Enhancing the Sensitivity of the EHT
Data Collection at Wide Bandwidths
One way to increase the sensitivity of the EHT is to capture more energy from the black hole targets at each EHT site. Since black holes emit radiation at many frequencies, we can do this by increasing the range of frequencies that are recorded during EHT observations. This, in turn, requires electronic systems and recording systems that operate at higher speeds. Industry trends that allow faster personal computers and higher capacity hard disk drives have enabled the EHT to leap forward to recording rates that are more than a factor of 10 faster than for any other global array. This is embodied in “Moore’s Law”, a heuristic coined in 1965 by Intel co-founder Gordon Moore, has predicted the exponentially increasing power of integrated circuits for the subsequent decades.
The effect of Moore’s Law has enabled the EHT to gather, record, and process much larger bandwidths at a fraction of the cost of earlier pioneering VLBI systems. The resulting increase in observing sensitivity has helped extend the EHT’s reach to longer baselines, and resulted in higher quality data sets with much better “signal-to-noise” ratio, or SNR.
The EHT equips each single dish site with specialized electronics designed and supplied by the collaboration. Though historically, analog VLBI equipment was used, in the modern era digital electronics is prevalent and has been the mainstay of the EHT. For single dish telescopes, the primary unit is called the VLBI “Digital Back End”, or DBE, which samples analog data from a radio receiver and feeds the formatted digital data to a data recorder.
Several different types of digital backend have been used in EHT observations, including the first-generation DBE1 system, the Digital Base Band Converter (DBBC) system, developed in Europe, and the ROACH Digital Backend (RDBE). The most recent incarnation is called the “R2DBE” or “ROACH2 DBE”, and has been deployed at all EHT sites. The R2DBE samples and processes data at a rate of 16 gigasamples-per-second, perfectly matched to the recording data rate of the Mark6 digital recorder, the latest generation of EHT VLBI Data Recorder. ROACH stands for “Reconfigurable Open Architecture Computing Hardware” and is shared by an open source astronomical instrument collaboration called “CASPER” the Collaboration for Astronomy Signal Processing and Electronics Research”.
Each Mark6 recorder receives digital data at a rate of 16 Gigabits/sec from the R2DBE and distributes it among a total of 32 hard disk drives grouped into 4 modules of 8 disks each. The EHT is scheduled to record an aggregate rate at each site of 64 Gigabits/sec by using 4 Mark6 units in tandem. This rate is matched to the maximum bandwidth current available from the key ALMA site (Atacama Large Millimeter/Submillimeter Array) that has the largest collecting area of all the EHT sites.
Recorded disk packs from each site are shipped back to two central locations, the Max Planck Institute in Bonn, Germany, and the MIT-Haystack Observatory in Westford, Massachusetts, for correlation. The DiFX, or “distributed F-X” software correlator is now used for EHT correlation. Among other advantages, software correlation clusters are scalable and the programs are easily customized. CPU-based processors are commodity products so in the processing domain as well as the recording the EHT take’s advantage of Moore's Law advances in processing power.
Increase In Telescope Aperture
The most straightforward way to boost the sensitivity of the EHT is to increase the net collecting area of the dishes in the array. New telescopes can be added— the 12m diameter Greenland Telescope, for example, is due to come on line in 2018— but larger dishes are especially valuable since collecting area scales as the square of the dish diameter. A larger collecting area means more photons emitted by hot gas near the black hole event horizon can be captured on the Earth. The Large Millimeter Telescope (LMT), an EHT station in Mexico, is 50 meters in diameter, making it the largest fully steerable millimeter/submillimeter wave telescope in the array.
But building large reflectors is an expensive and sometimes impractical proposition, especially at these short wavelengths, because the mechanical precision and rigidity of the dish has extremely tight tolerances, which are hard to meet.
Some of the EHT sites, such as ALMA, the SMA, and the planned IRAM NOEMA are themselves collections of smaller antennas. Like the EHT they are interferometers, but they operate on local short baselines up to hundreds of meters, rather than thousands of kilometers as for the EHT, and their dishes are connected by cables. To use these sites as EHT stations, the small dishes must be electronically phased together, which allows their collecting area to be combined. Phased ALMA, for example, combines up to 64 dishes, each with a 12m diameter, for a total collecting area of 7200 square meters, which is about three times the 2000m collecting area of LMT.
