design and development of an autonomous radar receiver for

Design And Development Of AnAutonomous Radar Receiver For The

Detection Of Ultra High Energy CosmicRays

By

Samridha Kunwar

Submitted to the graduate degree program in Physics and Astronomy and the GraduateFaculty of the University of Kansas in partial fulfillment of the requirements for the degree of

Doctor of Philosophy

(Chairperson) David Z. Besson

Christopher Allen

Philip S. Baringer

Douglas McKay

Daniel Tapia Takaki

Date Defended: 20th of February 2015

The Thesis Committee for Samridha Kunwarcertifies that this is the approved version of the following thesis:

Design And Development Of An Autonomous Radar Receiver For The Detection Of Ultra HighEnergy Cosmic Rays

Chairperson David Z. Besson

Date Approved: 23rd of March 2015

ii

Abstract

The detection of ultra-high energy cosmic rays is constrained by their flux, requiring detectors with

apertures of hundreds or even thousands of square kilometers and close to one hundred percent duty

cycle. The sheer scale that would be required of conventional detectors, to acquire sufficient statistics

for energy, composition or anisotropy studies, means that new techniques that reduce manpower and

financial resources are continually being sought. In this dissertation, the development of a remote

sensing technique based observatory known as bistatic radar, which aims to achieve extensive coverage

of the Earth’s surface, cf. Telescope Array’s 700 km2 surface detector, is discussed.

Construction of the radar projects transmitter station was completed in the summer of 2013, and

remote receiver stations were deployed in June and November of 2014. These stations accomplish

radar echo detection using an analog signal chain. Subject to less radio interference, the remote stations

add stereoscopic measurement capabilities that theoretically allow unique determination of cosmic ray

geometry and core location. An FPGA is used as a distributed data processing node within the project.

The FPGA provides triggering logic for data sampled at 200 MSa/s, detecting Cosmic Ray shower

echoes chirping at −1 to − 10 MHz/µs (depending on the geometry) for several µs. The data acqui-

sition system with low power consumption at a cost that is also comparatively inexpensive is described

herein.

iii

I would like to dedicate this dissertation to my loving parents, Dr. Uttam and Sharada Kunwar.

iv

Acknowledgements

And I would like to express my sincerest gratitude to my advisor Dave Besson. Dave, would have you

believe that acknowledgements are unnecessary and gratuitous, and while perhaps true, I don’t want to

leave Kansas without thanking everyone for making graduate school such a wonderful experience. So

thanks Dave, it has genuinely been a privilege working for you. Thinking about my time here, I don’t

think I would have rather been anywhere else despite the initial rough road towards my degree.

I’d also like to thank Ken Ratzlaff and Rob Young, without whom the project would have never gotten

off the ground.

Also thanks to Jess and Mark Stockham, Steven Prochyra and Jordan Hanson for numerous insightful

discussions but, more importantly for making working in Malott enjoyable, which I will remember with

fondness years from now.

My brother Saket Kunwar also needs to be thanked, not just because he’d be angry if I didn’t but, I am

forever indebted to him for his inspiration and getting me started with my career. ’Dhanyabad Dai’.

Finally, I’d like to thank you the reader for taking the first steps towards enlightenment.

v

Contents

Contents vi

List of Figures viii

List of Tables xiv

Glossary xvii

1 Cosmic Rays 11.1 Energy Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2

1.2 Anisotropies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6

1.3 Composition . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8

1.4 Candidate Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9

2 Radar Detection of Cosmic Rays 132.1 Extensive Air Showers (EAS) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13

2.1.1 Longitudinal Shower Development . . . . . . . . . . . . . . . . . . . . . . . 14

2.1.2 Lateral Shower Development . . . . . . . . . . . . . . . . . . . . . . . . . . . 14

2.2 EAS Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

2.2.1 Plasma Frequency . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

2.2.2 Plasma Lifetime . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

2.3 Radar Detection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22

2.3.1 Radar Detection of Ionization Trails . . . . . . . . . . . . . . . . . . . . . . . 22

2.3.2 Radar Detection of Relativistic Ionization Discs . . . . . . . . . . . . . . . . . 26

2.4 Bistatic Radar UHECR Experiments . . . . . . . . . . . . . . . . . . . . . . . . . . . 28

2.4.1 Mixed Apparatus for Radar Investigation of Atmospheric Cosmic-rays of High

Ionization (MARIACHI) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28

2.4.2 TARA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28

vi

Contents

3 Log Periodic Dipole Antenna 333.1 Theory of Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

3.2 Antenna Simulations with NEC . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

3.3 Antenna Characterization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 37

3.4 Effective Height . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40

3.5 Galactic Floor . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 42

4 Cosmic Ray Detector 464.1 Station Structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46

4.2 Mixer Module . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46

4.3 Triggering Mechanism . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

4.4 Chirp Acquisition Module (CAM) . . . . . . . . . . . . . . . . . . . . . . . . . . . . 52

4.4.1 Trigger Board . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 52

4.4.2 High-Speed Board . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56

4.4.3 FPGA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56

4.4.4 GPS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

4.4.5 Single Board Computer (SBC) . . . . . . . . . . . . . . . . . . . . . . . . . . 65

4.5 System Monitoring and Powering . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66

4.5.1 TDA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67

4.5.2 IVT Board . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67

4.5.3 SHM . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

4.6 Chirp Calibration Unit (CCU) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

5 TARA Remote Station Detection of UHECR’s 715.1 Snapshot Triggers . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71

5.2 Expected Event Rate . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72

5.3 Calibration Events . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73

5.4 Constraints on the Radar Cross Section . . . . . . . . . . . . . . . . . . . . . . . . . . 74

6 Conclusion 78

References 80

Appendix A Schematics 84

vii

List of Figures

1.1 Averages of Ionization measurements made by Hess on August 7th,1912. . . . . . . . 2

1.2 Energy Spectrum from 108 to 1021 eV. Taken from [1]. . . . . . . . . . . . . . . . . . 3

1.3 Panorama of interactions of possible cosmic primaries with the CMB. Taken from [2] . 4

1.4 Mean energy of protons as a function of propagation distance through the CMB. Curves

are for energy at the source of 1022 eV , 1021 eV , and 1020 eV. Taken from [2] . . . . . 5

1.5 Energy Spectrum from 1017 to 1021 eV. Taken from [1]. . . . . . . . . . . . . . . . . 6

1.6 Fraction of events correlating with AGNs as a function of the cumulative number of

events by the Auger Observatory. The dotted line shows the expected correlating frac-

tion for isotropic cosmic rays. Taken from [3]. . . . . . . . . . . . . . . . . . . . . . 8

1.7 Pierre Auger Observatory sky map plotted using 9 radius-circles drawn around known

2MRS objects within 90 Mpc. The dashed line is the field-of-view limit for the Auger

Observatory (for θ ≤ 80) and the blue solid line corresponds to the Super-Galactic

Plane. Taken from [4]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9

1.8 Telescope Array significance map plotted using 20 radius-circles. Both the signifi-

cance of the hotspot, as well as the number of events observed, are indicated by the

color coding. Taken from [5]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10

1.9 Hillas Plot. Sources above the top (red) line can accelerate protons up to 1021 eV and

sources above the bottom (green) can accelerate iron up to 1020 eV. Taken from [6]. . . 10

1.10 Unified scheme of AGN. Taken from [6]. . . . . . . . . . . . . . . . . . . . . . . . . 11

2.1 Longitudinal Distribution of particles in a ∼ 1018eV proton-initiated shower with an

incoming zenith angle of 60.4±0.3 in 870 Vertical Steps of 1g/cm2. . . . . . . . . . 15

2.2 Longitudinal Distribution of the energy content in different particle species for a ∼1018eV proton-initiated shower with an incoming zenith angle of 60.4± 0.3 in 870

Vertical Steps of 1g/cm2. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16

2.3 Number of shower electrons at varying depth. Blue: ∼ 1020eV, Green: ∼ 1019eV and

Magenta: ∼ 1018eV primaries. All Zenith angles are ∼ 60.4±0.3. . . . . . . . . . . 16

viii

List of Figures

2.4 Number of shower electrons at Shower Max at varying energies and zenith angle. Blue:

∼ 1020eV, Green: ∼ 1019eV and Magenta: ∼ 1018 eV primaries. . . . . . . . . . . . . 17

2.5 Lateral Distribution of shower electrons for primary proton-initiated shower of 1018,

1019 and 1020 eV incident at ∼ 60. . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

2.6 Longitudinal oscillations in a plasma. . . . . . . . . . . . . . . . . . . . . . . . . . . 18

2.7 Dispersion relation for waves in a plasma. . . . . . . . . . . . . . . . . . . . . . . . . 19

2.8 Plasma Frequency for primary proton initiated shower of 1018, 1019 and 1020 eV inci-

dent at ∼ 60. Also shown is a 54.1 MHz reference line, corresponding to our experi-

mental sounding frequency. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

2.9 Plasma frequency for a 1019eV primary proton-initiated shower at various zenith angles

and radii at shower maximum. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

2.10 The lifetime of plasma in the atmosphere at standard temperature and pressure due to

the attachment process discussed in the text. . . . . . . . . . . . . . . . . . . . . . . 22

2.11 Bi - static Radar Configuration. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24

2.12 RCS for 1019 and 1020 eV showers at varying zenith angles and a constant θ = 45. . . 25

2.13 Signal-to-Noise iso-contours for PT = 40 KW, GT = 22 dBi, GR = 12 dBi , λ = 5.545

m, σEAS = 400 cm2, B = 24 MHz, Ts = 300 K and separation between Transmitter

and Receiver 40 km. The dashed purple line is the Lemniscate of Bernoulli (figure

8-like curve) at ∼ 22.1 dB and the black dashed lines are the transmitter and receiver

horizontal beamwidths of 10 and 72 respectively. . . . . . . . . . . . . . . . . . . . 26

2.14 Bistatic Radar Chirp Rates. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

2.15 Schematic diagram representing the considered radar system and the reflection from a

plasma disk produced by a shower in the atmosphere. Taken from [7] . . . . . . . . . 27

2.16 The MARIACHI experiment. Shower scintillation detectors and Radar Cosmic Ray

Scattering (RCRS) stations are pictured. Taken from [8]. . . . . . . . . . . . . . . . . 29

2.17 The TARA experiment. Pictured are shower scintillation detectors and fluorescence

detector co-located with the radar echo station. Taken from [9] . . . . . . . . . . . . . 30

