ovsa expansion pipeline imaging. log-spiral array: “uv” distribution left: sampling of right:...

OVSA Expansion

Pipeline Imaging

Log-spiral array: “uv” distribution

Left: sampling of Right: inner region of spatial Fourier plane sampled spatial Fourier plane

The imaging problem• EOVSA will be imaging over probably the widest range of

spatial scales of any radio telescope: e.g., at 18 GHz, largest scale is 2000”, nominal resolution 3” requires < 1” cells, range is 2000. By contrast, in a typical VLA configuration the range is 20-40; Nobeyama Radioheliograph was about 200.

• I like to think about the problem in terms of the number of pieces of information we measure versus the number of resolution elements in each image: the more information per resolution element, the better the images will be.

• At 18 GHz, there are of order 105 resolution elements if the resolution is 3”, while a 13-element array has 78 visibilities in a snapshot: factor of 1000 short.

Sensitivity

TTA/(N(t))

TA is 10000 K for 2 m dishes

N = 13, = 0.12 MHz, t = 0.02 s

T 20 K

EOVSA configuration

Beam

Imaging mod-els: have 3 frequency-dep-endent comp-onents:

1. Coronal emission derived from an EIT image (2.6” pixels)

2. Chromo-spheric disk derived from Zirin et al.

3. Optically thick gyro-resonance sources derived from an MDI image.

5 GHz

Actual VLA Radio Images

VLA radio images on 1999 April 11. The 5 GHz image is a mosaic of 23 different fields and took all day to make (resolutions 40”/12”)

Imaging deconvolution• The raw product of an EOVSA observation is a set of

“visibilities”: complex numbers representing the spatial Fourier transform of the sky brightness distribution on each baseline between pairs of antennas in the array.

• These spatial Fourier terms can be combined to create an image (the “dirty map”) that represents the sky brightness convolved with the point source response of the telescope (the “dirty beam”). Deconvolution is the process of removing the effects of that convolution.

• Two most commonly used methods in radio astronomy:• “CLEAN”: go through the image sequentially removing

fractions of the dirty beam at the locations of peaks in the dirty image. Represent map as collection of point sources. Works best for compact sources.

• Maximum entropy: determine a model that, when convolved with the dirty beam and differenced with the dirty map, maximizes some measure of entropy (smoothness). Works best for large smooth sources.

Maximum entropy does a poor job on bright compact features (left): better technique is to clean bright emission, subtract it from data, run MEM on residual, then restore bright emission (right).

EOVSA real-time pipeline• Snapshot u,v coverage can image compact flares but not

solar disk• Full 8-hour synthesis can produce a reasonable map of the

full disk (at lower frequencies, probably not at higher)• General idea is to “maintain” full-disk images at each of

some set of frequencies, look for differences in real-time data from these default images

• Procedure for default synthesis images is to first run CLEAN to remove bright compact features, then MEM to get smoother structure

• To make a full synthesis, need to handle solar rotation: this needs a new routine (not currently possible in any package)

• For weak flares you would “subtract” preflare model from visibilities before mapping flare.

• MIRIAD (mostly FORTRAN) is the current deconvolution package in use, need to exploit multi-processor architecture

EOVSA imaging simulation

EOVSA simulations: residuals

EOVSA analysis …• “Default” images serve as our normal real-time image

sequence – make them on some suitable time cadence, completely automated, generate movies, etc

• Assume analysis code is in IDL, calibrated image sequences at multiple frequencies are the default data

• Daily summary “light curves” as dynamic spectra• Users can request image sequences for periods and

frequency range of interest • Flux versus frequency, selected regions, polarization, etc1

• …..

1(Most non-black-belt users will not do much else, eg, very few users do anything quantitative with EIT)

Full FASR Log-spiral array configuration

Log-spiral array: dirty beam

Imaging procedures• Mostly work with 5 GHz model (nominal 4” resolution) using

2048x2048 image with 1.2” cells4096x4096 dirty beam

Model converted from brightness temperature to Janskies for deconvolution steps, then back to K for analysis

Threshold for fidelity measure: 3000 K• For CLEAN, subtract model disk from visibilities first, then

restore at end (can also “fit” disk)• For MEM, use model disk as default image• “Weighting” of the visibilities (we regrid the data onto a regular

grid in order to exploit FFTs: relative weights ascribed to the gridded data can be based on local density, etc) plays an important role in point source response because of the wide range of spatial scales sampled.

• Simulations here in MIRIAD package (CARMA/ATNF/SMA), earlier in AIPS package (NRAO).

