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# Panoramic SETI: on-sky results from prototype telescopes and instrumental
design

Jérôme Maire  Center for Astrophysics & Space Sciences, University of
California San Diego, USA  Shelley A. Wright  Center for Astrophysics & Space
Sciences, University of California San Diego, USA  Department of Physics,
University of California San Diego, USA  Dan Werthimer  Space Sciences
Laboratory, University of California Berkeley, CA, USA  Department of
Astronomy, University of California Berkeley, CA, USA  Franklin P. Antonio
Qualcomm, San Diego, CA, USA  Aaron Brown  Center for Astrophysics & Space
Sciences, University of California San Diego, USA  Paul Horowitz  Department
of Physics, Harvard University, Cambridge, MA, USA  Ryan Lee  Department of
Astronomy, University of California Berkeley, CA, USA  Wei Liu  Space Sciences
Laboratory, University of California Berkeley, CA, USA  Institute of RF- & OE-
ICs, Southeast University, Nanjing, Jiangsu, China  Rick Raffanti  Techne
Instruments, Oakland, CA, USA  James Wiley  Center for Astrophysics & Space
Sciences, University of California San Diego, USA  Department of Physics,
University of California San Diego, USA  Maren Cosens  Center for Astrophysics
& Space Sciences, University of California San Diego, USA  Department of
Physics, University of California San Diego, USA  Carolyn M. Heffner
Astronomy Department, California Institute of Technology, Pasadena, CA, USA
Andrew W. Howard  Astronomy Department, California Institute of Technology,
Pasadena, CA, USA  Remington P. S. Stone  University of California
Observatories, Lick Observatory, USA  Richard R. Treffers  Starman Systems,
Alamo, CA, USA

###### Abstract

The Panoramic SETI (Search for Extraterrestrial Intelligence) experiment
(PANOSETI) aims to detect and quantify optical transients from nanosecond to
second precision over a large field-of-view ($\sim$4,450 square-degrees). To
meet these challenging timing and wide-field requirements, the PANOSETI
experiment will use two assemblies of $\sim$45 telescopes to reject spurious
signals by coincidence detection, each one comprising custom-made fast photon-
counting hardware combined with ($f/1.32$) focusing optics. Preliminary on-sky
results from pairs of PANOSETI prototype telescopes (100 sq.deg.) are
presented in terms of instrument performance and false alarm rates. We found
that a separation of $>$1 km between telescopes surveying the same field-of-
view significantly reduces the number of false positives due to nearby sources
(e.g., Cherenkov showers) in comparison to a side-by-side configuration of
telescopes. Design considerations on the all-sky PANOSETI instrument and
expected field-of-views are reported.

###### keywords:

High-time resolution, transients, Cherenkov shower, Cosmic Rays

##  1 INTRODUCTION

High-speed time-resolution astrophysics offers the unique opportunity to study
extreme environments of diverse objects such as cataclysmic variables,
pulsars, X-ray binaries, and stellar pulsations[1]. Some of the most energetic
and unusual phenomena in the universe, such as high-speed collisions related
to supernovae, blazars, and gamma-ray bursts, can be sources of gamma-rays and
cosmic rays which, when hitting Earth’s atmosphere, produce very brief flashes
of Cherenkov radiation generated by the cascade of relativistic charged
particles[2]. The Cherenkov radiation arrives at the ground in a flash of only
a few nanoseconds duration[3] and can thus be separated from the night-sky
background. Earth’s atmosphere is used as the detection medium by Imaging Air
Cherenkov Telescopes such as the Whipple telescope[4], H.E.S.S.[5], MAGIC[6],
FACT[7], VERITAS[8] and MACE[9] which have demonstrated the potential and
maturity of this detection technique.

High-time resolution instruments are of interest to search for
technosignatures by means of detecting nano- to milli-second light pulses that
could have be emitted, for instance, for the purpose of interstellar
communications or energy transfer. Several programs aimed to search for
technosignatures using optical wavelengths [10, 11, 12, 13, 14, 15] including
near-infrared[16] have been performed. The first optical SETI all-sky
surveys[17, 18] with 0.32 sq.deg of instantaneous field-of-view adopted a
transit observing strategy to cover the sky in 150 clear nights. However,
assemblies of single-aperture telescopes capable of observing different parts
of the sky instantaneously are still needed to survey the entire sky
efficiently and continuously.

