Radio interferometry at millimetre and sub-millimetre wavelengthsbn204/publications/2009/NikolicERIS... · 2018. 12. 14. · Radio interferometry at millimetre and sub-millimetre

Radio interferometry at millimetre and sub-millimetrewavelengths

Bojan Nikolic1 & Frederic Gueth2

1 Cavendish Laboratory/Kavli Institute for CosmologyUniversity of Cambridge

2 Institut de Radioastronomie MillimetriqueGrenoble

ERIS 2009Oxford, September 2009

Rev 33

B. Nikolic & F. Gueth (Cambridge/IRAM) (Sub-)mm Wave Interferometry ERIS 2009 1 / 62

Introduction

Outline

1 IntroductionScientific differences from the cm/m-wave bandObservational differences from the cm/w-wave bandScience examples

2 Current and forthcoming mm and sub-mm arrays

3 Atmospheric effects/other calibration uncertaintiesPhase fluctuationsAmplitude calibration uncertainties

4 Offline calibration/imaging

5 Summary


Introduction

MM/sub-mm bands (ALMA site, 1mm water)The numbers shown are the ALMA band designations (+ band 1 at 30-45 GHz)

Band 2

Band 3 Band 4

Band 5

Band 6

Band 7

Band 8 Band 9

Band 10

0

0.2

0.4

0.6

0.8

1

T xT x

100 200 500 1000

ν (GHz)ν (GHz)B. Nikolic & F. Gueth (Cambridge/IRAM) (Sub-)mm Wave Interferometry ERIS 2009 3 / 62

Introduction Scientific differences from the cm/m-wave band

Some of the fundamental science from other bands

Most science targets ‘cool’ and close to thermal equilibriumRotational lines of molecules, dust continuum, atomic carbon

Emission mechanisms are energetically significant for starformation both on local and galaxy-wide scalesRelatively low opacity except in the strongest molecular linesStrong positive cosmological ‘K-correction’

Continuum from star-forming galaxies does not dim from z = 1 toz ∼ 8









CO emission line ladder in Milky WayFixsen et al. (1999)

Lines show models S ∝ ν4 exp[−E/kT ]









Spectral energy distribution of galaxiesLagache et al. (2005)









Dust extinction model: UV to near-IRDraine (2003)



Dust extinction model: near-IR to mm-waveDraine (2003)









Positive cosmological K-correctionLagache et al. (2005)


Introduction Observational differences from the cm/w-wave band

Outline





5 Summary



Fundamental observational differences from cm/mwave

Small field of viewFiner resolution (yet to be fully realised)High cost per element of the arrayLack of zero-spacing (‘total-power’) and short-spacing informationSky has a small dynamic range, low surface brightness of typicalsourcesMechanical effects on antennas are importantTroposphere gets seriously in the way, ionosphere not importantLarge absolute – but small fractional – bandwidthsCurrent arrays do not have good instantaneous uv -coverage



Contrast vs single dish

Single dish continuum surveys are confusion limited –interferometry essential for really deep surveysMuch easier to integrate down:

Atmospheric brightness fluctuations are rejectedStanding waves are rejectedGain fluctuations of the receivers less important

Better astrometryGood surface brightness sensitivity is expensive


Introduction Science examples

Outline





5 Summary



Identifying the sub-mm galaxies (LH850.02)Younger et al. (2009)

Deep R-band SUBARU image with SCUBA centroid (dashed) and 2-σposition ellipsoid (dotted)



Identifying the sub-mm galaxies (LH850.02)Younger et al. (2009), SMA 890 µm and R-band SUBARU

SMA (colour scale) SUBARU R-band image+ SCUBA beam (dashed) + SMA position (yellow circle)+ SCUBA position (dotted)



Identifying the sub-mm galaxies (LH850.02)Younger et al. (2009), Spitzer and VLA

IRAC 3.6µm VLA 1.4 GHz+ SMA position (yellow circle) + SMA position (yellow circle)



Imaging of CO in a z = 6.42 quasar hostRiechers et al. (2009) [detection of C I also presented in the paper]

CO J 7→ 6 emission (contours) CO J 7→ 6 spectrumCO J 3→ 2 (colour scale)

PdB λ ∼ 3 mm observations



Resolving the [C II] emission in the z = 6.42 hostWalter et al. (2009)

Red− and blue−shifted [CII][CII] emissionContinuum emission

CO(3−2)

[C II] rest-frame wavelength is 158µmThese PdB observations at λ ∼ 1.1 mm


Current and forthcoming mm and sub-mm arrays

Outline





5 Summary



Existing (sub-)mm arrays

IRAM PdB: 6× 15-m antennasFunded by France, Germany and SpainOpen to all EU countries via RadioNet TNA

CARMA: 6× 10.4-m + 8× 6.1-m + 8×3.5-m antennasSMA: 8× 6-m antennasNobeyama Millimetre Array: 6× 10-m antennaSpecialised array for cosmic background: CBI, VSA, DASI, ...