2.18 Schematic of the transmitter hardware configuration. Taken from [9] . . . . . . . . . . 30

2.19 Simulated horizontal (top) and vertical (bottom) radiation pattern of a horizontally po-

larized TARA phased YAGI Array, shown in blue. Red points represent field data taken

using a dedicated receiver antenna. Agreement between simulation and data is observed

to be excellent. Taken from [9]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 32

3.1 TARA Log Periodic Dipole Antenna (LPDA) equivalent schematic. Virtual Apex angle,

α , of ∼ 4.8. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 34

ix

List of Figures

3.2 TARA LPDA showing the active region (red), inductive elements(blue) and capacitive

elements (green) in response to an incident excitation frequency, as a consequence of

phasing. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

3.3 Current Basis Functions and their sum on a four wire segment. Taken from Ref. [10] . 37

3.4 Overlay of SWR of a horizontally polarized TARA LPDA as measured in the KU ane-

choic chamber compared with NEC simulation. . . . . . . . . . . . . . . . . . . . . . 38

3.5 Spherical Coordinate System used for defining antenna geometry. . . . . . . . . . . . 39

3.6 Radiation pattern in the XY-plane of a horizontally polarized TARA LPDA at the trans-

mitter sounding frequency of 54.1 MHz. Beamwidths (3 dB below peak gain) are shown

with red lines. Black points show field measurements. . . . . . . . . . . . . . . . . . 39

3.7 Radiation pattern in the YZ-plane of a horizontally polarized TARA LPDA at the trans-

mitter sounding frequency of 54.1 MHz. Beamwidths (3 dB below peak gain) are shown

with red lines. Black points show field measurements. . . . . . . . . . . . . . . . . . 40

3.8 Beamwidth of a horizontally polarized TARA LPDA, as measured in the KU anechoic

chamber. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41

3.9 Front-to-Back ratio of a horizontally polarized TARA LPDA, as measured in the KU

anechoic chamber and an overlay of NEC simulation. The discrepancy is likely due

to differences in the height of the antenna while taking the measurement (∼ 6 f t.) vs.

simulation (12 f t.). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41

3.10 Effective Height of the TARA LPDA. . . . . . . . . . . . . . . . . . . . . . . . . . . 42

3.11 Apparent brightness temperature of the sky. . . . . . . . . . . . . . . . . . . . . . . . 43

3.12 An experimental setup to measure average system noise. Snapshots were taken every

30 seconds for 1 week. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 43

3.13 Attenuation in dB per 100 ft of the LMR - 600 and LMR - 400 transmission lines. . . 44

3.14 Average receiver system noise floor (green) Power Spectral Density (PSD) in dBm/Hz

superimposed with a fit to the known galactic background noise (red dashed line). Sys-

tem attenuation, filters, and amplifiers were accounted for in calculating the absolute

received power. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 45

4.1 Construction schematic showing the different Components of the TARA detector. Cour-

tesy Kenneth Ratzlaff, Instrumentation Design Lab, KU. . . . . . . . . . . . . . . . . 47

4.2 A linear down-chirp in the time domain (top) and frequency-time domain (bottom). . . 48

4.3 Radar block diagram showing the down-conversion process along with the filtering and

enveloping process. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49

x

List of Figures

4.4 Top: Power spectrum of a -10 MHz/µs chirp created by a signal generator, prior to

mixing. Bottom: power spectrum of a 1 MHz monotone signal after signal mixing

and passing through a low pass filter. The chirp is evident as the left-most peak in

this distribution. The 24 MHz peak and the 48 MHz harmonic is likely due to a Serial

Peripheral Interface(SPI). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50

4.5 Block diagram showing the different components of the mixer module. Courtesy Ken-

neth Ratzlaff, Instrumentation Design Lab, KU. . . . . . . . . . . . . . . . . . . . . . 51

4.6 Top: 0 dB SNR and 1 MHz/µs chirp embedded in noise prior to "de-chirping". Second

from top: The monotone signal after input chirp is mixed with delayed copy of itself

and passed through a low-pass filter. Bottom: Monotone passed through the Agilent

8471D power detector. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53

4.7 Elements of the CAM unit showing the communications protocols. . . . . . . . . . . 54

4.8 Trigger Board Schematics. Courtesy Rob Young, Instrumentation Design Lab, KU. . . 54

4.9 Bode plot for the four bandpass filters. . . . . . . . . . . . . . . . . . . . . . . . . . 55

4.10 A 50 KHz sine wave signal response from the envelope detector in each signal path. . 55

4.11 Read operation timing for the Analog Devices AD80066 ADC. 16 bits are multiplexed

as two 8-bit words at 3 MHz per channel. . . . . . . . . . . . . . . . . . . . . . . . . 56

4.12 SPI Timing for the Analog Devices AD80066 ADC. SCLK is the 24MHz transfer rate

clock; 16 bits are transferred while valid programming i.e while SLOAD is pulled low. 56

4.13 SPI Timing for the Analog Devices AD9634 ADC. SCLK is the 24MHz transfer rate

clock; 8 bits are transferred while valid programming, i.e. while SLOAD is pulled low. 56

4.14 Spartan 6 Programming via BPI port (BPI prom 28F128P30). Taken from [11]. . . . . 57

4.15 Nexsys 3 board’s 100 MHz CMOS oscillator and two 48 MHz clocks synthesized in the

FPGA. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

4.16 Schematic illustrating synthesis of secondary triggering clocks from a 48 MHz clock. 58

4.17 Synthesis of registry configuration and data transfer clocks from the 48 MHz clock. . 58

4.18 Pmod connector with eight logic signals, two grounds and two 3.3V VCC per Pmod. . 58

4.19 Event triggering logic for the four input signals (A,B,C,D) and Threshold (Th). . . . . 60

4.20 Triggering logic for a comparator based event and a GPS PPS based snapshot trigger . 60

4.21 Right: All FPGA pins routed to the VHDC connector are located in FPGA I/O bank0.

Left: 40 data signals, 20 ground signals, and 8 power signals are found on the VHDC

connector. Taken from [11]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60

4.22 High-Speed Board’s 200 MHz differential clocks and subsequent FPGA synthesis of

200 MHz clocks. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

4.23 Input Double Data Rate (IDDR2) is implemented to set even/odd bits on the clock

rising/falling edge. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

xi

List of Figures

4.24 IDDR2 when DDR alignment is C0. . . . . . . . . . . . . . . . . . . . . . . . . . . . 62

4.25 Differential Input Buffer Primitive. . . . . . . . . . . . . . . . . . . . . . . . . . . . 62

4.26 Block Memory Signals. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

4.27 FIFO Signals. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

4.28 Simple Dual Port RAM and FIFO Implementation . . . . . . . . . . . . . . . . . . . 64

4.29 Triggering Information and ADC Data. . . . . . . . . . . . . . . . . . . . . . . . . . 64

4.30 Implementation of D Flip Flop to assert an Interrupt signal. . . . . . . . . . . . . . . 64

4.31 FPGA - SBC Interface . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65

4.32 GPS to FPGA and SBC Interface. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65

4.33 Functional Block Diagram of the Raspberry Pi Model B Rev. 2 . . . . . . . . . . . . 66

4.34 Adafruit Prototyping Pi Plate showing the corresponding GPIO pins to FPGA and GPS 66

4.35 Transient Detector Apparatus (TDA). Courtesy Kenneth Ratzlaff, Instrumentation De-

sign Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67

4.36 Current Voltage Temperature (IVT) Board. Courtesy Kenneth Ratzlaff, Instrumentation

Design Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

4.37 System Health Monitor (SHM). Courtesy Kenneth Ratzlaff, Instrumentation Design

Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69

4.38 Threshold Settings on the System Health Monitor (SHM). . . . . . . . . . . . . . . . 70

4.39 Schematic of the Chirp Calibration Unit (CCU) Courtesy Steven Prohira, Dept. of

Physics and Astronomy, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

5.1 Modulation of the RS2 antenna at 70 MHz. The green line is a least squares fit to the

measured power while the orange line is a least squares fit to the expected modulation

due to the galactic center, chosen to be Sagittarius A*. The measurements were taken

from the horizontally polarized channel. . . . . . . . . . . . . . . . . . . . . . . . . . 72

5.2 Modulation of the RS2 antenna at 70 MHz and 77 MHz. The orange line is a least

squares fit to the expected modulation due to the galactic center, chosen to be Sagittarius

A*. The measurements were taken from the horizontally polarized channel. . . . . . . 73

5.3 The measured noise floor of the TARA LPDA using the Remote Station 2 DAQ, with

the expected noise floor superimposed (primarily galactic at these energies). The mea-

surements made were taken from the horizontally polarized channel. . . . . . . . . . 74

5.4 The cumulative distribution function for the TARA remote station receivers for an en-

semble of Monte Carlo CORSIKA-simulated events at 1019 eV (green) and 1020 eV

(blue). The pink band indicates a bandwidth of 24 MHz up to 200 MHz. Cuts at 0 dB

SNR are made on the received signal relative to the receiver noise floor, assuming this

is our future operational threshold. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75

xii

List of Figures

5.5 The expected integrated TARA remote station event rate based on the procedure de-

scribed in the text. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75

5.6 Remote Station triggered chirp (data). Calibration chirp sent from CCU in the field. . 76

5.7 Self and force-triggered events during the period from 20th to the 24th of January, 2015

for RS2, prior to software-filtering contamination by the carrier at 54.1 MHz. . . . . . 76

5.8 Self and Force-triggered events during the period from the 20th to the 24th of January,

2015 for RS2 after passing through a 5th order Butterworth bandpass filter (58 to 82

MHZ) to suppress out-of-band noise. . . . . . . . . . . . . . . . . . . . . . . . . . . 77

5.9 The Radar Cross Section limit given a certain number of observed events. . . . . . . . 77

A.1 Details of the power supply board for the CAM. Courtesy Rob Young, Instrumentation

Design Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85

A.2 Details of the Band-Pass-Filters, Envelope Detectors and the ADC (AD 80066, Analog

Devices). Courtesy Rob Young, Instrumentation Design Lab, KU. . . . . . . . . . . . 86

A.3 Details of the Triggering Board and connections to the FPGA and SBC. Courtesy Rob

Young, Instrumentation Design Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . 87

A.4 Adapter for the AD9634 Board and the FPGA. Courtesy Rob Young, Instrumentation

Design Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88

A.5 Control Circuit for the chirp generator. Courtesy Rob Young, Instrumentation Design