The Pulsed All-sky Near-infrared Optical SETI (PANOSETI) experiment[19, 20,
21, 22] is an all-sky observatory project aiming to detect transients that
will cover a wide range of timescales in search for nano- to second pulsed
light signals, across all optical wavelengths. Based upon two assemblies of
$\sim$45 0.46-m Fresnel-lens telescopes equipped with fast low-noise MMPC
(Multi-Pixel Photon Counter) detectors operating in the 0.32–0.85 $\mu$m
spectral range, the PANOSETI instrument provides sufficient sensitivity to
detect petawatt pulsed signals that could have been sent from kiloparsec
distances and beamed toward our direction, that would be distinguishable from
most known astrophysical sources from our perspective. PANOSETI will be
capable of detecting gamma-rays and cosmic rays with energy above some tens of
TeV. Each part of the sky will be observed simultaneously from two locations
for direct detection and confirmation of transients.

We describe PANOSETI telescopes and their required calibrations in Sec. 2. The
deployment of two pairs of telescopes at Lick Observatory and Palomar
Observatory, with the use of a 1-km baseline, is described in the following
sections along with obtained results (Sect.2-5). Considerations of the
production PANOSETI experiment are reported in Sec. 6.

##  2 PANOSETI module

Affordable, lightweight, refractive Fresnel lenses have been used for the
detection of nanosecond optical showers generated by high-energy cosmic-rays
striking the top of the Earth’s atmosphere [23, 24, 25, 26]. Even though these
lenses have moderate angular resolution (a few arcminutes[19]), their
relatively large-collecting apertures ($>$0.4m) and small focal ratios makes
them optimal for a low-angular resolution wide-field survey. Each PANOSETI
module (or “telescope”) focuses the incoming light using a 0.46-m $f/1.32$
Fresnel lens (Orafol SC214) built with concentric 0.5-mm equal-width grooves
which replace the curved surface of a conventional optical lens, acting as
individual refracting surfaces, and bending parallel light rays to a common
focal point. The lens is made of clear optical 1.8-mm thick acrylic material,
with a high transmittance in the optical ($>96\%$ in the 0.3–1.6$\mu$m in
visible and near-infrared bands) with an anti-reflective coating applied to
the lens surfaces improving transmittance by $\sim 2\,\%$. The lens frame has
been designed to accommodate thermal expansion[21] and includes a 2-mm thick
clear acrylic plate to protect the grooved part of the lens against dust and
condensation. The small focal ratio of the Fresnel lens allows formation of
wide field-of-view (9.97$\times$9.97 ∘) images of distant objects at the focal
plane of the instrument where detectors are located. The telescope Point-
Spread Function (PSF) is slightly under-sampled on the optical axis, one-
fourth of a pixel full-width-half-maximum (FWHM), and more extended off-axis
(up to $\sim$1 pixel FWHM in the corner of the field-of-view).

Three baffles of graduated sizes are attached along the inside of the
telescope tube frame to limit reflections from the inner surfaces of the
telescope[27] and to block stray light resulting from the faceted and
discontinuous surface of the lens. The shapes of the baffles are designed
using superellipses, i.e., intermediate shapes between the circular shape of
the lens and the square shape of the detector array. The telescope assembly,
tube, frame, and baffles are painted in a black matte finish to further absorb
unwanted light[27]. The prototype telescopes are oriented using altitude-
azimuthal mounts. Once set, their pointing directions are fixed during an
observation night.