Spatial resolution versus frequency• Scattering in the solar atmosphere limits the achievable

resolution: Scattering angle ~ 2

3 arcmin at 500 MHz, 15 arcsec at 1.5 GHz, 2 arcsec at 5 GHz

• At low frequencies (below 2 GHz) arcsecond resolution is not possible. This means that we do not need to design for high resolution at low frequencies.

• This also puts an upper limit on the number of antennas we need: N antennas give N(N-1)/2 pieces of information. Ideally we have one piece of information for each resolution element. E.g., at 1 GHz resolution ~ 25”, 6400 resolution elements fill disk, 113 antennas give this many baselines.

• At high frequencies finances will prevent us from adequate sampling: 1” beam => 4. 106 res el => 2800 antennas

Antenna size

• Need to be able to image the full disk of the Sun from 100 MHz to 20 GHz: range of 200 in frequency

• A single antenna of diameter D meters has a field of view

FWHM = 1030 arcmin / D (m) / freq (GHz)

at frequency f GHz. The Sun is 30 arcminutes across.• At 20 GHz need D < 1.7 m• At 2 GHz need D < 17 m• No single dish can handle the wide frequency range: a 2

m dish is not sensitive enough at lower frequencies.• This drives us to using three separate arrays to cover

the frequency range.

Frequency dependence of solar radio flux and of disk brightness temperature

Solar TB diluted by primary beam size

Imaging measures• The conventional measure used in imaging simulations is

“image fidelity”: (image – model)/model (many people use the inverse; here small numbers mean good imaging fidelity).

• We implement this by looking at:the distribution of

(image – model)/model for all pixels exceeding a threshold in the model

• We calculate the mean of the absolute value over all ratios, the median of the absolute value, and the standard deviation of the differences.

• Eventually we will calculate these measures for different spatial scales within the image.

Coronal magnetography

• Need to cover 2-20 GHz.

• Assume that we have 4 × 500 MHz IFs, each 2 polarizations, covering 2 GHz at a time: need 9 settings to cover frequency range.

• 5 second time resolution would give 0.5 seconds per setting

• 32 MHz spectral resolution is about 10 G magnetic field resolution: roughly 500 channels from 2-20 GHz.

• Need images of high quality to be able to use spatially-resolved spectroscopy.

Synchrotron emission from flares

• Trigger flare mode when flux on compact spatial scales exceeds specified level

• Time resolution of prime importance: bright emission confined to small spatial region.

• If we have 4 × 500 MHz IFs, each 2 polarizations, covering 2 GHz at a time, then 0.05 second integrations at each setting gives 0.5 second resolution.

Solar disk radio brightness temperatures

Disk brightness temperatures: 600000 K at 0.3 GHz - optically thick corona

150000 K at 1 GHz - optically thin corona

20000 K at 3.75 GHz - upper chromosphere + thin corona

13000 K at 9.4 GHz - upper chromosphere

10000 K at 17 GHz - chromosphere

7000 K at 86 GHz - chromosphere

Time sharing and time resolution

• Assuming a correlator with finite instantaneous bandwidth, need different observing modes to handle different science goals.

• This forces trade-offs between spectral coverage and time resolution

Coronal Coronal magnetographymagnetography

High spectral resolution from 2-High spectral resolution from 2-20 GHz, low time resolution, high 20 GHz, low time resolution, high image quality, continousimage quality, continous

Synchrotron in flares Synchrotron in flares (> 3 GHz)(> 3 GHz)

Low spectral resolution, high time Low spectral resolution, high time resolution (1 second), only during resolution (1 second), only during flaresflares

Low frequency radio Low frequency radio bursts (< 1 GHz)bursts (< 1 GHz)

High time and spectral resolution High time and spectral resolution all the timeall the time

Low frequency radio bursts

• To maintain high temporal resolution continuously, need 2 × 500 MHz IFs, each 2 polarizations, constantly dedicated to 100-1100 MHz.

• 1 MHz spectral resolution would be nice, 4 MHz would get most of the science.

• Will probably only have sensitivity to see coherent bursts at this resolution.

FASR Instrument (Antennas)

Three arrays, 3 km baselines for HFA,Three arrays, 3 km baselines for HFA,

6 km baselines for LFA and LPDA6 km baselines for LFA and LPDA

Array Designation

Number of Antennas

Frequency Range

Antenna Size

HFA High Frequency Array

~100 2-24 GHz 2 m

LFALow Frequency Array

~60 0.2-3 GHz 6 m

LPDALog-Periodic Dipole Array

~40 20-300 MHzLog-

dipole