Lab, KU. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89

xiii

List of Tables

2.1 Length and relative boom position of antenna elements of the TARA Yagi Antennas.

All elements have a outer diameter of 3/4′′. . . . . . . . . . . . . . . . . . . . . . . . 31

3.1 Length and relative boom position of antenna elements of the TARA Log Periodic

Dipole Antennas. All elements have a diameter of 1/4 ′′. . . . . . . . . . . . . . . . . 35

3.2 Number of wire segments for each wire and the corresponding length per segment. . . 36

4.1 Passband of the Butterworth Pi Filters and the corresponding Channels. . . . . . . . . 52

4.2 Truth Table for the four input channel triggering (see Fig. 4.9). . . . . . . . . . . . . . 59

xiv

Glossary

Acronyms / Abbreviations

2MRS 2 Mass Redshift Survey

ADC Analog-to-Digital Converter

AGN Active Galactic Nuclei

BRAM Block Random Access Memory

CAM Chirp Acquisition Module

CCU Chirp Calibration Unit

CMB Cosmic Microwave Background

CMOS Complementary metal–oxide–semiconductor

CORSIKA COsmic Ray SImulation for KAskade

DCM Digital Clock Manager

DDR Double Data Rate

EAS Extensive Air Shower

EFIE Electric Field Integral Equation

FFT Fast Fourier Transform

FIFO First In First Out

FPGA Field Programmable Gate Array

FSM Finite State Machine

xv

Glossary

FSRQ Flat Spectrum Radio Quasar

FWFT First Word Fall Through

GPIO General Purpose Input/Output

GPS Global Positioning System

GRB Gamma Ray Burst

GZK Greisen, Kuzmin and Zatespin

HPBW Half Power Beam Width

IBUFDS Input BUFfer Differential Signaling

IVT Current Voltage Temperature

LPDA Log Periodic Dipole Antenna

LVDS Low Voltage Differential Signaling

MARIACHI Mixed Apparatus for Radar Investigation of Atmospheric Cosmic-rays of

High Ionization

MOM Method of Moments

NEC Numerical Electromagnetics Code

NKG Nashimura-Kamata-Greisen

PAO Pierre Auger Observatory

PCM Phase Change Memory

PLL Phase Locked Loop

Pmod Peripheral Module

PPS Pulse Per Second

PV Photo-Voltaic

QGSJET Quark Gluon String with JET

RAM Random Access Memory

xvi

Glossary

RF Radio Freqquency

SBC Single Board Computer

SD Surface Detector

SHA Sample Hold Amplifier

SHM System Health Monitor

SSOP Small-Shrink-Outline-28-Package

SSR Solid-State Relay

TA Telescope Array

TARA Telescope Array RAdar

TDA Transient Detector Apparatus

UART Universal Asynchronous Receiver/Transmitter

UHECR Ultra-High Energy Cosmic Ray

VCV Véron-Cetty and Véron

VHDC Very High Density Cable

VSWR Voltage Standing Wave Ratio

xvii

Chapter 1

Cosmic Rays

The Earth is continuously being struck by particles with kinetic energies ranging from a few MeVs

to at least 3 ·1020 eVs. Over the past several decades there has been a concerted effort in understanding

their flux, their arrival energies, composition and origin in the celestial sphere. The answer to these

questions along with correlations with known astronomical maps will solve one of the biggest mysteries

of the 21st century in understanding the most violent processes in our universe.

The first discovery of cosmic ray particles is attributed to Victor Hess, a passionate balloonist,

who undertook ten ascents (five at night) between 1911 and 1913. On these balloons, measurements

were made using electroscopes. The potential drop as a function of time, measured approximately

every hour, was converted into ion-pairs created in the gas of the electroscope per cubic centimeter per

second [12]. Fig. 1.1 shows averaged measurements made from two detectors [13] on board the balloon

Böhmen on the 7th of August, 1912.

This ascent was his seventh attempt, for which coal-gas used in previous attempts was replaced

with hydrogen, enabling Hess to reach an altitude of 5350 m. At the time, conventional wisdom was

that it was the Earth’s radioactivity that was responsible for all the measured terrestrial radiation and,

as a result, should decline with altitude. Hess’s balloon experiments disproved this and demonstrated,

that following an initial decrease, there is an increase in ionization as one ascends higher into the atmo-

sphere. Additionally, following measurements on a separate ascent during a quasi-total solar eclipse on

the 7th of April, 1912; Hess claimed that not only was the ionization extraterrestrial in nature, but also

that the sun was not the only source.

Werner Kolhörster, a German Physicist, with modified temperature and pressure stable electro-

scopes verified Hess’s results in 1914 and was successfully able to measure ionization at up to 9 km

altitudes. A decade later, Millikan and Cameron in 1926, using special water resistant electroscopes

1

1.1 Energy Spectrum

0 1000 2000 3000 4000 5000Altitude(m)

0

5

10

15

20

25

30

35

40

Ionization

(ions/cm

3/s)

ascentdescent

Fig. 1.1 Averages of Ionization measurements made by Hess on August 7th,1912.

made ionization measurements at Muir Lake (alt. 11,800 ft) and Arrowhead Lake (alt. 5,100 ft). They

found that the ionization decreased steadily with depth in the water, and the difference in readings be-

tween the two lakes at different elevations were "the exact water equivalent of the absorption of the

atmosphere" [14]. These results confirmed that the ionization was indeed extraterrestrial in nature and

subsequently led to Millikan coining them as Cosmic Rays. Initially, these particles were believed to be

γ − rays. However, by 1934, an East - West asymmetry observed in their arrival directions due to de-

flections caused by the Earth’s magnetic field provided evidence that they are predominantly positively

charged particles or nuclei.

1.1 Energy Spectrum

Primary cosmic rays - those accelerated at astrophysical sources, including protons, helium, iron and

other nuclei synthesized in stars have an energy spectrum that has been studied by a variety of experi-

ments. Fig. 1.2 summarizes these results.

Apart from those associated with solar flares, the flux of such cosmic rays, J(E), below 1010eV, are

modulated by solar magnetic fields. These fields decelerate and partially deflect cosmic rays from the

inner solar system. In addition, cosmic ray arrival intensities are also subject to geomagnetic fields.

At higher energies, the flux is non-thermal and follows a broken power law for about ten orders of

magnitude,

J(E) =d2φ(E)

dEdΩ∼(

EeV

)−γ

(1.1)

2

1.1 Energy Spectrum

Energy (eV)

910 1010 1110 1210 1310 1410 1510 1610 1710 1810 1910 2010

-1 s

r GeV

sec

)2

Flux

(m

-2810

-2510

-2210

-1910

-1610

-1310

-1010

-710

-410

-110

210

410

-sec)2(1 particle/m

Knee-year)2(1 particle/m

Ankle-year)2(1 particle/km

-century)2(1 particle/km

FNAL Tevatron (2 TeV)CERN LHC (14 TeV)

LEAP - satellite

Proton - satellite

Yakustk - ground array

Haverah Park - ground array

Akeno - ground array

AGASA - ground array

Fly’s Eye - air fluorescence

HiRes1 mono - air fluorescence

HiRes2 mono - air fluorescence

HiRes Stereo - air fluorescence

Auger - hybrid

Cosmic Ray Spectra of Various Experiments

Fig. 1.2 Energy Spectrum from 108 to 1021 eV. Taken from [1].

where γ is the spectral index. There exist interesting features in this energy spectrum: at ∼ 1015.5 eV

the spectral index changes from 2.7 to ∼ 3 and is referred to as the knee, at ∼ 1018.5 eV the spectrum

becomes harder again, with γ changing to ∼ 2.6 (the ankle) and finally the flux steeply decreases above

3 ·1019eV. The flux drops from 1 particle/m2 − year at the knee to 1 particle/km2 − year at the ankle

and eventually to 1 particle/km2 − century towards the end of the observed spectrum.

The change in the spectral index, at the knee, is likely due to sources within the galaxy unable to

accelerate cosmic rays to higher energies, cosmic rays escaping our galaxy during diffusive propagation

processes or both. Equivalent to the latter, is that the galaxy confines cosmic rays with energies below

3

1.1 Energy Spectrum

the knee and they are unable to leak out (the ’leaky box’ model). [15] provides a detailed discussion on

the origin of the knee in the cosmic-ray energy spectrum.

The cut-off at high energies can be explained by the interaction of primary cosmic rays with the

remnant cosmic microwave background (CMB) radiation. If these cosmic rays are heavier nuclei, they

can subsequently photo-disintegrate [16] or pair-produce resulting in an energy loss of the primary. For

heavier nuclei like iron [17],

A + γCMB → (A−1)+ N

→ (A−2)+ 2N

→ A + e+ + e−(1.2)

results in a lower observed energy spectrum and lighter nuclei than at the source.

Fig. 1.3 shows possible interactions with the CMB.

Figure 3: Panorama of the interactions of possible cosmic primaries with the CMB.Curves marked by “p+γCMB → e+e−+p” and “Fe+γCMB → e+e−+p” are energy losslengths (the distance for which the proton or Fe nucleus loses 1/e of its energy dueto pair production). The curve marked by “p+γCMB → π+n or πp” is the meanfree path for photo-pion production of a proton on the CMB. The curve marked“Fe+γCMB → nucleus + n or 2n” is the mean free path for a photo-nuclear reactionwhere one or two nucleons are chipped off the nucleus. The curve marked “γ +γCMB

→ e+e−” is the mean free path for the interaction of a high-energy photon with theCMB. Added for reference is the mean decay length for a neutron indicated by “n →peν”.

are not easily met, which has stimulated the production of a large number of creativepapers.

In Figure 1 I plot the number of theoretical papers, mostly speculative, writtenon the subject of the highest-energy cosmic rays as a function of time, as found on theLos Alamos server as astro-ph papers. Over the last three years the average has beenone paper per week. In Figure 2 I list a random sample of the titles. The authors ofthese papers deserve a strong response from the experimental community.

2 Propagation of the highest-energy cosmic rays

The interaction of the particles with the cosmic microwave background (CMB) andmagnetic fields plays an important role in their propagation. All possible species ofcosmic rays with the exception of neutrinos interact with the CMB. A panorama ofthe various interactions is given in Figure 3.

4

Fig. 1.3 Panorama of interactions of possible cosmic primaries with the CMB. Taken from [2]

For protons, the high energy cut-off of the spectrum was predicted independently by Greisen,

Kuzmin and Zatespin in 1966 [18] [19] and is now known as the GZK limit . This limit primarily

arises due to the interaction of protons with the cosmic microwave background, producing delta reso-

nances that subsequently decay into nucleons and pions,

p + γCMB → ∆+(1232 MeV ) → n + π

+

→ ∆+(1232 MeV ) → p + π

0(1.3)

4

1.1 Energy Spectrum

Figure 4: Mean energy of protons as a function of propagation distance through theCMB. Curves are for energy at the source of 1022 eV, 1021 eV, and 1020 eV.