The telescope’s 32x32-pixel detector array is made of four adjacent detectors,
each one subdivided into four adjacent 8x8-pixel MMPC (Multi-Pixel Photon
Counter, Hamamatsu S13361-3050AE-08) detector arrays which operate in the
spectral range 0.32 to 0.85 $\mu$m with a peak sensitivity at 0.45 $\mu$m
wavelength. These silicon photomultipliers (SiPM) are comprised of Geiger-
mode-operated avalanche photodiodes with a high internal gain to enable single
photon detection while featuring low dark count ($<$1 Mcps), high photon
detection efficiency (45$\%$), and excellent timing resolution. Each 3mm pixel
is made of 50 $\mu$m micro-cells detecting photons identically and
independently. The sum of the discharge currents from each of these individual
binary micro-cells combines to form a summed photon output, thus giving
information on the magnitude of an incident photon flux. Each of the $1,024$
detector signals is amplified, pulse shaped, and peak detected using an
application-specific integrated circuit (ASIC) WEEROC-MAROC3A read-out
chips[28, 29] controlled by FPGAs (field-programmable gate arrays) also used
to modify the detector setup, to timestamp frames, and send data over a 1 GB
fiber-optics communications system to a 10 GB network switch, itself relaying
data from all telescopes to the central computer[30]. The timing
synchronization between detectors and telescopes is maintained by a fiber-
connected White Rabbit system[31] providing nanosecond synchronization
accuracy, and absolute timing through the use of a GPS disciplined time-
frequency reference[30].

Highly-stabilized power-supplies have been specifically designed for PANOSETI
in order to bias detectors as well as to power electronic boards, activate the
internal pulser light source and mechanical parts of the telescopes, such as
the focus stage and detector shutter.

In order to examine pulse widths over a large range of time scales (nanosecond
to second), the instrument has two observing modes that can be run
simultaneously, each one covering a different pulse width range. For detecting
pulses shorter than the amplifier shaping time ($<$ 200 ns), the instrument
measures the pulse height of the signal at each pixel (“PH mode”). For time
scales larger than the amplifier shaping time, the instrument counts the
number of photons in each pixel (“imaging mode”) at a programmable frame rate
end exposure time. This mode is useful for detecting transients with pulse
widths larger than 10$\mu$s, as well as for detector calibrations.

The discriminator thresholds (DAC) applied to the pulse height detection can
be modified to change the sensitivity of the detectors to a given minimal
number of photo-electrons (p.e.) per pulse. Fig. 2-a shows the number of
pulses obtained under dark conditions as a function of the DAC discriminator
threshold for two different pre-gain settings in both imaging and PH modes.
The steps in the curves are used to establish a relationship between DAC value
and p.e. thresholds. In PH mode, the ADC intensity of each channel is also
characterized from the identification of p.e. steps in triggered frames at
different DAC thresholds (Fig. 2-b).

Figure 1: PANOSETI detector calibrations: panel (a) represents the number of
pulses counted per second as a function of the DAC threshold for a given
detector channel in both imaging and pulse height modes at two different gain
settings. Individual photo-event (p.e.) levels can be identified and are used
to calibrate the DAC threshold setting. Sweeping the DAC threshold in PH mode
is used to calibrate the ADC pulse height output (panel b) as p.e. levels can
also be identified in triggered ADC values. Panel (c): Histogram of time of
arrival difference while observing the flasher source internal unit to verify
the nanosecond timing synchronization between telescopes. (d) and (e) panels:
number of pulses as a function of DAC threshold for all 64 pixels of a
detector quadrant, before and after gain adjustment, showing the equalization
of pixel sensitivity at a given threshold.

Since the same discriminator threshold value is applied to each different
pixel of a 8x8-pixel array, the sensitivity of each individual channel is
slightly different and is therefore adjusted. The variable gain pre-amplifier
of the MAROC3A allows the gain of each channel to be tuned and is able to
compensate the non-uniformity between detector channels. Fig. 2-d-e show the
pixel-to-pixel variations of sensitivity before and after performing a pre-
amplifier compensation where p.e. steps have been aligned to the same
threshold values. In addition to telemetry metadata which includes detector
currents and voltages, the instrument monitors each detector’s temperature to
adjust the optimal detector biases (with adjustments of 54mV/∘C ) in real-time
to maintain a constant detector gain.

Coincident detection between two or more sites requires precision time
stamping of events at each site. Prior to observations, a light source unit is
used to send flashes on both detectors at various repetition rates to check
the timing synchronization between telescopes. The light pulses from the
flasher are sent through fibers of equal length to both telescopes. A
histogram of measured time of arrival differences between telescopes are
represented on Fig. 2-c showing a $\pm$ 2 ns accuracy in timing
synchronization.

