High resolution/high frequency High resolution/high frequency radio interferometryradio interferometry

Anita RichardsUK ALMA Regional Centre

Jodrell Bank Centre for Astrophysics University of Manchesterthanks to fellow tutors, ALMA and JBCA colleagues

and IYAS organisersand IYAS organisers

Outline● Dish-based arrays – sub-mm to m

– Earth Rotation Aperture Synthesis● High-frequency considerations

– Field of view– Brightness temperature - observability– Correlation and spectral windows– Atmosphere

● Calibration– Instrumental measurements– Astrophysical standards and phase referencing– Self-calibration– Mostly ALMA examples

● Principles same for NOEMA– Differences in techniques summarised

● Data products

Dish-based arrays: (sub-)mm

(a) Steerable dishes(b) Subreflector(c) (Multiple)

receivers cabin(d) Optical fibre links

to correlator(e) Reconfigurable

(b)(a)(d)

(c)

(e)

Also SMA ≿0.45 mm; VLA ≿6 mm; ATCA ≿3mm; KVN ≿2.5mm to to 0.8 mm 0.8 mm

to to 0.35 mm0.35 mm

(e)

(sub-)mm windows & bands

99 5 5 33 2 2 1.61.6 1.2 0.9 1.2 0.9 0.6 0.6 0.45 0.45 0.350.35 (mm) (mm)

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NOEMANOEMA1 2 3 4 1 2 3 4

ALMA

cm-wave arrays: connected, ~open access(a) e-MERLIN (UK); 217 km; 1.3 – 22 cm

- highest resolution(b) Westerbork (NL); 2.7 km; 3.5 – 92 cm

- largest field of view (c) Australia Telescope Compact Array;

6 km; 27 – 0.3 cm- shortest wavelengths

(d) VLA (USA); 36 km; 0.6- 22 cm;- most sensitive at short

(e) Giant Metrewave Radio Telescope (India); 25 km; 18 – 500 cm

- most sensitive at long (+ dipole etc. arrays e.g. LOFAR, MWA)

217 km

& many more

(a)

(b)

(c) (d) (e)(e)

● Link single dishes & arrays – And in space

● cm-wave now xtending to mm– KVN, VERA, VLBA, new EVN ...

● Global (sub-)mm VLBI with ALMA and NOEMA– Event Horizon telescope

RadioAstron SatelliteVery Long Baseline

interferometry-arcsec resolutionOften recorded pre-correlaton

Coming: Coming: MeerkatMeerkatSKA SKA + AVN!+ AVN!

Point source overhead

Correlator

● Signals in phase, interfere constructively– Combined phase 0o, amp constant

Resolved source overhead

– Combined depends on s– Complex visibility amplitude is sinusoidal function of

● Signals from overhead vector s0

● Signals from lobe– Angular offset vector s– Path length s = s ∙ s0

Earth rotates

● Telescopes separated by baseline Bgeom

● Earth rotates– Projected separation

b = Bgeom cos 0

● Samples different scales of source

● Additional geometric delay path – Remove in correlator

Earth rotation aperture synthesis

– Combined depends on ss(time)– Complex visibility amplitude is sinusoidal function of

Interferometry to imagesl

m

● Sensitivity: – Number antennas – d freq. width per image, t total time on source

u

vCorrelator

bb

σrms∝T sys

√N (N−1)/2 δ ν Δ t N pol

● Correlation makes complex visibilities V(u,v)– One V, Aei, per -chan per baseline per int per pol.

● Fourier transform gives sky brightness distribtion∑V(u,v) e[2i(ul + vm)] dudv = I(l,m)

or V(u,v) ⇔ I(l,m)* for short

*I(l,m) can bewritten T(x,y)

High-frequency considerations● Same principles as any radio interferometry● You are unlikely to be bothered by:

– Ionosphere (delayionosphere ∝ 2)

– Confusion (PrimaryBeam~/B ~ 55” @ 3mm, 12-m dish)

● Most bright extragal. sources have <0 where S ∝

● You will suffer from:

– Small field of view (PrimaryBeam ~9” @ 0.45 mm)

– Tropospheric refraction (troposphere ∝ 1/ )● phase affected, amps if signal decorrelates

– Tropospheric absorption and emission● amplitudes affected

● Same principles as any radio interferometry● You are unlikely to be bothered by:

