Chapter 2
Pulsar searching
“In every branch of knowledge the progress is proportional to the amount of facts on which to build, and therefore to the facility of obtaining data.”
James Clerk Maxwell
While the fundamental principles behind pulsar searching have not changed for the past 40 years, the techniques, algorithms and hardware available to perform such searches have advanced significantly. The ever increasing power and decreasing cost of comput-ing resources, in combination with advances in data storage and transport techniques, have allowed modern-day pulsar surveys to explore the sky at much higher time and frequency resolution than has previously been achievable. In this chapter we describe the instrumentation and algorithms used in a modern-day pulsar search, namely the High Time Resolution Universe North survey (see Chapter 4).
2.1 Instrumentation
Pulsars are very weak radio sources, with observed flux densities in the Jy toµJy range at 1.4 GHz (Manchester et al., 2005). While it is possible to observe pulsars using small radio telescopes and long integration times, to achieve sensitivity to the weakest pulsars within reasonable timeframes1, we require telescopes with large collecting areas. The high instantaneous sensitivity of such systems is also vital for the observation of weak radio bursts of astrophysical origin, a major aspect of current pulsar surveys (see Section
1Being able to achieve high signal-to-noise ratios in short observations is important in simplifying the task of searching for pulsars in binary systems (see Section 2.2.4.7).
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Figure 2.1: Schematic representation of a generic single-dish pulsar observing system.
2.2.6). Several large telescopes exist that are well suited to pulsar search observations.
These include the 100-m Effelsberg radio telescope, the 64-m Parkes radio telescope, the 105-m Robert C. Byrd Green Bank Telescope (GBT) and the 300-m Arecibo radio Telescope, all of which are currently performing large-scale pulsar searches (see Table 1.1).
Figure 2.1 shows a schematic of a generic observational setup for single-dish pulsar observations. In the following sections, we will review each stage of the signal chain.
Pulsar searching 31 2.1.1 Frontend
The term frontend generally refers to those parts of the observing system which deal with the signal at the incoming radio frequency (RF). For most observing systems, the frontend is composed of a feed horn, low-noise amplifier, bandpass filter and mixer (L´ena et al., 2012).
Radiation focused by the dish onto the receiver enters the feed horn, where standing waves at select RFs are allowed to form. The electric fields created by these standing waves are then sampled by a probe in two orthogonal polarisations (see e.g. Burke and Graham-Smith, 2009). As the signal paths for the two sampled polarisation senses (usually orthogonal linear or dual circular) are identical, we shall from here on refer to only a single signal path.
After being sampled by the probe, the weak radio signal is amplified in a low-noise amplifier (LNA) with a specific frequency response centred on the middle of the observing band. This amplification stage represents the major noise contribution in high-frequency radio observations (&1 GHz)2. To limit the thermal noise in the system, the LNAs and the probe are usually housed in a cryogenic Dewar flask which is kept at temperatures of the order of a few tens of Kelvin. The amplified signal is then passed through a filter that suppresses known strong, persistent radio-frequency interference (RFI) signals in the band.
The last stage in the frontend involves frequency conversion from the incident RF to a lower intermediate frequency (IF). Frequency conversion is achieved through the process of heterodyning, wherein the RF signal is ‘beat’ with a lower-frequency signal from a local oscillator, in a device known as a mixer. As the characteristic impedance of a transmission line decreases with decreasing frequency (see e.g. Huba, 2011), the drop in frequency acts to reduce losses during transmission of the signal from the frontend to where the backends are housed.
Between the frontend and the backend, the IF signal undergoes a final amplification stage. Like the RF amplifier the IF amplifier also has a specific frequency response range, which is the width of the desired band. Before entering the backend, the signal is passed through an equaliser to remove any frequency-dependent artifacts caused by propagation effects in the cabling used to carry the signal from the frontend.
For more information on frontend systems for radio telescopes, see e.g. L´ena et al.
(2012).
2At lower observing frequencies, the major noise contribution comes from Galactic synchrotron emis-sion. For this reason the true switch frequency between the two regimes is dependent on Galactic longitude and latitude.
Pulsar searching 32 2.1.2 Pulsar search backend
The term backend, refers to the end of the signal chain, in which the signal is digitised, processed and stored. Here we will specifically look at the type of backend that is most commonly used in modern pulsar search observations, the polyphase filterbank.
As the IF signal enters the backend, it is Nyquist-sampled and digitised using an analogue-to-digital converter (ADC). The ADC maps out the IF signal to an n-bit number with nano-second-scale sampling before passing the digitised data stream to a digital filterbank.
In the simplest terms, the digital filterbank contains a field-programmable gate array (FPGA) which performs fast Fourier transforms (FFT, see Section 2.2.4.2) of discrete blocks of data to produce a spectrum every few microseconds. Spectral leakage in the FFT stage can cause strong signals at specific frequencies to spread their power over several adjacent Fourier bins in the power spectrum. To reduce the effects of spectral leakage, modern-day backends employ polyphase filtering techniques (see e.g. Lyons, 2010) . These techniques purify the signal response by weighting overlapping sections of data with a sinc function pre-FFT. The combination of polyphase filter and FPGA is referred to as a polyphase filterbank3.
The length and sampling rate of the data undergoing FFT are determined by the number of filterbank channels desired and the overall bandwidth. To achieve a 512-channel filterbank over a bandwidth of 300 MHz, would require a 1024-point FFT of data sampled at 600 MHz, i.e. a 1024-point FFT every ∼1.7 µs. Such a high rate of sampling is not desired for search observations, as we only require a sampling rate two or more times higher than the frequency of the fastest pulsar for which we are searching. This means it is theoretically possible to detect a pulsar with a spin frequency of 1 kHz using a sampling of only 2 kHz. In reality we require a larger number of samples across the pulse period, as the period isa priori unknown and any strong blue-noise4 components in the data can potentially mask out the pulsar signal. For this reason the spectra produced in the polyphase filterbank are integrated to give sampling rates of several tens of microseconds, more than enough to detect the fastest pulsars known.
As we do not require any polarisation information for the purposes of pulsar searching, a quadrature sum of the output from each polarisation channel is performed.
In the final step performed by the backend, the data are formatted and the resultant files are written to disk for processing.
3A short, useful guide to polyphase filterbank techniques can be found athttps://casper.berkeley.
edu/
4i.e.noise which is stronger at higher temporal frequencies.
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