For each telescope pointing an astrometric solution is determined using on-sky
imaging-mode observations of bright stars (V$<$4). Since the telescope
pointings are fixed in time it is convenient to obtain the altitude-azimuth
coordinates for each pixel. Known stellar coordinates are used during the time
of observations to give astrometric reference points before performing a
gnomic projection to obtain the coordinates of each pixel centers[32, 33].
Observations of several stars at different orientations in the field-of-view
are used to generate a distortion solution and plate constants. The
“distortion map” is then used to correct the pixel coordinates with respect to
the gnomic projection. These astrometric calibrations are performed for each
telescope and are then used to determine residual misalignment offsets between
each telescope’s field orientation on-sky.

##  3 Observations

###  3.1 Lick Panograph Experiment

Two identical PANOSETI telescopes have been deployed side-by-side in the
Astrograph dome[34] at Lick Observatory (Fig. 2-left) since January 2020 to
search for transient signals. It also has been used as a test bed instrument
to verify instrumental performance and to characterize false alarm rates. The
pointing directions of the two telescopes were set to be identical to provide
direct confirmation of detected transient signals. Observations with the
PANOSETI telescopes in the Astrograph, coined “Panograph” instrument, are
performed remotely which allows our team to observe and test new software as
the project develops.

###  3.2 Palomar baseline experiment

In order to verify the expected number of false alarms generated by
instrumental noise and Cherenkov showers (see Sec.5 below), successive
experiments using two telescopes separated by 250-m and 1-km baselines were
performed at Palomar Observatory on Aug. 21st and Oct. 9th 2020, respectively.
The long baseline separation introduces a measurable parallax angle
(illustrated on Fig. 2, right-panel) for flashes of light generated in Earth’s
atmosphere and a time arrival offset compared to astronomical sources (see
Sect.4). The two telescopes were deployed in the field using a long 1.2km
duplex fiber for timing synchronization and data transfer. Co-alignment of
telescope pointing direction was performed by adjusting pointing directions
while observing in real-time bright planets (Jupiter and Mars) in imaging
mode.

Figure 2: LEFT: Two side-by-side PANOSETI telescopes pointing in the same
direction have been deployed inside the Astrograph dome at Lick Observatory
(photo courtesy: Laurie Hatch Photography). RIGHT: 1-km baseline experiment at
Palomar Observatory, the two telescopes are located on sites A and B and
pointing at the same star. This baseline introduces a measurable parallax
angle for flashes of light generated at finite distances in Earth’s atmosphere
and a time of arrival difference offset with respect to astronomical sources.

##  4 Experimental design

###  4.1 Time of Arrival difference

Accurate time stamping of detected events at multiple sites enables the
ability to measure the distance of nearby phenomena in Earth’s atmosphere or
low-Earth orbit, e.g., false alarms from atmospheric Cherenkov radiation or
optical transients from satellite glints. For a given baseline AB (illustrated
on Fig. 2-right), the time of arrival difference between two sites depends on
the pointing direction of each telescope and distance of the detected event
between the two sites. For any event C occurring at coordinates
$(\mathrm{lat}_{C},\mathrm{lon}_{C},\mathrm{alt}_{C})$ with respect to the
World Geodetic System of 1984 (WGS84), one can calculate the distances CA and
CB of a event from each telescope considering the known WGS84 coordinates of
sites A and B. We then determine the optical path length difference by
considering the change of refractive index of air along the lines of sight CA
and CB given the changes in pressure and temperature with altitude with
respect to the 1976 Standard Atmosphere. We used the Ciddor method[35] to
calculate the change of refractive index of air along the path as a function
of pressure and temperature (for a 0.5 $\mu$m wavelength, 50$\%$ humidity and
a CO2 concentration of 450 ppm). The expected time of arrival difference
$\Delta t=t_{A}-t_{B}$ is then determined from the calculated optical path
length difference and speed of light in that medium. The time of arrival
difference is negligible for side-by-side telescopes but increases with the
baseline distance. If the baseline is long enough, the difference of $\Delta
t$ between a nearby source and a distant astronomical sources in any pointing
directions becomes measurable considering the timing accuracy of the
instrument. Fig. 3-right represents the difference of $\Delta t$ between a
phenomena occurring at 10km above sea level (typical Cherenkov shower maximal
altitude is 12.8k̃m [36]) and a source at infinity for a 1km baseline at
Palomar Observatory in any pointing directions above 30∘ elevation. The
minimum of this difference (37 ns) occurs at lower elevations for pointing
directions along the baseline direction and a maximum of $\sim$200 ns close to
zenith.