– Ionosphere (delayionosphere ∝ 2)

– Confusion (PrimaryBeam~/B ~ 55” @ 3mm, 12-m dish)

● Most bright extragal. sources have <0 where S ∝

● You will suffer from:

– Small field of view (PrimaryBeam ~9” @ 0.45 mm)

– Tropospheric refraction (troposphere ∝ 1/ )● phase affected, amps if signal decorrelates

– Tropospheric absorption and emission● amplitudes affected

High-frequency considerationsVLA beam

18cm ~25'

ALMA primary beam 0.45mm ~9”; 3mm ~55”

5 arcmin

Estimate observability from low-res image

• Brightness temperature ightness temperature TTbb = = SSsourcesource 10 10-26-2622 / 2 / 2kkBB – emitting area (sr),emitting area (sr), (m), (m), S S (Jy)(Jy)

• Resolved? Use Resolved? Use SS per measured per measured • Unresolved by low-resolution obs. e.g. single dish?Unresolved by low-resolution obs. e.g. single dish?

– Estimate actual source size forEstimate actual source size for

– Use Use SS per best estimate of per best estimate of to find to find TTbb

– Predict ALMA flux density Predict ALMA flux density SSALMAALMA = = TTbb 2 2kkBB bb22 / 10 / 10-26-26 22

• Substitute Substitute == bb22 ( (ALMA synthesised beamALMA synthesised beam))

– Use Use Sensitivity Calculator Sensitivity Calculator - n- need >5eed >5rmsrms on peak on peak

• e.g. total source 1 Jy, e.g. total source 1 Jy, √√=15”, =15”, TTbb~0.07 K at ~0.07 K at 1mm 1mm

• Brightness temperature ightness temperature TTbb = = SSsourcesource 10 10-26-2622 / 2 / 2kkBB – emitting area (sr),emitting area (sr), (m), (m), S S (Jy)(Jy)

• Resolved? Use Resolved? Use SS per measured per measured • Unresolved by low-resolution obs. e.g. single dish?Unresolved by low-resolution obs. e.g. single dish?

– Estimate actual source size forEstimate actual source size for

– Use Use SS per best estimate of per best estimate of to find to find TTbb

– Predict ALMA flux density Predict ALMA flux density SSALMAALMA = = TTbb 2 2kkBB bb22 / 10 / 10-26-26 22

• Substitute Substitute == bb22 ( (ALMA synthesised beamALMA synthesised beam))

– Use Use Sensitivity Calculator Sensitivity Calculator - n- need >5eed >5rmsrms on peak on peak

• e.g. total source 1 Jy, e.g. total source 1 Jy, √√=15”, =15”, TTbb~0.07 K at ~0.07 K at 1mm 1mm

– bb=0”.15, flux density 1x(15/0.15)=0”.15, flux density 1x(15/0.15)22 = 0.0001 Jy/beam = 0.0001 Jy/beam

Estimate observability from low-res image

Correlator configurations● Two sidebands, fixed spacing depending on band

– e.g. B7, B3 sideband centres separated by 12 GHz● Max. continuous b/w <4 GHz (8 GHz Bands 9, 10)

332 333 GHz 343 345

338

● See documentation and OT for full details

● Four basebands (BB), max. width 2 GHz – 128 chans per BB (dual pol) TDM– 4096 chans per BB (~0.5 km/s at 300 GHz) FDM

● Useful max. 1.875 GHz (so 3840 channels usable) ● Narrower spectral windows for higher resolution

– Factors of two down to 62.5 MHz (15.25 kHz chans)– Higher spectral resolution in single polarization

Correlator configurations● Two sidebands, fixed spacing depending on band

– e.g. B7, B3 sideband centres separated by 12 GHz● Max. continuous b/w <4 GHz (8 GHz Bands 9, 10)

332 333 GHz 343 345

● See documentation and OT for full details

● Four basebands (BB), max. width 2 GHz – 128 chans per BB (dual pol) TDM– 4096 chans per BB (~0.5 km/s at 300 GHz) FDM

● Useful max. 1.875 GHz (so 3840 channels usable) ● Narrower spectral windows for higher resolution

– Factors of two down to 62.5 MHz (15.25 kHz chans)– Higher spectral resolution in single polarization