IRAM PdB array



IRAM PdB array – max baseline ∼ 800 m



Current: Sub-Millimetre Array (SMA)



Near Future: ALMA50×12-m+ 12×7-m antennas, currently being commissioned


Atmospheric effects/other calibration uncertainties

Outline





5 Summary



Troposphere

(Sub-)mm radiation ispredominantly affected by thetroposphere→ (Sub-)mm telescopes aresited at high elevations(Mauna Kea, Chajnator, SP),airborne observatories(SOFIA) or spaceLittle effect on polarisation



Two molecular species most significant: H2O and O2

Atmospheric transmission broken down into contributions from: H2O (blue)and O2 (red), and total (black)

0

0.2

0.4

0.6

0.8

1

Tx

Tx

0 200 400 600 800 1000 1200

ν (GHz)ν (GHz)



H2O is not well mixed



Atmospheric transparencyModel of atmospheric conditions at summit of Mauna Kea

Sky Transparency Sky brightness (in K)

0

0.2

0.4

0.6

0.8

1

Tx

Tx

0 200 400 600 800 1000 1200

ν (GHz)ν (GHz)

50

100

150

200

250

300

TB

(K)

TB

(K)

200 400 600 800 1000

ν (GHz)ν (GHz)

Loss of astronomical signal + Additional noise


Atmospheric effects/other calibration uncertainties Phase fluctuations

Atmospheric path fluctuations

Refractive index n 6= 1:

n − 1 ≈10−6[α

Pd

T+ β

Pw

T+ γ

Pw

T 2

]

Pw : Partial pressure of the water vapourT : Temperature of the water vapourFurthermore, the refractive index is a function of frequency (i.e.,the atmosphere is dispersive), especially at sub-mm frequenciesand close to the edges of the bandsHorizontal and line of sight variation in atmospheric propertieslead to phase errors and phase fluctuations



Example of path fluctuationsSMA, Mauna Kea, Hawaii

−750

−500

−250

0

250

500

750

p(µ

m)

p(µ

m)

16.8 17 17.2 17.4 17.6 17.8 18

t (hours UT)t (hours UT)

Measured pathwhile observinga quasar200 m baselineAbout 3.5 mmline-of-sightwaterσφ = 207µm.



Correlation between baseline length and fluctuation

The phase fluctuation measured at 22 GHz at the VLA by observing a quasarfor about thirty minutes. Correlations along one arm of the VLA only shown.

10

20

50

100

Pha

seflu

ctua

tion

(deg

)P

hase

fluct

uati

on(d

eg)

1 · 102 2 · 102 5 · 102 1 · 103 2 · 103 5 · 103 1 · 104

baseline length (m)baseline length (m)



Effect of uncorrected phase errorsFrom simulations in ALMA Memo # 582

Point-source sensitivity Gaussian beam size

0.1

0.2

0.5

1

SS

0.05 0.1 0.2 0.5 1 2 5

φrms (rad)φrms (rad)

0.1

0.2

0.5

1

D(a

rcse

cs)

D(a

rcse

cs)

0.05 0.1 0.2 0.5 1 2 5


‘Decorrelation’ Limit on possible resolution



Effect of uncorrected phase errors on ‘snapshots’From simulations in ALMA Memo # 582

Positional error Fractional flux error

0.001

0.01

0.1

1

√ 〈∆P

2〉(

arcs

ecs)

√ 〈∆P

2〉(

arcs

ecs)

0.01 0.02 0.05 0.1 0.2 0.5 1 2 5


0.001

0.01

0.1

1

√ 〈S2〉−

〈S〉2

/〈S

〉√ 〈S

2〉−

〈S〉2

/〈S

〉

0.01 0.02 0.05 0.1 0.2 0.5 1 2 5


Astrometric errors Flux calibration errors



Correcting the phase fluctuations

(Wait for stable weather)The magnitude of fluctuations can vary by a factor of five

Switch to a quasar and measure the phase, apply to scienceEssentially the same as normal phase calibrationTo be effective v∆t

2 < BCorrection with Water Vapour Radiometers (WVRs)

Measure water vapour along line of sight of each antennaInfer the path fluctuation on one second timescaleCorrect for the resulting phase errors