Figure 3: LEFT: Expected parallax angle of a Cherenkov shower occurring at 10
km above sea level for a 1 km PANOSETI baseline in any pointing direction
above 30∘ elevation. RIGHT: Expected difference between time of arrival
difference for a source at infinity and a phenomena occurring 10 km above sea
level (e.g., Cherenkov shower maximum) using the same 1 km Palomar baseline,
in any pointing direction above 30∘ elevation.

###  4.2 Coincidence detection

A time coincidence occurs when both telescopes pointing in the same direction
detect an event above our sensitivity threshold (e.g., above 15 p.e.) during a
time interval $\Delta t_{\mathrm{window}}$. To consider the delays in time of
arrival due the geometrical orientations of the telescopes, the offset in time
of arrival $\Delta t_{\mathrm{offset}}$ for an object at infinity in a given
pointing direction is included, such that a time occurrence occurred when

$|t_{A}-t_{B}-\Delta t_{\mathrm{offset}}|<\Delta t_{\mathrm{window}}$ (1)

For distant astronomical sources, $\Delta t_{\mathrm{window}}$ could be chosen
such that it will discriminate nearby events from the coincidence detection
process. However, the value of $\Delta t_{\mathrm{window}}$ should be large
enough to take into account timing accuracy and uncertainty of the instrument
($\sim$5ns at 30∘ elevation up to $\sim$10ns at zenith for the 1km baseline).

###  4.3 Parallax angle

The Palomar 1 km baseline yields large parallax angles for optical flashes
occurring in Earth’s atmosphere that are observed by two separated telescopes.
The parallax of such an event is the difference in the apparent positions of
the shower maxima on both telescopes, and is measured by the angle of
inclination between those two lines. The parallax angle of an event can be
calculated in the WGS84 system as the angle $\Lambda$ between two vectors CA
and CB, where C represents the location of the shower maximum (see Fig. 2) of
geodetic coordinates (latC,lonC,altC). For any event C, the parallax angle,
i.e. the angle $\Lambda$ between the two vectors CA and CB, can be determined
using

$\Lambda=\mathrm{atan2}(|\textbf{CA}\times\textbf{CB}|,\textbf{CA}.\textbf{CB})$
(2)

The parallax angle of a shower maximum can thus be determined at any pointing
directions and shower altitudes as shown in Fig. 3-left, in the case of the
Palomar 1km baseline for a shower occurring at 10km above sea level. The cap
represented on the figure represents the pointing directions for any elevation
greater than 30∘. The detected shower maximum as seen by the two separated
telescopes will have different coordinates and will be separated by at least
$\sim 2^{\circ}$ (6 pixels) at low elevations in the baseline direction up to
$\sim 7^{\circ}$ (21 pixels) close to zenith.

##  5 Results

We measured the false alarm rate (FAR) at Lick Observatory during a 3 h period
of continuous on-sky observations in PH mode on April 2nd 2020\. Considering
coincident events detected by the two telescopes above 15.5 p.e., we found a
false alarm rate of 0.31 coincidence/s/pair of telescopes (on average 1
coincident event each 3.2 s per pair of telescopes) with respect to the
$9.9^{\circ}$-wide field-of-view. This 15.5 p.e. threshold was chosen to be
above the average highest threshold needed to detect pulses generated by sky
background (5̃ to 11 p.e. depending on lunar illumination). An example of
Cherenkov shower detected at Lick Observatory is represented on Fig. 4. The
spread and elongation of the shower over many degrees can be used to
discriminate this specific event from an astronomical source. However, in the
case of showers developing along the telescope line-of-sight, it is difficult
to differentiate them from other point-like source. About 20 - 30 % of the
detected coincidences could not be rejected on criteria based on the shape or
spread or angular separation. Since the telescopes are located side-by-side,
no rejection based on parallax or time of arrival analysis could be obtained.
As can be seen on Fig. 4, almost all coincidences have a time of arrival
difference between telescopes included in a $\pm$ 10 ns interval.

