////// R15(H14)////// R16(H15)R12(H11) R11(H10) R10(H09) R09(H08)R05(H04) R06(H05) R07(H06) R08(H07)R04(H03) R03(H02) R02(H01) R01(H00)

The tuning range is 325-375 GHz and the median receiver temperature is about 170K and system temperature is about 350K in band 2 weather.

Tsys359 353 N/A 397 375 383 395 336 292 277 375 235N/A 440 327 394

Radio Astronomy: JCMT -HARP Observing Techniques with Heterodyne Receivers:

Autocorrelation receiver is a radio frequency spectrometre which is essentially the analogue of the Michelson or Fourier Transform Spectrograph at optical and IR wavelengths.

The radio signal (which is frequency dependent) is multiplied by itself with a sequence of equally spaced spaced delays introduced (usually at a power of 2 (like 256 or 1024)

Integrating over time on all these channels produces a Digital Fourier Transform of the spectrum of the input radiation.

ACSIS: it is the digital autocorrelation spectrometer used as the backend for the receivers.

It has a total down converter modules (DCMs) and correlator cards. A DCM 'picks out' a 250 or 1000MHz wide subband from an IF input and feeds the signal to a correlator card.

With HARP having only 16 IF outputs while there are 32 DCMs and correlator cards these can be used:

Chaining 2 correlator cards together and doubling the number of channels in a subband.

Use two DCMs to extract 2 frequency subbands from each of the 16IF inputs and feed each subband to a correlator card. This creates two spectral windows or spectral regions.

Radio Astronomy: Data Reduction

Radio Astronomy: Data Reduction

Radio Astronomy: Data Reduction To reduce JCMT HARP data we use the Starlink software. (Make sure you have the libg2c.so library for linux users)

First download the appropriate file for your system (32bit; 64bit; OSX) from:

http://starlink.jach.hawaii.edu/starlink/KapuahiDownload

Download (or ask me for) the libexpat-??.so.0 file (for linux systems) and rename it to libexpat.so.0

Make a working directory i.e.

>mkdir HARP

Move the starlink software package (.tar.gz) to this directory and enter it too

>mv starlink-kapuahi-Linux-??.tar.gz HARP >cd HARP>ls you should see the starlink*.tar.gz file→

Extract the software from the tar.gz file

>tar -zxvf starlink-kapuahi-Linux-??.tar.gz

Wait until it finishes unpacking

Radio Astronomy: Data Reduction Move the libexpat file in the star-kapuahi/lib directory

>mv file/path/libexpat.so.0 star-kapuahi/lib/

Let's fix the bash file.

>cd>gedit .bashrc

If black write on top:

# ~/.bashrc: executed by bash(1) for non-login shells.

In the end of the file write (MAKE SURE THE SPACING IS AS HERE):

if [ "$STARLINK_DIR" == "" ]then export STARLINK_DIR="/home/user/HARP/star-kapuahi" source $STARLINK_DIR/etc/profilefi

Save and close the file.

Congratulations, you should now have the starlink software installed and ready to use.

Let's test this, shall we?

Radio Astronomy: Data Reduction Move the data (that I will give you) to the HARP directory in a new working directory:

(in HARP dir do:)>mkdir data

>mv /file/path/data.tar.gz file/path/HARP/data/

Extract the data:

>tar -zxvf harp_lecture_data.tar.gz

You should now have files called a2010004_000??_01_000?.sdf.gz & a20100920_000??_01_000?.sdf.gz

Unzip the files:

>gunzip *gz

And we are now ready to start.

Load the starlink packages we will need (kappa to begin with)

>kappa

You can use kaphelp to see all the scripts/procedures that are there (a lot).

Radio Astronomy: Data Reduction We start from the raw time ordered data (i.e. file contains the data as they were taken in time, rather than how they should be in space)

Take a look to one of this files. Remember this data is a cube, i.e. 3 axes (x,y,spectrum)

>gaia a20100920_00035_01_0001.sdf

We need to check if there are any power spikes (e.g. from the instruments behaving weirdly). To do this we need to collapse the data in 1D:

>collapse a2010????_000??_01_000?.sdf axis=3 estimator=rms out=temp1.sdf>collapse temp1.sdf axis=2 estimator=rms out=spike_??_00?.sdf

Again with gaia display the new image and check if there are any spikes (I will show you how).

Do the same for the rest of the files.

Once that is done we may or may not need to remove spikes.

Next we need to identify bad receptors. For this we need the actual map in proper latitude-longitude-velicity axes.

>makecube “a2010????_000??_01_000[1-4].sdf” out=temp_?? outcat=temp_32.FIT autogrid=false specbounds=\"-60.0 60.0\" pixsize=7.3

Radio Astronomy: Data Reduction Again use gaia to determine the bad detectors

>gaia temp_??.sdf

Go to Data_Servers Local Catalogues Load from file→ →

Choose the temp_??_1.FIT

In the new window go to Options set plot symbols→ and change the symbol/size/colour are you prefer. Now check which of these receptors are on bad data (brighter stripes of one pixel width usually) by clicking on the symbols and seeing the name of the receptor in the other window. Note it down.

We need to remove the bad receptors so in the time ordered data we do:

>makecube “a2010????_000??_01_000[1-4].sdf” out=a2010????_??.sdf autogrid=false specbounds=\"-60.0 60.0\" pixsize=7.6 detectors='"-H??"'

Next we need to baseline the data (i.e. make sure that all spectra have noise around 0)To do this we fit the edges of the spectrum (in the velocity axes, again use gaia to open the new cubes we made and click on a spectrum to see where there is signal and where not. We want spectrum chunks without signal).

>mfittrend a2010????_??.sdf ranges=\"?? ?? ?? ??\" order=3 axis=3 out=baseline.sdfsub a2010????_??.sdf baseline.sdf out=ba2010????_??_3.sdf

Radio Astronomy: Data Reduction Congratulations! You now have a chunk of the map ready for analysis.

Repeat these steps for as many map segments as you have and then join the maps:

wcsmosaic in="a2010????_3*.sdf" out="12co_??_2010????.sdf" acc=0.001 method=bilinear ref="" LBND="" UBND=""

And now you are done for real with data reduction.

Well.unless you notice something wrong in the reduced maps in which case you need to figure out what went wrong and do it all again.

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