Figure 5: Fluctuation of the energy of a proton propagating through the CMB.

5

Fig. 1.4 Mean energy of protons as a function of propagation distance through the CMB. Curves are forenergy at the source of 1022 eV , 1021 eV , and 1020 eV. Taken from [2]

For a temperature of ∼ 2.73K, γCMB has number density 411cm−3, total energy density 0.26eV cm−3,

and average energy 6.34×10−4eV . Following the discussion in [20], the threshold energy for protons

interacting with a photon to produce π0,+ is ∼ 1020eV. Interactions at smaller energies by a factor of 3

or 4 are possible, as the microwave spectrum extends to higher energies (In addition to its higher-energy

thermal tail, the microwave background increases with redshift due to the expansion of the Universe).

Given the delta resonance cross-section of ∼ 5× 10−28 cm2 and the γCMB number density, the ultra-

high energy proton mean free path is ∼ 8 Mpc [21]. Each time a photo-pion interaction occurs, the

subsequent proton loses about 20% of its initial energy. Fig. 1.4 shows that after each interaction, as the

proton continually loses energy to the microwave background, the energy of the proton eventually falls

below the delta resonance threshold. As shown, for observation of trans-GZK protons, sources must be

within ∼ 100 Mpc (the GZK horizon).

Other interaction modes of cosmic protons with the cosmic microwave background are possible,

including pair production,

p + γCMB → p + e+ + e− (1.4)

which has a threshold of E ∼ 8×1017 eV, and along with redshift propagation loss, is the primary energy

loss mechanism for protons that fall below the delta resonance threshold. GZK attenuated protons

produce a ’pile-up’ while pair production and red shift propagation losses cause the steepening in the

spectrum.

5

1.2 Anisotropies

Energy (eV)1710 1810 1910 2010 2110

)-1

sec

-1 s

r-2

m2 (e

V24

J(E)

/10

3 E

-110

1

10

Yakustk - ground arrayHaverah Park - ground arrayAkeno - ground arrayAGASA - ground arrayFly’s Eye - air fluorescenceHiRes1 mono - air fluorescenceHiRes2 mono - air fluorescenceHiRes stereo - air fluorescenceAuger - hybrid

J) of Various Experiments3Cosmic Ray Spectra (E

Fig. 1.5 Energy Spectrum from 1017 to 1021 eV. Taken from [1].

For iron nuclei dominated or mixed composition models, galactic-extragalactic transitions are likely

responsible for the ankle to the left of the cosmic ray energy spectrum at an energy of approximately

1018.5 eV.

Fig. 1.5 shows the pileup and the ankle in the cosmic ray energy spectrum.

1.2 Anisotropies

One of the highest energy particles detected was the Oh-My-God (OMG) particle [22]. With an energy

of ∼ 3.2±0.9×1020 eV, the particle was detected in 1991 by the Fly’s Eye Detector in western Utah.

However, the source of such Ultra High Energy Cosmic Rays (UHECR) are still under discussion. It is

widely accepted that local galactic sources are unlikely to be cosmic accelerators producing > 1019 eV

particles, as the galactic magnetic field would not be able to confine these particles. Observations also

show that these cosmic rays have a relatively isotropic distribution and aren’t concentrated in the galactic

plane [5].

As the sources are likely to be extragalactic, deflections caused by galactic and intergalactic mag-

netic fields complicate obtaining directional information. The amount of deflection δ experienced by

any charged particle of charge q and energy ε traveling through a magnetic field B for a distance L is

6

1.2 Anisotropies

proportional to [23],

δ ∝qBL

ε, (1.5)

but, if such deflections are small, then it may be possible to determine the source location in the UHECR

sky. [24] provides a more detailed discussion regarding mapping deflections of extragalactic UHECR’s.

If trans-GZK protons and light nuclei like iron have powerful astrophysical sources, then due to the

inhomogeneity of these sources within the GZK horizon in our local universe, they should have minimal

deflections and are expected to exhibit an anisotropic arrival distribution. Therefore, it is of interest to

correlate the UHECR sky with a catalog of known energetic sources.

The recent discussion regarding anisotropies in the UHECR sky began with the Pierre Auger Obser-

vatory (PAO) in 2007. They made correlations with the arrival directions of > 55 EeV events with the

VCV(Véron-Cetty and Véron 2006) Catalog of Active Galactic Nuclei(AGN) [25]. Fig. 1.6 shows the

"most likely value of the fraction of correlated events, plotted with black dots, as a function of the total

number of time-ordered events (the events used in exploratory scans excluded)" [3]. Here 33± 5% of

events are correlating cosmic rays, compared with 21% expected for an isotropic hypothesis. The most

updated fraction of these correlating cosmic rays is 28.1+3.8%−3.6% for > 53 EeV . The energy threshold was

revised after accounting for "measurements of the fluorescence yield, a better estimate of the invisible

energy, a deeper understanding of the detector, and an improved event reconstruction" [4] resulting in a

modified correlation fraction.

Fig. 1.7 shows the Auger sky map with ten years of data from January 1st, 2004 to March 31st, 2014

for ⩾ 52EeV arrival directions. Here circles are drawn around all of the 2 Mass Redshift Survey (2MRS)

objects within 90 Mpc. Each circle has radius 9, which is the value for which the cross-correlation has

maximum significance [4]. They reported no statistically significant evidence of correlation in UHECR

distribution.

In 2014, the Telescope Array (TA) analyzed five years of Surface Detector (SD) data for cosmic

rays with energies > 57 EeV. Summing events over 20 radii circles [5] ("an angular scale comparable

to the clustering length of AGN’s within 85 Mpc" [26]), TA found deviations from isotropy, as shown

in Fig. 1.8. There were, however, no clearly identifiable sources reported. The TA suggested that, if the

hot spot is indeed real, then it could be due to the "closest galaxy group, galaxy filaments connecting

us with the Virgo Cluster" or both [5]. Alternatively, "if cosmic rays are heavy nuclei, then they may

have originated in the supergalactic plane containing the Ursa Major cluster and then been deflected by

galactic and intergalactic fields" [5].

While both experiments suggest anisotropy in the distribution of UHECRs, more statistics are nec-

essary for confirmation.

7

1.3 Composition

with an antisymmetric halo magnetic field (A) is already excluded by the upper limits. In Fig. 2the phase measurement as a function of the energy shows an interesting pattern: it suggests asmooth transition between a phase of ∼ 270 (consistent with the right ascension of the galacticcenter) below 1×1018 eV and another phase of ∼ 90 (consistent with the right ascension of thegalactic anti-center) above 5×1018 eV. This is interesting since a real anisotropy would need lessevents to be established with high statistical confidence from the phase consistency in orderedenergy intervals than by amplitude measurements [9]. New data will show if this feature stillstands.

3. Correlation with celestial objects

Figure 3. Fraction of events correlatingwith AGNs as a function of the cumulativenumber of events, starting after theexploratory data. The expected correlatingfraction for isotropic cosmic rays is shownby the dotted line.

Energy Threshold (EeV)50 55 60 65 70 75

Log1

0(p_

valu

e)

-3

-2.5

-2

-1.5

-1

-0.5

0

Data2pt-L2pt+3pt

Figure 4. Pvalue of the Auger datafor the 2pt-L, 2pt+ and 3pt methods.The minimum in Pvalue is at 100 eventsfor the 2pt+ and 3pt methods and cor-responds to an energy of about 51 ≈ EeV .

For the highest energy cosmic rays, possible candidates able to accelerate them up to1020 eV are jets of AGN, as mentioned before. The Pierre Auger Collaboration reported[11, 12] a correlation of events with the AGNs in the VCV catalogue. The first 14 events wereused for an exploratory scan that yielded the following search parameters: energy threshold(Eth = 55 EeV), maximum angular separation (Ψ ≤ 3.1) and maximum redshift (z ≤ 0.018).Those parameters minimize the probability that the correlation with AGN could result from abackground fluctuation if the flux were isotropic. The subsequent 13 events established a 99%confidence level for rejecting the hypothesis of isotropic cosmic ray flux. The reported fractionof correlation events was (69+11

−13)%. By adding data with Eth = 55 EeV up to the end of 2009

(69 events in total), the amount of correlation decreased to (38+7−6)% [13]. For this dataset, we

show in Fig. 3 the most likely value of the fraction of the correlated events plotted with blackdots as a function of the total number of time-ordered events (the events used in the exploratoryscan are excluded). The most updated estimate of the fraction of correlating cosmic rays is(33± 5)%, with 21% expected under the isotropic hypothesis [14].

A posteriori studies [13] showed that the distribution of arrival directions of the 69 cosmicrays is compatible with models (for suitable values of two parameters, the smoothing factor σand an isotropic fraction fiso) based on populations of nearby extragalactic objects, such asgalaxies in the 2MRS and AGNs in the SWIFT-BAT catalogues. The models fit the data forsmoothing angles around a few degrees and for correlating fractions of order 40% (fiso ≈ 0.6).The data does not fit either the isotropic expectation or the predictions of the models with

Fig. 1.6 Fraction of events correlating with AGNs as a function of the cumulative number of events bythe Auger Observatory. The dotted line shows the expected correlating fraction for isotropic cosmicrays. Taken from [3].

1.3 Composition

In addition to the energy spectrum and anisotropy, the composition of cosmic rays is another important

factor in identifying the origins of UHECR. The composition of cosmic rays with energy ≤ 1016eV can

be measured directly with space-based experiments [27] equipped with spectrometers and calorimeters.

However, for higher energy cosmic rays, the small flux means that indirect detection techniques are

necessary.

Indirect measurement of the composition of cosmic rays is achieved by measuring the depth of

shower maximum (Xmax) as the primary cosmic ray interacts with the atmosphere to produce secondary

particles [17]. Sec. 2.1 provides a detailed discussion of proton-initiated Extensive Air Showers. While

the penetration of the primary in the atmosphere increases with energy for all nuclei, heavy vs. light

nuclei also vary in penetration, and it is this characteristic that enables composition measurements. In

addition, fluctuations in Xmax and the number of secondary muons may be used in the determination

of the primary’s composition. These indirect measurements, rather than on an event-by-event basis, are

done statistically by a comparison of measured and simulated distributions [28].