Figure 4: LEFT: An example of a Cherenkov shower detected at Lick Observatory
with side-by-side telescopes (slightly misaligned by an offset of 1.5 pixels).
This specific shower extends to several degrees and can be distinguished from
a source at infinity. Right: Histogram of time of arrival difference for
detected coincidences above 15.5 p.e. at Lick Observatory, on April 2nd 2020
UT showing that these events are detected within less than 10 ns difference.

The false alarm rate was measured at Palomar Observatory (Table 1) using the
same sensitivity threshold with two telescopes separated successively by three
different baselines: modules side-by-side, separated by a 250 m, and by a
1,012 m baseline. The false alarm rate decreases drastically with baseline
distance since the baseline eventually becomes larger than the typical shower
detection area ($\sim$ 240-m diameter[36]). However, a few rare showers are
still detected with the 1-km baseline. We observed 1 event per 18 min 45 s on
average, using a large $\Delta t_{\mathrm{window}}$ (200 ns) at a zenith angle
of 35∘. All detected events are identified as Cherenkov showers since their
time of arrival differences and parallaxes are consistent with phenomena
occurring between 9 and 12 km above sea level. Fig. 5 shows a typical example
of these events detected by both telescopes pointing on the same field-of-view
with the 1 km Palomar baseline on October 9th, 2020. Due to the long baseline,
the parallax is so large (4.75∘ at shower maximum) that the event is detected
on different quadrants of the two detectors. A 1km baseline decreases the
false alarm rate since the apparent angular separation of the detected shower
may result in showers falling outside of the second-telescope field-of-view.
The time of arrival differences of showers detected with a 1-km baseline are
significantly different from an astronomical source (92 to 141 ns of
difference with a source at infinity in their pointing directions).

Figure 5: LEFT: Example of a Cherenkov shower detected by the two PANOSETI
telescopes at Palomar Observatory using a 1-km baseline on Nov. 9th, 2020. The
images show the triggered events on both detectors and their respective
positions on the entire telescope field-of-view (top). The 4.75∘ offset in
position of the shower maximum is due the parallax that the 1km baseline
provides for nearby objects. Both the time of arrival difference (1913 ns) and
the parallax (4.75∘) are in agreement with an event that occurred at about 10
km above sea level. Table 1: Measured PANOSETI false alarm rates (FAR) at Lick
Observatory (1,284 m altitude) and Palomar (1,680 m & 1,698 m altitudes of
sites A and B) considering coincident events detected by the two telescopes
above 15.5 p.e. (without rejection based on parallax or shape of the event).

|  | Eleva- | Obs. | False Alarm Rate | False Alarm Rate
---|---|---|---|---|---
Location | Baseline | tiona | Duration | [coincidences/s/pair of telesc.] | [coincid./s/pair of telesc.]
|  |  |  | ($\Delta t_{\mathrm{window}}$ =200 ns) | ($\Delta t_{\mathrm{window}}$ =10 ns)
Lick | Side-by-side | 60∘ | 3 h 2 min | 0.31 (1 coinc. per 3.3 s) | 0.3 (1 coinc. per 3.2 s)
Palomar | Side-by-side | 30∘ | 8 min 34 s | 0.2 (1 coinc. per 5 s)b | 0.2 (1 coinc. per 5 s)b
Palomar | 250 m | 28∘ | 13 min 39 s | 0.034 (1 coinc. per 29.3 s)b | 0.023 (1 coinc. per 42.8 s)b
Palomar | 1,012 m | 55∘ | 1 h 15 min | 8.8 10-4 (1 coinc. per 18 min 45 s) | 0 (no coincidence)

a: related to field-of-view center. b: extrapolated from 3/4 of the detectors
(full detector otherwise).

In summary, these results highlight the importance of the long baseline in
excluding false alarms such as Cherenkov showers, as the baseline introduces a
large parallax and a timing offset between nearby Earth objects, compared with
astronomical sources. The 1-km baseline reduces the number of Cherenkov
showers detected by both telescopes due to the angle at which events are
detected. Even though a large part ($\sim$70- 80 %) of showers can be
identified from their elongated or widely-spread shapes using side-by-side
telescopes, the long 1 km baseline enabled the ability to distinguish all
events from astronomical sources at infinity during our spaced-telescope
observations. We plan to still use the side-by-side modules for confirmation
follow-up and tracking if an interesting candidate source is detected.