Self-calibrationSmall field of view, small dynamic range of sky→ only possible inspecialised projectsExample: quasar absorption lines

‘Paired-antenna’ technique: Dedicated antennas continuouslymonitoring a quasar



IRAM PdB 22 GHz WVRs



ALMA 183 GHz WVRsBlue rectangles are the WVR filters

0

50

100

150

200

250

T b(K

)T b

(K)

175 177.5 180 182.5 185 187.5 190

ν (GHz)ν (GHz)



How WVR phase correction worksALMA WVRs + SMA: Channel 3 data from two telescopes

150

175

200

225

250

T B,3

(K)

T B,3

(K)

16.75 17 17.25 17.5 17.75 18




How WVR phase correction worksALMA WVRs + SMA: Channel 3 difference & phase

−1000

−500

0

500

1000

L(µ

m)

L(µ

m)

16.75 17 17.25 17.5 17.75 18


−7

−6

−5

−4

−3

∆TB

,3(K

)∆T

B,3

(K)



How WVR phase correction worksALMA WVRs + SMA: Channel 3 phase prediction and residual

Uncorrected path fluctuation: 157µm RMS

Estimated Optimum

−750

−500

−250

0

250

500

750

p(µ

m)

p(µ

m)

16.8 17 17.2 17.4 17.6 17.8


−200

0

200

∆p(µ

m)

∆p(µ

m)

−750

−500

−250

0

250

500

750

p(µ

m)

p(µ

m)

16.8 17 17.2 17.4 17.6 17.8


−200

0

200

∆p(µ

m)

∆p(µ

m)

Residual RMS 74µm Residual RMS 71µm



WVR correction in practice at the PdB



WVR correction in practice at the PdBExample of point source observation

Turbulent conditions, 4.4 mm precipitable water vapour, long baselinesNo WVR correction With WVR correction

×2.5 improvement in signal/noise


Atmospheric effects/other calibration uncertainties Amplitude calibration uncertainties

Amplitude calibration uncertainties

Fundamental uncertainties1 Atmospheric transparency varies with time and with frequency2 Receiver gain is variable3 Antenna gain is difficult to measure and sometimes variable

Difficult to inject a signal of known strengthQuasars are highly variable at (sub-)mm wavelengthsSolar system bodies (e.g., Mars, Neptune) also variable, but canbe modelled

Accurate models for the radiometric brightness are requiredMay be resolved, especially at sub-mm wavelengths

→ A calibration chain is required


Atmospheric effects/other calibration uncertainties Amplitude calibration uncertainties

Flux calibration chain

Build reasonably stable receiver systemsCalibrate receiver gain using hot and ambient load

every ∼ few to tens of minutesCalibrate atmospheric absorption through a combination of:

Tipping scans (once ∼ 1 hour)Atmospheric models and WVRs (could go as short as ∼ 1 second)Total power atmospheric emission[Quasar observations (once ∼ 3 mins)]

Calibrate antenna gains using primary calibration standardsonce per session – once a yearat short wavelengths only planets may be suitable

Calibrate antenna primary beam shape through directinterferometric measurement


Offline calibration/imaging

Outline





5 Summary



Bandpass calibration

Principle:Frequency-dependence of gain is independent of time

Calibration steps:Observe a strong quasar at beginning of each session(Need high SNR since can not combine the channels)Fit (complex) gain vs frequency

If SNR is high solve for each channel individuallyOtherwise fit a smooth function of frequency

Apply this bandpass solution to all other data in the session



Bandpass calibration: PdB example – amplitude



Bandpass calibration: PdB example – phase



Bandpass calibration: SMA Data + CASAAmplitude – two out of 24 spectral windows shown

Antennas 1&2 Antennas 1&3 Antennas 1&4

-12400 -12350 -12300 -12250 -12200

Channel Velocity (km/s)

0

5

10

15

20

25

Vect

or

Avera

ge A

mplit

ude

-12400 -12350 -12300 -12250 -12200


5

10

15

20

25

Vect

or

Avera

ge A

mplit

ude

-12400 -12350 -12300 -12250 -12200


5

10

15

20

25

Vect

or

Avera

ge A

mplit

ude

Antennas 1&5 Antennas 1&6

-12400 -12350 -12300 -12250 -12200


5

10

15

20

25

Vect

or

Avera

ge A

mplit

ude

-12400 -12350 -12300 -12250 -12200


5

10

15

20

25

Vect

or

Avera

ge A

mplit

ude



Bandpass calibration: SMA Data + CASAPhase – two out of 24 spectral windows shown