Shower observations suggest light primaries dominate at the knee and change to heavier primaries

up to ∼ 1017 eV [27]. A reversal follows with lighter nuclei up to ∼ 1018eV. The PAO finds that

8

1.4 Candidate Sources

– 18 –

1e-05

0.0001

0.001

0.01

0.1

1

0 50 100 150 200

f min

, P

D [Mpc]

Cross-correlation with 2MRS

Pfmin

Cross-correlation 2MRS, D=90 Mpc

0 5 10 15 20 25 30

Ψ [deg]

40

45

50

55

60

65

70

75

80

Eth

[EeV

]

0.001

0.01

0.1

1

f

+

Fig. 5.— Cross-correlation of events with the galaxies in the 2MRS catalog. The top-left

panel shows the values of fmin and P as a function of the maximum distance D to the

galaxies considered. The top-right panel shows the results of the scan in ψ and Eth for the

value D = 90 Mpc corresponding to the minimum values in the top-left plot. The bottom

plot shows the sky distribution (in Galactic coordinates) of the events with E ≥ 52 EeV

(black dots). Blue fuzzy circles of 9 radius are drawn around all of the 2MRS objects

closer than 90 Mpc. The dashed line is the field-of-view limit for the Auger Observatory

(for θ ≤ 80) and the blue solid line corresponds to the Super-Galactic Plane.

penalization due to the scan performed in D, the probability of obtaining a value of P

smaller than 1% from isotropic distributions for any value of D is P ≃ 6%. Finally, we

show the map of events and objects in the bottom panel. Given the minimum found, we

include events with E ≥ 58 EeV and draw circles of 1 radius around the BAT AGNs

closer than 80 Mpc.

The results of the cross-correlation with jetted radio galaxies are shown in Figure 7.

The minimum value fmin ≃ 2×10−4, with P ≃ 1.4%, is obtained for D = 10 Mpc (see

top-left panel). The only object included in this catalog within such a distance is the

Centaurus A galaxy. Since the correlation with Cen A is discussed separately in the next

section, we consider here the second minimum, which is found for D = 90 Mpc. This

Fig. 1.7 Pierre Auger Observatory sky map plotted using 9 radius-circles drawn around known 2MRSobjects within 90 Mpc. The dashed line is the field-of-view limit for the Auger Observatory (for θ ≤80) and the blue solid line corresponds to the Super-Galactic Plane. Taken from [4].

the spectrum is dominated by light nuclei from 1018eV to 1018.5eV [29] after which a heavier nuclear

component becomes significant [30]. By contrast, the TA can explain their data over this entire energy

range without a heavy nuclear admixture [31], leading to a discrepancy in these energy regimes between

the two observatories.


While the source of such UHECRs has not been pinpointed, there are various candidate sources that are

best described by the Hillas plot in Fig. 1.9, that may provide a window into understanding the most

violent processes in the universe. In this Figure, the sources above the top (red) line indicate those

sources that can accelerate protons to energies up to 1021 eV, while those above the bottom (green) line

can accelerate heavier iron nuclei up to 1020 eV.

The Hillas criterion upon which Fig. 1.9 is based places a limit on the maximum energy of a particle

of electric charge q, leaving the source with magnetic field B and radius Rs:

εmax = qBRs (1.6)

This constraint is obtained by demanding that the Larmor radius [6], RL = ε/qB is contained within

9


– 12 –

Fig. 1.— Aitoff projection of the UHECR maps in equatorial coordinates. The solid curves

indicate the galactic plane (GP) and supergalactic plane (SGP). Our FoV is defined as the

region above the dashed curve at Dec. = −10. (a) The points show the directions of the

UHECRs E > 57 EeV observed by the TA SD array, and the closed and open stars indicate

the Galactic center (GC) and the anti-Galactic center (Anti-GC), respectively; (b) color

contours show the number of observed cosmic ray events summed over a 20-radius circle;

(c) number of background events from the geometrical exposure summed over a 20-radius

circle (the same color scale as (b) is used for comparison); (d) significance map calculated

from (b) and (c) using Equation 1.

Fig. 1.8 Telescope Array significance map plotted using 20 radius-circles. Both the significance of thehotspot, as well as the number of events observed, are indicated by the color coding. Taken from [5].

4 Sources and acceleration of high energy cosmic rays

pulsars

Galactic halo

AGN cores

GRB

cluster

SNR

radio galaxies

AU Mpckpcpc

−9

−7

−5

−3

−1

1

3

5

7

9

11

13

0 2 4 6 8 10 12 14 16 18 20 22log(R/km)

log(B/G

)

Figure 4.1: Magnetic field strength versus size of various suggested cosmic ray sources.

Blandford argument The acceleration of a proton to the energy E = 1020 eV by regularelectromagnetic fields requires the potential difference U = 1020 V. What is the minimal powerP dissipated by such an accelerator? In order to use the basic equation P = UI = U2/Rknown from high-school physics, we have to know the appropriate value of the resistance R.Since the acceleration region is in most cases nearly empty, we use R ∼ 1000 Ω (lead by the“impedance of the vacuum”, R = 4πk0/c = 1/(ǫ0c) ≈ 377Ω). Hence a source able to produceprotons with E = 1020 eV by regular acceleration in electromagnetic fields has the minimalluminosity [12]

L = U2/R >∼ 1037 W = 1044 erg/s . (4.1)

This can be transformed into an upper limit on the density ns of ultrahigh energy cosmicrays (UHECR) sources, since the observed UHECR intensity fixes the required emissivity L,i.e. the energy input per volume and time, as L ∼ 3× 1046erg/(Mpc3yr). Hence, the densityof UHECR sources able to accelerate protons to E = 1020 eV should be smaller than ns =L/L ∼ 10−5/Mpc3, if the acceleration is by regular electromagnetic fields. For comparison, thedensity of normal galaxies is ns ≈ 10−2/Mpc3, while the most common type of active galacticnuclei in the nearby Universe, Seyfert galaxies, has the density ns ≈ (1 − 5) × 10−5/Mpc3

within redshift z <∼ 0.02.

4.1.2 Specific sources

Most galactic astrophysical sources are connected with type II (or core-collapse) supernovae(SN) and their remnants (SNR): Examples are the direct acceleration in the magnetosphere of

28

Fig. 1.9 Hillas Plot. Sources above the top (red) line can accelerate protons up to 1021 eV and sourcesabove the bottom (green) can accelerate iron up to 1020 eV. Taken from [6].

10

1.4 Candidate Sources4.1 Sources of high energy cosmic rays

Figure 4.2: The unified scheme of AGN.

an extremely efficient energy generation mechanism. For accretion on a BH, the maximalenergy gain is Emax ∼ GmM/RS , where the Schwarzschild radius is RS = 2GM/c2, and thusEmax = mc2/2. A large part of this energy will be lost in the BH, while the remainder heatsup via friction an accretion disc around the black hole. Modeling the accretion process givesan efficiency of ǫ = 10%–20%. Thus the luminosity from accretion is

L =ǫc2dm

2dt. (4.9)

For a rather modest mass consumption of the BH, dm/dt = 1M⊙/yr, one obtains 6×1045erg/sor L ∼ 1012L⊙.

The unified picture of AGNs is illustrated in Fig. 4.2. The different AGN types are onlyfacets of the same phenomenon–accretion on a SMBH—viewed from different angles, at differ-ent stages of activity (small or large dm/dt) and evolution in time (e.g. from a quasar phaseat redshift z ∼ 1–4 towards a Seyfert galaxy at present). Blazars are AGN with a relativisticjet that is pointing in the direction of the Earth and are therefore often ranked among themost promising sources of ultrahigh energy cosmic rays. According Fig. 4.2, blazars are eitherFlat Spectrum Radio Quasars (FSRQ) or BL Lac objects.

Gamma Ray Bursts come in two (main) sub-varieties, depending on their duration thatvaries from fraction of a second to many minutes. While short duration GRBs are most likelythe result of binary collisions between e.g. neutron stars, long duration GRBs which makeup about 2/3 of all GRBs are associated with supernova events in extremely massive stars.GRBs are highly beamed sources of gamma-rays and perhaps also of high energy neutrinosand cosmic rays. A distinctive feature is the high Lorentz factor of shocks in GRBs.

31

Fig. 1.10 Unified scheme of AGN. Taken from [6].

the accelerator of size Rs [6],

RL = ε/qB ≤ Rs (1.7)

Conceptually, to produce the highest energy particles, small accelerators will need to have large

magnetic fields to trap charged particles in the acceleration region or, vice-versa, a cosmic accelerator

with a small magnetic field requires an immense acceleration region of the order of Mpc.

Possible candidates for UHECR origin include (a) Gravitational Accretion Shocks where dark mat-

ter accretion can cause shocks around large-scale structures in the Universe, (b) Gamma Ray Bursts

(GRB) where multiple shock regions can form potential zones of acceleration, (c) Neutron Stars that

have extremely strong dipole fields of order 10−15G [27] and finally, the most discussed source: (d)

Active Galactic Nuclei (AGN).

Active Galaxies have a total luminosity exceeding the thermal emission of the individual stars com-

prising the galaxy due to accretion of supermassive black holes located at their center. The type of AGN

observed is dependent on the observing angle, stage of activity and evolution in time [6], as shown in

Fig. 1.10. Amongst these AGN’s, Blazars have relativistic jets pointing towards Earth and are consid-

ered strong candidates as UHECR sources [32]. These Blazars are either Fanaroff-Riley I galaxies and

their associated BL Lac objects or Fanaroff-Riley II galaxies and their associated Flat Spectrum Radio

Quasars (FSRQ) .

Understanding the energy, composition and anisotropy of cosmic rays will provide answers to some

11


of the most intriguing questions in our Universe. The next chapter discusses physics of cosmic ray

interactions and their detection using a novel radar technique.

12

Chapter 2

Radar Detection of Cosmic Rays

In this chapter, we discuss the properties of Extensive Air Showers (EAS) and the technique of

detection by radio interrogation.

2.1 Extensive Air Showers (EAS)

Cosmic rays with energies per nucleon in excess of 1014 eV [33] create cascades of particles with

electromagnetic and hadronic components in the atmosphere, known as Extensive Air Showers (EAS).