The measure of the operational quality of a phased array system is the “phasing efficiency” , which ideally should be greater than 90% for typical EHT sources. The plot here shows the phasing efficiency of SWARM over a track during the 2016 EHT campaign. Credit: Primiani et al., 2016, JAI, Vol 5, No. 4 (2016), page 13.
Improving the Resolution of the EHT
Building A Larger Array
An array of telescopes becomes able to discern finer and finer features in the emission as the separation between the telescopes is increased. We have been adding new telescopes to the EHT to build an array with a footprint approaching the size of the Earth.
The Schwarzschild radius for Sgr A* is 10 microarcseconds, an exceedingly small size even by astronomical standards. EHT observations to date have achieved a resolution of better than 60 microarcseconds— about the angular size of an orange on the moon. We are working to include other millimeter telescopes into the EHT array in order to improve the resolution of the EHT and eventually be able to produce images of the black holes in Sgr A*, M87, and other sources.
The angular resolution of a baseline is given by λ/D, where D is the projected separation between the antennas. We also refer to this calculation by λ/B to specify that D is the Baseline, the separation of EHT sites used for Very Long Baseline Interferometry. Finer angular resolution can be obtained by observing at shorter wavelength and by increasing the distance between telescopes. Therefore, we actively attempt to get the best resolution, gained from our ability to resolve miniscule objects with smallest resolution achievable.
We have obtained detections to date on baselines between Hawaii (SMA + JCMT), and the continental United States (LMT + SMT), with an angular resolution corresponding to 6 Schwarzschild radii for Sgr A*. ALMA will double this angular resolution. Future EHT observations may be able to obtain a resolution as fine as 1.5 Schwarzschild radii.
|Baseline||Resolution at 230 GHz||Resolution at 345 GHz|
|LMT - SMT||140 μas/ 14 RSch||93 μas/ 9.3 RSch|
|Hawaii - SMT||58 μas/ 5.8 RSch||39 μas/ 3.9 RSch|
|Hawaii - ALMA||28 μas/ 2.8 RSch||19 μas/ 1.9 RSch|
|Plateau de Bure - South Pole||23 μas/ 2.3 RSch||15 μas/ 1.5 RSch|
Table showing the resolution (both in Microarcseconds and Sgr A* Schwarzschild Radii) achievable on Sgr A* with current and future EHT baselines.
The improving uv coverage as telescopes are added EHT array.
Image fidelity: As more telescopes are added to the EHT, we will be able to produce images of the emission around black holes. In general, the fidelity of images produced by an interferometric array increases as additional telescopes are added to the array.
Radio astronomers use the term "uv coverage" to refer to the projected baseline lengths and orientations for which data are obtained. The east and north projection of each baseline as measured in units of the observing wavelength are referred to as "u" and "v", respectively. As the Earth rotates, the projection of each baseline in the plane normal to the direction to the source changes such that each baseline sweeps out an arc in the uv plane. Each location in the uv plane corresponds to one Fourier component of the image on the sky. The ability to reconstruct the sky image improves with increasing uv coverage.
An additional consideration when scaling a project to include a wide array of telescopes is the weather and observation variability for each observatory. Having sites in deserts, tropics, mountains, and ice fields reinforces the need to account for climate differences, which we continue to monitor and analyze to find key times throughout the year for optimal visibility. Here, we take key measurements of atmospheric opacity and transmittance to find correlations among all our array sites, thus discerning the best observation periods when we can most effectively gather light with all our telescopes. Opacity and transmittance measure the ability of electromagnetic signals to pass through the atmosphere and arrive at our receivers while still encoding the desired information. At high atmospheric opacities, the molecules in the air diffract (change the direction of photon paths) the light so that the original signal decays. Therefore, we continue to monitor weather patterns to minimize the signal decay and observe the most complete wave.
As we build a larger array, more challenges and considerations arise, but with them greater opportunities to increase our baseline, and resolve the structure of black holes.