GAP

Absorption and emission● The atmosphere both absorbs the astrophysical

signal, and adds noise

where the source would provide temperature T if measured above the atmosphere and z is the zenith distance

● Same source, same baselines– Raw amps lower at higher Preciptable Water Vapour

T received=T sourcee τatm /cos z+T atm(1−eτatm /cos z

)

PWV 0.25mm

ShortShort baseline baselineLong baselineLong baseline

Am

plit

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PWV 0.6mm

Position per integration, raw IRAM PdBI data (Krips)

Phase errors cause position errors● Averaging over phase fluctuations causes

decorrelation of amplitudes– Visibility V = Voe

i

rms in radians

Lose 9% amplitude for 5o rms ● Fluctuations on time-scales

of few sec: raw data position jitter ● Water Vapour Radiometry

– Measure PWV, each antenna, every ~second– Calculate phase delay

● Apply corrections to all observed data

⟨V ⟩=V o ⟨ei ϕ⟩=V o e−(ϕrms

2 )/2

WVR before & after

PhaseLong baselineLong baselineShort baselineShort baseline WVR correctionsLongLong Short Short

RAW PHASE

WVR CorrectionsWVR-Corrected PHASE

Measure total noise Tsys

● Measure Tsys using hot & cold loads– ALMA every few min, 128

channels per 2-GHz baseband

– NOEMA per baseband● Tracks atmospheric &

instrumental fluctuations– Applied during data

reduction

Tsys corrections applied

Raw data, coloured by baseline Tsys correctionbefore & after

Tsys corrections

Frequency

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Phase referencing • Observe phase-ref source close to target

– Point-like or with a good model– Close enough to see same atmosphere

● ~2-15 degrees (isoplanatic patch)

– Bright enough to get good SNR much quicker than atmospheric timescale

● 10 min/30 s short/long B & low/high – Nod on suitable timescale e.g. 5:0.5 min

● Derive time-dependent corrections to make phase-ref data match model

● Apply same corrections to target– Correct amplitudes similarly

• Self-calibration works on similar principle

Astrophysical calibration sources• Use flux standard (e.g. asteroid) to set scale in Jy• Calibrate amplitude and phase as a function of :

– Frequency (delay and BandPass calibration)• Bandpass (mostly instrumental errors) stable for hours

– Time (all calibrators; phase referencing applied to target)• Atmospheric fluctuation scale minutes (after WVR)

• Need signal/noise Scalsource/ant > 15 if possible • ~1/15 radians ~5o phase error, 10% amp decorrelation

– per calibration interval per antenna per polarization

• array is noise in all-baseline data per interval

• e.g. 30 antennas, ant ~ 4 x map noise per interval/pol.

σant (δ t ,δν) ≈ σarray √ N (N−1)/2N−3

Pallas~270 km

Time

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Phase-reference visibilities

Time

Vis

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Corrections for point model

TimeTime

Corrections

Corrected phase-ref visibilities

-30o p

hase

30o

Scatter reducedPhases aligned

close to 0o

-150

o Vis

ibili

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150

o

Raw phase (all baselines)

Time

Corrected phase for point source

Corrections (per antenna)

NOEMA calibration

Frequency

amp

phas

e

● Phase calibration– Fit to series of phases

● Suitable for higher noise situations● Controls for multiple calibrators, jumps...

● RF (bandpass) cal– Phases 20th order– Amps 12th order

● Bandpass and phase-ref sources– Often too low S/N for self-cal

Time

phas

e

Effects on imaging

No astrophysicalcalibration: no source seen

Bandpass, phase-refetc. solutions applied:S/N 20Anti-symmetric artefacts dominate- residual phase errors

Phase-only self-cal:S/N 40

Residual, symmetricamplitude errors

Amp & phase

self-cal:S/N 50

VectorsVV isibility = f(u,v)I I mage to be calculated

AAdditive baseline error

ScalarsScalars

SS (mapping I to observer pol.)

l,m image plane coordsu,v Fourier plane coordsi,j telescope pair

Calibration: Measurement Equationbasis for calibration in CASA

VVijij = MMijijBBijijGGijijDDijij∫∫EEijijPPijijTTijijFFij ij S S IInn (l,m)(l,m)ee-i-i22(uijl+vijm)(uijl+vijm)dldm dldm +AAijij

Jones MatricesMMultiplicative baseline error

BBandpass response

GGeneralised electronic gain

DDterm (pol. leakage)

EE (antenna voltage pattern)

PParallactic angle

TTropospheric effects

FFaraday rotation etc.