Antennas 1&2 Antennas 1&3 Antennas 1&4

-12400 -12350 -12300 -12250 -12200


-150

-100

-50

0

50

100

150

Phase

Avera

ge

-12400 -12350 -12300 -12250 -12200


-160

-150

-140

-130

-120

-110

Phase

Avera

ge

-12400 -12350 -12300 -12250 -12200


-160

-150

-140

-130

-120

Phase

Avera

ge

Antennas 1&5 Antennas 1&6

-12400 -12350 -12300 -12250 -12200


-60

-50

-40

-30

-20

-10

0

10

20

Phase

Avera

ge

-12400 -12350 -12300 -12250 -12200


-100

-90

-80

-70

-60

-50

-40

Phase

Avera

ge



Phase calibration

Principles:Observed phase of a point source at phase centre should be zeroThe phase response of telescope will change very little for smallangular changes on the skyMost causes of errors are antenna-based and independent ofbaseline

Calibrations steps:Observe quasars every 10 seconds to 20 minutesFit (complex) gain vs time to estimate phase variation

If SNR is low and you think phase should be varying slowly, fitsmooth functionsIf SNR is high or there are jumps, use linear interpolation

Apply phase solution to science data



Phase calibration: PdB Example – smooth phasevariation



Phase calibration: PdB Example – jump ignored bysmooth fit



Phase calibration: PdB Example – use higher orderfunction



Phase transfer

Principles:Low frequency receivers are more sensitive, atmosphere moretransparent, telescopes more efficientQuasar spectra often roughly ∝ ν−0.7

→ easier to make phase calibration observations at lowerfrequency

Calibration steps:Observe phase calibration targets at 3 mmScale phase solutions to science bands and apply to data



Phase transfer: PdB Example



Phase transfer: PdB Example – with transferredcorrection



Amplitude/Flux calibration









Imaging

Imaging generally tractable with established techniquesAdding short/zero-spacing one important challenge but nowalmost routine in some systems

Comparison

(sub-)mm observations vs cm/m Science Imaging

Resolution (∼ λ/B) (current arrays) § ©Field of View (∼ λ/D) § ©Few pixels in image (∼ B/D) (current arrays) § ©Many spectral channels © §Low signal-to-noise § ©



Short-spacing informationBelloche & Andre (2004), Class 0 proto-star observations

PdB only – dashed circle is the primary beam



Short-spacing informationBelloche & Andre (2004), Class 0 proto-star observations

PdB + short spacing from IRAM 30-m


Summary

Outline





5 Summary


Summary

Summary

The physics of (sub-)mm emission means observing it allowsunique scienceInterferometers open the possibility of high-resolution and deepobservationsTroposphere has a big effect on (sub-)mm radiation butcombination of excellent sites and new techniques can/will resolvemost of theseMany of the techniques are the same as traditional cm-waveinterferometryIRAM mm-interferometry summer schools: next year


Summary

Resources on the web

Brogan et al: CASA training pages:http://casa.nrao.edu/casatraining.shtml

IRAM MM-Interferometry Summer School:http://www.iram.fr/IRAMFR/IS/presentations.html

Schilke, P: “ Interferometric Calibration & Imaging” http://www.astro.uni-bonn.de/˜bertoldi/wiki/RadioInterferometry


http://casa.nrao.edu/casatraining.shtml

http://www.iram.fr/IRAMFR/IS/presentations.html

http://www.astro.uni-bonn.de/~bertoldi/wiki/RadioInterferometry

http://www.astro.uni-bonn.de/~bertoldi/wiki/RadioInterferometry

Summary

References

Belloche A., Andre P., 2004, A&A, 419, L35Draine B. T., 2003, ARA&A, 41, 241Fixsen D. J., Bennett C. L., Mather J. C., 1999, ApJ, 526, 207Lagache G., Puget J.-L., Dole H., 2005, ARA&A, 43, 727Riechers D. A., Walter F., Bertoldi F., Carilli C. L., Aravena M., Neri R.,

Cox P., Weiss A., Menten K. M., 2009, ArXiv e-printsWalter F., Riechers D., Cox P., Neri R., Carilli C., Bertoldi F., Weiss A.,

Maiolino R., 2009, Nature, 457, 699Younger J. D., Omont A., Fiolet N., Huang J.-S., Fazio G. G., Lai K.,

Polletta M., Rigopoulou D., Zylka R., 2009, MNRAS, 394, 1685


Documents

Radio interferometry at millimetre and sub-millimetre wavelengthsbn204/publications/2009/NikolicERIS... · 2018. 12. 14. · Radio interferometry at millimetre and sub-millimetre