Even though the study of Cherenkov showers is not the PANOSETI primary
objective, the long baseline gives the capability to characterize these events
in finer details. Given the different orientations of the elongated shower
measured with our data (see for instance Fig. 5), we can use the stereoscopic
effect given by the long baseline to deduce the original direction of the
cosmic-rays in Earth’s atmosphere. For showers induced by cosmic-rays (i.e.,
electrically-charged particles such as nuclei), their trajectories in space
are bent by interstellar magnetic fields making it difficult to identify the
true direction from which they originate. However, for particles with the
highest energy ($>$10 EeV, i.e., $>$1019eV), their deflection can be small
enough that it is indeed possible to identify their sources[37]. Unlike
cosmic-rays, $\gamma$-ray photons travel in straight lines and are not
affected by interstellar magnetic fields, making it possible to deduce the
source from which they originate directly from their Cherenkov shower axis. As
future work, we plan to determine the energy of the particle or $\gamma$-ray
emission that initiated the shower, and, if possible, to identify their
sources. All Cherenkov events will be time stamped and recorded in the
PANOSETI software system.

##  6 Dome assembly

Assemblies of high-time-resolution telescopes capable of observing different
parts of the sky instantaneously are still needed to survey the entire sky
quickly and repeatedly. We presented a conceptual design for a PANOSETI
observatory based upon an assembly of refracting $\sim 0.5$ m Fresnel
telescopes tessellating two geodesic domes[19, 20, 21]. This design produces a
spherical layout of collecting apertures that optimizes the instrument
footprint, aperture diameter, instrument sensitivity, and total field-of-view
coverage. While our initial design of the entire assembly had a larger total
collecting area, we further optimized it by reducing the tessellation
frequency of the geodesic dome structure, and opted for a telescope design
that could fit into all triangles of a single-layered geodesic dome. This
lower-frequency design significantly minimizes the amount of material required
for building the frame structure and telescope support compared to our
previous design[19]. Reducing the tessellation frequency also reduces the
number of required customized parts, which gives the capability to use off-
the-shelf struts and joints for the geodesic structure that are relatively
common and inexpensive. Other considerations, such as the size of available
astronomical enclosures capable of protecting the telescope equipment from
inclement weather have imposed constraints on the assembly size, telescope
apertures, and number of telescopes per assembly.

The geodesic assembly design proposed for the PANOSETI spatial configuration
of telescopes is generated from the triangular subdivision of a truncated
spherical icosahedron. Among the five Platonic solids, it has the largest
number of identical faces, making it a natural choice for designing an all-sky
instrument structure supporting a collection of identical telescopes. Any grid
on one face of the spherical icosahedron can be replicated, thus covering the
entire sphere with a pattern of regular shapes covering the $\pi$-sr spherical
cap that the PANOSETI instrument will be surveying at elevation greater than
30∘. The tessellation frequency defines the number of segments that the grid
divides into the polyhedron edges. Various subdivision schemes have been
developed to form uniform grids of nearly-identical flat triangles[38],
differing in the choice of great-circle arcs to divide the face. Each
subdivision method gives a grid with different chord lengths, triangle areas,
and shapes. We used the Equal-arcs (three great circles) subdivision
method[38] to design a uniform grid which has, among the other tested
subdivision methods[19], the largest diameter of the inscribing circle that
fits into all triangles, as well as the smallest variance in strut length and
triangle area, and the most uniform distribution of face orientations[38].
Telescopes will be attached to the geodesic dome struts using brackets that
are adjustable to a few cm irregularities to define the telescope placement
with respect to the incenter of each triangle. Attachment points will be added
to the telescope tube to interface with the geodesic dome at an adjustable
orientation[27].

Table 2: PANOSETI geodesic dome main parameters and effective sky coverage
with redundancy for different configurations of telescope assemblies.