The development of the shower can be defined by the first interaction in the atmosphere. The interaction

of a primary proton with an absorber nucleus A can be represented as,

p + A → p + X + π0,± + K0,±, (2.1)

where X is the fragmented nucleus. Several particles such as K, Λ, η , ...... are produced. However,

the pions carry away roughly half of the energy (dependent on interaction elasticity) and the remainder

by the remnant cosmic ray, which continues to interact. The entire shower development can then be

described in terms of the behavior of the pions produced in the first few interactions.

Neutral pions (π0) with a rest lifetime of ∼ 10−16s have the decay modes,

π0 → γγ, π

0 → γe+e−, and π0 → e+e−e+e−, (2.2)

where the two-photon mode dominates. These photons go on to produce e+e− pairs that then emit

bremsstrahlung photons after traveling some distance. When

−dEdx

|bremsstrahlung = −dEdx

|ionization, (2.3)

13

2.1 Extensive Air Showers (EAS)

the electron is said to be at its critical energy (∼ 84 MeV in air) and the shower evolution at that point

described as the shower maximum.

Primary cosmic ray energy affects the decay time and interaction length of the secondary charged

pions (π±). These pions can interact to initiate subsequent pion showers or can decay to produce muons

through their most common decay mode,

π+ → µ

+ + νµ and π− → µ

− + νµ (2.4)

2.1.1 Longitudinal Shower Development

The Gaiser-Hillas functional form [34] has proven to be useful in fitting the measured and simulated air

shower development with various hadronic models like QGSJET (Quark Gluon String with JET) and

QGSJETII [35]. The form is given by,

N(X) = Nmax(X −X0

Xmax −X0)

Xmax−X0λ

eX0−X

λ , (2.5)

where Nmax is the maximum number of charged particles, Xmax is the atmospheric depth where the

maximum occurs, X0 is the atmospheric depth of the first interaction and λ is a fit parameter.

The cascade of these secondary particles produced in the atmosphere have been investigated using

Monte Carlo Simulations such as CORSIKA 6.50 (COsmic Ray SImulation for KAskade) [36] with the

QGSJET II hadronic interaction subroutine.

Figs. 2.1 and Fig. 2.2 show CORSIKA simulation of the longitudinal distribution of the number

of charged particles and energy deposition, respectively. Shown is a ∼ 1018eV proton-initiated shower

at a zenith angle of 60.4± 0.3, using a realistic curved Earth surface model. Shower maximum is at

∼ 400g/cm2.

Fig. 2.3 shows a comparison of the number of shower electrons for showers initiated by 1018eV,

1019eV, and 1020eV cosmic ray protons at different shower depths.

Several hundred simulations were run for proton energies 1018 ≤ E ≤ 1020.5, showing the relation-

ship between depth at shower maximum and zenith angle of the arriving primary proton. The number

of electrons produced at shower max does not vary significantly with the incoming angle of the primary

proton (see Fig. 2.4).

2.1.2 Lateral Shower Development

As the secondary particles have transverse momentum with respect to the shower axis, the lateral shower

profile for the electromagnetic particles is approximated by the Nishimura-Kamata-Greisen (NKG)

14

2.2 EAS Plasma

0

100

200

300

400

500

600

700

800

900

Depth (g/cm2 )

104

105

106

107

108

109

1010

Numberof

Particles

Longitudinal Distribution

γ

e+e−µ+µ−hadrons

Fig. 2.1 Longitudinal Distribution of particles in a ∼ 1018eV proton-initiated shower with an incomingzenith angle of 60.4±0.3 in 870 Vertical Steps of 1g/cm2.

function,

ρ(r) =N(s)

rm2sm

2(r/rmsm)s−2(1 + r/rmsm)s−4.5

2πβ (s,4.5−2s)(2.6)

where ρ(r) is the charged particle density at a distance r to the shower axis in particles per unit area,

and N(s) is the total number of charged particles at shower age s (s = 31+2Xmax/X such that s = 1 at Xmax).

The Moliere radius is given by

rm = 70(ρ0

ρ) m, (2.7)

where ρ is the density of the air at the altitude under consideration, and ρ0 is the density at sea level.

Fig. 2.5 shows the lateral distribution of electrons at various primary energies incident at ∼ 60. It

is essential to note that these are shower electrons as opposed to ionization electrons that contribute to

the plasma frequency discussed in the next section.

2.2 EAS Plasma

Electrons and ions close to the shower core form a plasma that on average can be assumed to be electri-

cally neutral. Such a plasma has excitation modes that are longitudinal or transverse. We now discuss

15

2.2 EAS Plasma

100

200

300

400

500

600

700

800

900

Depth (g/cm2 )

103

104

105

106

107

108

Energy

Deposition

(GeV

)


sum

neutrino

γ

emioniz

emcut

µioniz

µcut

hadrionizhadrcut

Fig. 2.2 Longitudinal Distribution of the energy content in different particle species for a ∼ 1018eVproton-initiated shower with an incoming zenith angle of 60.4±0.3 in 870 Vertical Steps of 1g/cm2.

0

100

200

300

400

500

600

700

800

900

Depth (g/cm2 )

106

107

108

109

1010

1011

1012

Numberof

Shower

Electrons


1018 eV

1019 eV

1020 eV

Fig. 2.3 Number of shower electrons at varying depth. Blue: ∼ 1020eV, Green: ∼ 1019eV and Magenta:∼ 1018eV primaries. All Zenith angles are ∼ 60.4±0.3.

the formation and excitation of the plasma. In the discussion that follows, the mass of protons is as-

16

2.2 EAS Plasma

Depth (g/cm2)

300 400 500 600 700 800 ZenithAngle

( )

1020

3040

5060

log 1

0(No.

ofShower

Electrons)

8.5

9.0

9.5

10.0

10.5

No. of Shower Electrons at Shower Max

Fig. 2.4 Number of shower electrons at Shower Max at varying energies and zenith angle. Blue: ∼1020eV, Green: ∼ 1019eV and Magenta: ∼ 1018 eV primaries.

100 101 102 103 104

r(cm)

102

103

104

105

106

107

108

Numberof

shower

electronsper

cm2

Shower Max (∼60 Zenith Angle)

1018 eV

1019 eV

1020 eV

Fig. 2.5 Lateral Distribution of shower electrons for primary proton-initiated shower of 1018, 1019 and1020 eV incident at ∼ 60.

17

2.2 EAS Plasma

sumed to be infinite compared to that of electrons.

2.2.1 Plasma Frequency

Conceptually, if atomic electrons are displaced by some distance ∆x, as shown in Fig. 2.6, regions of

positive and negative charge due to the ions and electrons are, respectively, created.

+

+

+

+

-

-

-

-

Δx Δx

E

Fig. 2.6 Longitudinal oscillations in a plasma.

The longitudinal electric field in this intervening region is

E = −nee∆xε0

, (2.8)

and the force restoring the electrons to their original position is then,

F = eE = −nee2∆xε0

= med2∆xdt2 , (2.9)

which is the equation of a simple harmonic oscillator such that the electrons oscillate coherently in the

longitudinal direction with the frequency,

ω2

p =nee2

ε0me(2.10)

Additionally, as the transverse fields are translationally invariant, the transverse canonical momen-

tum of a particle in the field is conserved such that we have the dispersion relation [37],

ω2 − k2c2 = ω

2p (2.11)

where k is the propagation constant.

18

2.2 EAS Plasma

For a wave with a frequency above the plasma frequency (ω > ωp), the propagation constant k is

real. The wave propagates through the plasma with a phase velocity, vφ = ω/k> c, as shown in Fig. 2.7.

The individual electrons scatter independently according to the Thompson cross-section,

k

ωωp

ωk

=c

Transverse

Longitudinal

Fig. 2.7 Dispersion relation for waves in a plasma.

σT =8π

3(

e2

mec2 )2 = 6.69×10−29 m2, (2.12)

When ω < ωp, the propagation constant is imaginary (i.e., k < 0), and the wave is evanescent with

no energy transmitted through the plasma. The plasma rather reflects the wave.

Fig. 2.8 shows the plasma frequency at various primary energies incident at ∼ 60 as determined

from CORSIKA simulations. Ionization electrons are calculated by measuring the ratio of the energy

deposited in the atmosphere and the mean energy per ion pair production of 33.8 eV [38].

Fig. 2.9 shows the plasma frequency for a 1019eV primary incident at different zenith angles and

various radii at shower maximum.

2.2.2 Plasma Lifetime

Various factors contribute to the lifetime of plasmas, such as diffusion through the ambient air, attach-

ment, and recombination of electrons [39].

At lower EAS shower max altitudes, the ambipolar diffusion constant Di is given by,

Di ∝ Tiνi−1, (2.13)

19

2.2 EAS Plasma

100 101 102 103 104

r(cm)

106

107

108

109

ν p(M

Hz)

Shower Max (∼60 Zenith Angle)

1018 eV

1019 eV

1020 eV

54.1 MHz

Fig. 2.8 Plasma Frequency for primary proton initiated shower of 1018, 1019 and 1020 eV incident at∼ 60.Also shown is a 54.1 MHz reference line, corresponding to our experimental sounding frequency.

where Ti, νi are the kinetic temperature and collisional frequency, respectively. The ambipolar constant

has been found to be ∼ 5 cm2s−1 [39] at EAS altitudes due to higher collisional frequency as compared

to meteor shower altitudes at ∼ 80 to 120 km. Assuming that the diffusion of EAS showers are similar

to that of meteor ionization, where for the underdense regime, the power decays exponentially with a

time constant, τm = λ 2/32π2Di, then diffusion alone would lead to a lifetime of ≃ 200 s for a frequency

of 54 MHz. Additionally, it has also been shown that for plasmas due to lightning strikes, (which have

higher ionization densities by several orders of magnitude than EAS) at ∼ 10 km and 11 cm wavelength

the time constant is 240 ms. For these reasons, diffusion is expected to have minimal effect on the

lifetime of EAS plasmas. Similarly, recombination times of electrons and ions in the troposphere are

expected to be several minutes [39] [40] and so likely even less significant than diffusion.

Attachment is the primary factor affecting EAS plasma lifetime [41]. These processes are

Two Body Attachment:

e− + O2 → O− + O, (2.14)

20

2.2 EAS Plasma

1 2 3 4 5 6 7 8 9r (cm)

10

20

30

40

50

60ZenithAngle

()

Energy∼1019 eV

20

30

40

50

60

70

80

90

100

Plasm

aFrequency

(MHz)

Fig. 2.9 Plasma frequency for a 1019eV primary proton-initiated shower at various zenith angles andradii at shower maximum.