Polarization

Diagram thanks to Wikipaedia

Most (sub-)mm Rx use linear polarization X and Y combined in correlator

ALMA Band 3 (~3mm) uses Ortho Mode Transducer

Band 9 (~0.45mm) uses wire grid

XX, YY for total intensity

XX, YY, XY, YXused to provide Stokes IQUV

Parallactic angle

Rea

lVXY VYX

VXX

VYY

ALMA polarization calibration● Calibration source has unknown linear polarization

– Gain solutions decomposed into X and Y per antennagx' = gx(1+Q/I)0.5 gy' = gy(1 – Q/I)0.5

● Leakage between X and Y feeds ('D'-terms)● Feeds rotate on sky as alt-az antenna tracks source

– Parallactic angle rotates– Qobs time-dependent

● 3+ scans, >3hr HA coverage– Solve for leakage and

source polarization● Known feed orientation

– Directly correct pol. angle ● For high pol. purity correct 2nd

order terms– Need source model (external or self-cal)

ALMA data processing● Each SB self-contained

– Pointing, flux scale, BP, (pol.) cal scans

– ~20-40 mins target(s)/phase ref

– Execute (EB) until sensitivity is reached

● same spectral & array configuration

● Initial data processing (pipeline or staff script)– Convert ASDM to

Measurement Set– Calibrate

● Combine EBs for target imaging

What ALMA data do you get?

Images/cubes for principalscience target channels

Data processing/ pipeline scripts so you could tweak imaging/ self-cal using:

Processingsummary

ASDM (one per EB)

Flag tables

Calibration tables✉ ✉ ✉ ✉ ✉ ✉ ✉ ✉

Science products & scriptsScience products & scriptsmight be all you need

or you could useCASA scripts to regenerate CASA scripts to regenerate edited, calibrated edited, calibrated uvuv data data

• Self-calibrate if bright enough (S/N >~20 per scan)• Re-image

– Weighting for higher spatial sensitivity/lower resolution or v.v.– Change spectral resolution, make spectral index image etc. etc.....

Observing support• www.almascience.org

– Observing Support Tool for quick simulations– Splatalogue for line frequencies etc.– Sensitivity Calculator– Observing Tool for preparing proposals, observations– Archive including public-domain data– On-line Helpdesk– ALMA Regional Centre Nodes face-to-face support

• Community days, schools• RadioNet* MARCuS support to visit nodes for data redn

• http://www.iram.fr/IRAMFR/– IRAM Interferometry & single dish schools

• Biennial, alternate – next mm-dedicated ERIS 2016– RadioNet* TNA support for observations

* No RadioNet support in 2016

(sub-)mm Interferometry Summary• Atmospheric refraction/absorption dominates quality

– Cold dry sites OK ≲370 GHz, exceptional sites ≲1 THz– Troposphere affects phase & amp on ≳1s timescales

• Instrumental calibration (WVR, Tsys) etc.

– ALMA / NOEMA Jy sensitivity, 10 / 100 mas resolution• Good sites, many or large antennas

– Sub-mJy sensitivity at sub-arcsec resolution• Extended sources need multiple arrays/SD fill in• Large fields need mosaicing

• Normally observe two separate sidebands– May have 'mirror' or noise

• Plan night-time, dry season observing at sub-mm• ALMA delivers calibrated data, sample images

– May want to self-calibrate, reimage changing resolution

ALMA Hands-on sessions• Pre-main sequence TTauri star TW Hya

– ALMA ~345 GHz, here dust continuum & line NH2+

– Visibility data have instrumental, bandpass and phase-reference calibration applied already

• List what is in the visibility data set and plot• Make a clean image of TW Hya continuum

– Use this as a model for self-calibration• Make final 'best' continuum image

– Split out calibrated visibility data for TW Hya• Identify line channels and subtract continuum

– Make clean line cube• Image/spectral analysis• Ready-made calibrated data/images available

– Don't worry if you don't quite finish Day One

Day O

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ay One

Day T w

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ay T wo