Assembly Configurations | A | B | C
---|---|---|---
Dome(s) per observatory | 1 | 1 | 2
Nb modules per observatory | 80 | 45 | 90
Tessellation frequency | 4ν | 3ν | 3ν
Geodesic dome sizes:
Dome Diameter (incl. telescopes) | 6.2m (20.5ft) | 4.84m (15.8ft) | Same than 3ν
Dome Height (incl. telescopes) | 2.2m (7.3ft) | 1.8m (5.6ft) | Same than 3ν
Possible Enclosures | AstroHaven 22ft | AstroHaven 18ft | Same than 3ν (x2)
| LosBerger 10m | LosBerger 6m |
|  | Baader 6.5m |
Geodesic dome components:
No. of struts | 130 | 75 | 2x 75
No. of unique strut length | 5 | 3 | Same than 3ν
Struts lengths [min - max] | 0.92m - 1.07m | 0.93m - 1.07m | Same than 3ν
No. of unique face panels | 5 | 3 | Same than 3ν
Diameter of largest inscribing | 0.52m (20.4in) | 0.53m (20.5in) | Same than 3ν
circle which fits in all panelsa |  |  |
No. of joints | 51 | 31 | 2x 31
No. of pixels per dome | 81,920 | 46,080 | 92,160
PANOSETI | 7,449 sq.deg. | 4,441 sq.deg. | 7,502 sq.deg.
instantaneous | 2.27 sr | 1.35 sr | 2.28 sr
field-of-view | 18.0% of the sky | 10.8% of the sky | 18.2% of the sky
Redundant coverage | 422 sq.deg. | 0 sq.deg. | 1,442 sq.deg.
Redundant pixels | 6.7% | 0% | 19.2%

a: using 1.3-in strut diameter

Figure 6: PANOSETI all-sky telescope assemblies and projected field-of-views
in different configurations: 80 Fresnel telescopes assembly using a 4th-
frequency geodesic dome layout (top-left panel), and a configuration of 45
telescopes using a 3rd-frequency geodesic layout (top-middle panel). Top-right
panel illustrates the clearance space of the 45-telescopes configuration
inside a 18-ft diameter dome enclosure. Bottom panels represents the
respective instantaneous field-of-view coverage (blue squares) with an
additional configuration (bottom-right panel) of two 3rd-frequency domes
tilted by 14.4∘ and rotated with respect to each other (blue and green squares
in field-of-views).

Table 2 summarizes and compares PANOSETI main assembly parameters in three
possible configurations: a 4th-frequency assembly of 80 telescopes resulting
in 7,450 sq.deg. of instantaneous field-of-view coverage with 442 sq.deg. of
redundant coverage, a 3rd-frequency assembly of 45 telescopes (4,441 sq.deg.
instantaneous field-of-view coverage without redundancy), and a double 3rd
frequency assembly of 90 telescopes separated in two separate domes (7,502
sq.deg. instantaneous field-of-view coverage with 19.2$\%$ redundancy). These
assemblies are illustrated in Fig. 6 along with a representation of their
instantaneous field-of-view coverage. The configuration B using 45 telescopes
for a total of 4,441 sq.deg. of instantaneous field-of-view has been chosen
for the PANOSETI dual observatory, given that its smaller number of telescopes
allows faster deployment and ease in construction with readily available
enclosures.

In summary, this PANOSETI experiment has shown proof-of-concept for detecting
astronomical transients over a wide field-of-view with high sensitivity. The
deployment of pairs of telescopes has confirmed the benefit of using a long
baseline ($>$1 km) to identify and characterize false alarms generated by
nearby transient phenomena occurring in Earth’s atmosphere.

###### Acknowledgements.

We would like to acknowledge the Lick Observatory staff and engineers for the
installation of the two PANOSETI telescopes into the Astrograph dome, with a
special thanks to Adam Nichols. We would also like to thank the Palomar
Observatory staff in their support of our on-site baseline observations. The
PANOSETI research and instrumentation program was made possible by the support
and interest of Franklin Antonio. We thank the Bloomfield Family Foundation
for supporting SETI research at UC San Diego in the CASS Optical and Infrared
Laboratory. Harvard SETI was supported by The Planetary Society and The
Bosack/Kruger Charitable Foundation. UC Berkeley’s SETI efforts involved with
PANOSETI are supported by NSF grant 1407804, the Breakthrough Prize
Foundation, and the Marilyn and Watson Alberts SETI Chair fund.

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