Three Body Attachment:

e− + O2 O−∗2 ,

O−∗2 + O2 → O−

2 + O2 + energy.(2.15)

At low altitudes, the three-body attachment dominates the de-ionization process [40] due to its quadra-

ture dependence on oxygen concentration. At higher altitudes, as the concentration of oxygen dimin-

ishes, two-body attachment becomes more important in the de-ionization process. For calculating the

plasma lifetime at EAS altitudes, three-body attachment is the most significant process.

The three-body attachment rate is dependent on the particle density of the atmospheric elements

involved and given by [42],

∂Ne

∂ t= −katt1NeN2

m[O2]2 − katt2NeN2m[O2][N2], (2.16)

where Ne is the electron number density, Nm is the total number density of atmospheric molecules

(Nitrogen, Oxygen, Carbon-Dioxide and trace gases), [O2] is the fraction of Oxygen and [N2] the

fraction of Nitrogen in the atmosphere. These parameters are, [O2] = 0.209476, [N2] = 0.78084,

katt1 = 2 ·10−30cm6s−1, and katt2 = 8 ·10−32cm6s−1.

The lifetime of a plasma at sea level is found to be ∼ 15 ns and at 10 km to be ∼ 130 ns [42] as

21

2.3 Radar Detection

shown in Fig. 2.10.

2 4 6 8 10altitude (km)

0

20

40

60

80

100

120

140

lifetime

(ns)

Fig. 2.10 The lifetime of plasma in the atmosphere at standard temperature and pressure due to theattachment process discussed in the text.

2.3 Radar Detection

2.3.1 Radar Detection of Ionization Trails

The development of the ionization trail induced by cosmic rays is believed to be analogous to that of

micro-meteors that form similar ionization columns. As meteoroids penetrate the Earth’s atmosphere,

gradual ablation results in a cylindrically ionized channel along their path. Radar detection of these

meteor plasma columns have been observed for trails between altitudes of 70 and 120 km, with most

observed between 80 and 105 km [43].

Typically the smallest detectable meteors have a mass of 10−6g and dimensions of 10−3 to 10−2m.

At typical velocities of ∼ 4−5 ·104m/s, they have kinetic energy of ∼ 1J, most of which is dissipated

to ionization. However, as discussed earlier, in an EAS, the lateral distribution of particles spreads out

as the shower progresses and propagates at relativistic speeds. Subsequently, air showers result in an

ionization column with a different initial distribution than that of a meteor.

In 1941, Blackett and Lovell [44] were the first to propose cosmic rays as an explanation of anoma-

lies observed in atmospheric radar data. They suggested a point cluster approximation, with n ions

22

2.3 Radar Detection

within a Fresnel length, L f , of the ionized column,

L f =√

λR, (2.17)

where λ is the transmitter wavelength and R the distance from the transmitter to the ionized column. In

this approximation, the maximum electronic ions per cm is,

n =12

10−7 pE (2.18)

where E is the incident energy and p the pressure in atmospheric units. Consequently, in an equivalent

point cluster, the total number of electrons is given by,

N = nL f =12

10−7 pE√

λR (2.19)

They suggested that via their model, reflections observed at altitudes as low as 10 km for E ≥ 1015eV

were possibly from such cosmic ray ionization trails. However, no definitive observations were ever

reported.

In later discussions, it was suggested that the point cluster method would not lead to coherent scat-

tering due to the spread of shower electrons in the cascade process. In addition, and more importantly

attenuation is caused by the scattering of the oscillating electrons from neutral molecules that interfere

with the coherence oscillation of the plasma. This decrease in scattered energy is given by [45] and

claimed to reduce the detectable cosmic ray energy by as much as 30 dB [38],

11 +( νc

πν0)2 , (2.20)

where νc is the collision frequency and ν0 the plasma frequency.

It has now been shown that the low-energy threshold initially suggested by Blackett and Lovell is

unlikely, and an energy threshold of E ∼ 1019eV is more realistic.

While dampening can cause substantial signal attenuation, its effect may be mitigated by using a

bistatic configuration (see Fig. 2.11). In this case, the radar cross-section (RCS) is of order 4πA2/λ 2

higher than that of the monostatic radar [38]. Here λ is the electromagnetic wave wavelength, and A is

the target objects cross-sectional area.

While the radar cross section, σEAS is unknown, the bistatic radar equation yields,

PR

PT=

GT

4πR2T

σEASGR

4πR2R

λ 2

4π, (2.21)

where PT and PR are the transmitted and received power, GT and GR the transmitter and receiver antenna

23

2.3 Radar Detection

β/2δ

RT

RXTX

V

RR

θT θR

β = θT + θR

α = δ + β/2Target

Fig. 2.11 Bi - static Radar Configuration.

gains, RT and RR the distance to the ionization trail from the transmitter and receiver respectively, λ the

wavelength of the transmitted signal and σ the radar cross-section (RCS).

As the over-dense regime is likely to be the most significant contributor for radio reflection of EAS,

and since rc << λ (Rayleigh regime), where λ is the radar wavelength, the shower can be approximated

as a ’thin wire’, where the RCS is given by,

σEAS =πL2sin2θ [

sin( 2πLλ

cosθ)2πL

λcos

]2

(π

2 )2 +(ln λ

γπrcsinθ)2

cos4φ (2.22)

with L the length of the wire, rc the radius of the wire, γ = 1.78, θ the angle between the wire and the

direction of incidence, and φ the angle between the polarization direction and the plane defined by the

wire and the direction of incidence [46].

Fig. 2.12 shows the RCS for 1019 and 1020 eV showers at varying zenith angles and a constant θ =

45 assuming a vertically polarized incident signal, as calculated using CORSIKA simulations.

In addition, at the receiver, the noise power is given by PN = KBTsB, where KB is Boltzmann’s

constant, Ts the system noise temperature and B is the noise bandwidth. The received signal-to-noise

ratio at the receiver is then,

S/N =k

R2RR2

T(2.23)

where,

k =PT GT σEASGRλ 2

(4π)3KBTsB(2.24)

Assuming, PT = 40 KW, GT = 22 dBi, GR = 12 dBi , λ = 5.545 m, σEAS = 400 cm2 , B = 24 MHz, Ts

= 300 K and taking the separation between the Transmitter and Receiver to be 40 km, Fig. 2.13 shows

24

2.3 Radar Detection

0 10 20 30 40 50 60Zenith Angle ( )

101

102

103

104

σ(cm

2)

θ=45

1019 eV fitted

1020 eV fitted

Fig. 2.12 RCS for 1019 and 1020 eV showers at varying zenith angles and a constant θ = 45.

the isocontour lines for the signal-to-noise ratio in the bi-static plane in this particularly optimistic case

without considering attenuation due to collisional dampening. Here the cross-section refers to a physical

cross-section of ∼ 2 cm× 200 cm, or roughly the dimensions of the over-dense regime. The effective

area is limited by the beamwidth of the transmitter and receiver antennas.

This bi-static configuration results in a Doppler shift,

f =1λ

ddt

[RT + RR] (2.25)

where RT and RR are the line of sight from the transmitter and receiver to the target respectively and λ

the wavelength of the transmitted signal. While the target moves near the speed of light, the distances

RT and RR evolve slowly over time.

This can be expressed as,

f = 2 f0cosδcosβ/2 (2.26)

where f0 is the frequency of the transmitted signal, β the bistatic angle and δ the angle between the

direction along which the target is moving and the bisector of the bistatic angle.

As the transmitting frequency is constant and since the path lengths are time-dependent the received

signal is a Doppler like signal (chirp). The rate of change in frequency (chirp rate) is subsequently given

by,

κ =∂ f∂ t

= − f0c[1

RT+

1RR

]sinα (2.27)

25

2.3 Radar Detection

30 20 10 0 10 20 30X (km)

15

10

5

0

5

10

15

Y(km

)

TX RX

σEAS = 400 cm2

Band−width = 24 MHz

29

25

23

22

Signal−to−Noise Iso−contours (dB)

Fig. 2.13 Signal-to-Noise iso-contours for PT = 40 KW, GT = 22 dBi, GR = 12 dBi , λ = 5.545 m, σEAS =400 cm2, B = 24 MHz, Ts = 300 K and separation between Transmitter and Receiver 40 km. The dashedpurple line is the Lemniscate of Bernoulli (figure 8-like curve) at ∼ 22.1 dB and the black dashed linesare the transmitter and receiver horizontal beamwidths of 10 and 72 respectively.

where α = δ +β/2, i.e., the angle between the line of sight between the receiver and the target and the

direction of the target.

Fig. 2.14 shows the dependence of the chirp rates between the transmitter and receiver as a function

of the horizontal distance from the receiver and the angle α .

2.3.2 Radar Detection of Relativistic Ionization Discs

Recent discussions have addressed the feasibility of radar detection of the relativistic shower front [7], [47]

and [48]. In this case, the shower front is treated as a disc rather than a cylinder. The side of the disc

has a smaller cross-section than the front; consequently the primary contribution to the scattered signal

is considered to be from this shower front as shown in Fig. 2.15.

Here a transmitter (T) illuminates the shower front and is scattered from an element of the plasma

in the shower front and received at the receiver (R). The geometry of the radar system is determined by

the distances from the shower core to the transmitter (dT ) and to the receiver (dR), together with their

azimuth angles (φT and φR) and the altitude of the transmitter (hT ) [7].

Given such a geometry, and since the shower front is continuously being created and destroyed, the

interference from the reflected signal from different stages of the shower results in a frequency up-shift.

26

2.3 Radar Detection

X(km)

0 5 10 15 20 25 30α(deg)

020

4060

80100

120140

160

||(M

Hz/µs)

0

1

2

3

4

5

0.6

1.2

1.8

2.4

3.0

3.6

4.2

4.8

5.4

Fig. 2.14 Bistatic Radar Chirp Rates.

Θ

jRrR

Θs

dT

X’

Z’ Y’

XR

GT

r®

sc

HrL

L

ds

dR

r®

Y’

Z

s a

R

nejt

Α

h

h = 0

hTh

T

jTRrR

Figure 1: Schematic diagram representing the considered radar system and reflectionfrom the plasma disk produced by a shower in the atmosphere. A ground-based radiotransmitter (T) emits a radio signal, which is scattered off an element of the plasma diskand subsequently observed by the receiver antenna (R). The geometry of the radar systemis determined by the distances from the shower core to the transmitter (dT ) and to thereceiver (dR) together with their azimuth angles (ϕT and ϕR) and the altitude of thetransmitter (hT ).

transmitter (T) irradiates a disk-like static plasma left behind the showerfront. The radio signal is scattered by free electrons in the ionization trailand subsequently received by the ground-based antenna (R). The geometry ofsuch a radar system is described conveniently by the cylindrical coordinatesof the transmitter and the receiver, i.e. by the distances from the showercore to the transmitter (dT ) and to the receiver (dR), and by the angles ϕT

and ϕR. The altitude of the transmitter is given by hT , whereas the receiveris at the ground level.

The system of coordinates XYZ is chosen in such a way that the planeconstructed by the X-axis and the shower axis is perpendicular to the ground.Moreover, the X and Y axes lie at the ground level. The center of the coor-dinate system is placed at the shower core, i.e. at the point of intersectionof the shower axis with the ground. The coordinate system of the disk-likestatic plasma X ′Y ′Z ′ is simply created by rotating the XYZ frame of refer-ence around Y-axis by the shower inclination angle |θs − π/2| and translating

5

Fig. 2.15 Schematic diagram representing the considered radar system and the reflection from a plasmadisk produced by a shower in the atmosphere. Taken from [7]

27

2.4 Bistatic Radar UHECR Experiments


A handful of attempts have been made to use a bi-static radar configuration to detect the forward scatter

of radio waves from UHECR induced air showers.

2.4.1 Mixed Apparatus for Radar Investigation of Atmospheric Cosmic-rays of HighIonization (MARIACHI)

Located in Long Island, United States, MARIACHI consisted of 4 radar stations designed to detect radio

signals transmitted by local TV stations and reflected off of cosmic ray showers [8]. In addition, 12 mini

conventional scintillator detectors enabled coincidence measurements between radar and scintillator

detector triggers. Fig. 2.16 illustrates the schematic of this experiment.

The scintillation detector stations were located at high schools, with each detector comprising four

scintillators, each of 0.25m2, in classroom corners along with a counter in one of these corners. A Field

Programmable Gate Array provided logic enabling the triggering of a cosmic ray event and a Global

Positioning System with 100 ns timing accuracy provided an event time-stamp. Orthogonal Inverted

VEE dipole antennas provided directional information for radar echo candidate events, and data was

subsequently read out using PCR 1000 radio receivers.

2.4.2 TARA

The Telescope Array Radar (TARA) project is similar to the MARIACHI experiment but uses a dedi-

cated transmitter at 54.1 MHz as the radar source. In this bistatic configuration, the transmitter emits

the sounding signal, and a distant radar receiver detects radar echoes (chirp). TARA operates in con-

junction with a set of conventional cosmic ray detectors in a quiet noise environment at the Telescope

Array in Millard County, Utah (www.telescopearray.org/tara/).

Fig. 2.17 shows the layout of the TARA experiment, together with the Telescope Array’s fluores-

cence and surface array detectors.

The TARA receiver operates in Long Ridge, UT (39-14.560 N 113-05.291 W). The receiver antenna

is a dual polarized Log Periodic Dipole Antenna designed to match the expected radar echo in the

passband from 50 to 80 MHz. Details of the receiver antenna can be found in Chap. 3 and specifics of

the radar detector are given in Chap. 4.

The Transmitting facility is just outside the city limits of Hinckley, UT (39 200′ 19.824” N, 112420′ 3.24”W ).

The station comprises a phased high-gain Yagi antenna array and the transmitter.

28


Fig. 2.16 The MARIACHI experiment. Shower scintillation detectors and Radar Cosmic Ray Scattering(RCRS) stations are pictured. Taken from [8].

Transmitter

In June 2009, conventional television programming transitioned from analog to digital in the United

States and subsequently vacated the analog TV channel 2 band. TARA received a 20 KW transmitter

from KUTV Salt Lake City and obtained a second one from KTVN Reno. TARA re-purposed this

transmitter equipment and currently broadcasts continuous wave signals as WF2XZZ at 54.1 MHz, just

above the amateur radio band ending at 54 MHz.

Fig. 2.18 shows a functional block diagram of the TARA transmitter.

29


Telescope Array Radar Remote Station Design & Development

Samridha Kunwar The University of Kansas

for the Telescope Array Radar (TARA) project APS Savannah Meeting, 2014

Fig. 2.17 The TARA experiment. Pictured are shower scintillation detectors and fluorescence detectorco-located with the radar echo station. Taken from [9]

Fig. 2.18 Schematic of the transmitter hardware configuration. Taken from [9]

Transmitting Antenna

The TARA transmitting antenna array consists of 8 narrow-band Yagi antennas, designed and manufac-

tured by M2 Antenna Systems, Inc. Each Yagi, constructed of aluminum, is capable of handling 10 kW

of continuous RF power and consists of 5 elements per antenna: a driven element, a reflector, and three

director elements. The elements are mounted on a 21.6’ long and 2" diameter boom. All elements are

constructed of aluminum tubing of 3/4" outer diameter. Each element, except for the driven element, is

constructed of two equal sections that are joined at the boom via 7/8" outer diameter sleeve elements.

Tab. 2.1 shows the dimensions of these elements and the spacing along the boom.

A balanced t-match is fed from a 4:1 coaxial balun that transforms the unbalanced 50 Ω input to the

30


Element Length (in) Position (in)Reflector 107.625 -44.375Driven 100.500 0.000Director 1 99.500 51.125Director 2 97.250 131.625Director 3 97.00 193.625

Table 2.1 Length and relative boom position of antenna elements of the TARA Yagi Antennas. Allelements have a outer diameter of 3/4′′.

balanced 200 Ω used to drive the antenna. A 50 Ω 7/8" coaxial waveguide connects the balun to four

port power dividers. Further details on the transmitter power delivery can be found in [9].

The eight yagis are phased such that the forward gain is ∼ 22 dBi and the front to back ratio is

∼ 18 dB. Fig. 2.19 show measured points overlaid on an NEC model [10] for the radiation pattern of

the array in the XY and YZ-plane, respectively.

31


0° 15°30°

45°

60°

75°

90°

105°

120°

135°

150°165°180°-165°

-150°

-135°

-120°

-105°

-90°

-75°

-60°

-45°

-30°-15°

−3

3

9

15

21

0° 15°30°

45°

60°

75°

90°

105°

120°

135°

150°165°180°-165°

-150°

-135°

-120°

-105°

-90°

-75°

-60°

-45°

-30°-15°

−3

3

9

15

21

Vertical Gain (dBi)Fig. 2.19 Simulated horizontal (top) and vertical (bottom) radiation pattern of a horizontally polarizedTARA phased YAGI Array, shown in blue. Red points represent field data taken using a dedicatedreceiver antenna. Agreement between simulation and data is observed to be excellent. Taken from [9].

32

Chapter 3

Log Periodic Dipole Antenna

In this chapter, we discuss the TARA Log Periodic Dipole Antenna’s (LPDA’s) used to capture EeV

cosmic ray radar reflections.

3.1 Theory of Operation

The TARA LPDA’s consist of a sequence of parallel linear dipoles close together. The lengths (Ln’s)

and spacings (Rn’s) of these dipoles logarithmically increase, as defined by the inverse of the geometric

ratio τ [49],1τ

=L2

L1=

Ln+1

Ln=

R2

R1=

Rn+1

Rn(3.1)

with the spacing constant σ ,

σ =Rn+1 −Rn

Ln+1(3.2)

These constants determine the frequencies of operation for the LPDA: for TARA these are τ ∼ 0.82

and σ ∼ 0.54.

Successive dipoles are connected alternately to opposite sides of a transmission line with the feed

point at the end with the shorter elements. Fig. 3.1 shows an equivalent schematic of the receiver LPDA.

Tab. 3.1 gives the lengths and positions of the antenna elements on the boom from the front edge to the

back. All elements are constructed of aluminum tubing of 1/4 ′′ outer diameter (d).

Straight lines joining the antenna elements meet at a point called the virtual apex, forming an angle

2α , characteristic of antennas whose pattern and impedance is practically independent of frequency in

the band of operation. The angle α , the geometric ratio τ , and the scaling constant, σ are related by,

σ =1− τ

4tanα(3.3)

33

3.1 Theory of Operation

L1

L2

Ln

2α

Rn

R2

R1d

Element diameter

Rn-1

Ln-1

Feed Point

Fig. 3.1 TARA Log Periodic Dipole Antenna (LPDA) equivalent schematic. Virtual Apex angle, α , of∼ 4.8.

As an electric field is incident on the antenna, near-resonant length elements of the exciting fre-

quency form the active region of the antenna and can be described as being some number of wavelengths

from the virtual apex. As the excitation frequency changes, the active region consequently changes but

remains the same number of wavelengths from this virtual apex (See Fig. 3.2). That is,

Ri

λi= const., i = 1....n. (3.4)

As successive dipole elements are connected 180 degrees out of phase, relative to the active region,

the closely spaced shorter elements (< λ/2) have negligible current and provides capacitive impedance.

The larger spaced and longer elements (> λ/2) have small currents and the phase reversal provides

inductive impedance.

34

3.2 Antenna Simulations with NEC

Element Length (in) Position (in)1 21.875 3.6252 26.625 18.06253 32.5 35.6254 39.625 57.05 48.3125 83.1256 58.3125 115

Table 3.1 Length and relative boom position of antenna elements of the TARA Log Periodic DipoleAntennas. All elements have a diameter of 1/4 ′′.

L/λ

R/λ

Active Region

Virtual Apex

2α

Fig. 3.2 TARA LPDA showing the active region (red), inductive elements(blue) and capacitive elements(green) in response to an incident excitation frequency, as a consequence of phasing.


Several parameters characterize an antenna: Antenna geometry, operating frequency, and the environ-

ment, particularly the type of ground. To get information on characteristics like directionality one must

solve the electric field integral equation (EFIE),

ε (r) = −ıωµ

∫V

G(r, r′) j(r′)dr′ (3.5)

where j(r′) is the current density located in the volume defined by the antenna geometry, µ = µ0µr

the permeability, and G(r, r′) is the dyadic homogeneous Green’s function. [50] provides a detailed

treatment.

35


Numerical Electromagnetic Code (NEC) [10] is a simulation tool that employs the Method Of

Moments (MOM) to solve this EFIE numerically.

Equation 3.5 appears similar to the generalized form,

L(I) = E (3.6)

where L is a linear operator, E is the excitation source (electric field), and I is an unknown function

(current).

In MOM, I is first expanded into a sum of N weighted basis functions,

I ≈N

∑n=1

anIn (3.7)

where an are unknown weighting coefficients, and since L is a linear operator, one can write,

N

∑n=1

anL(In) ≈ E (3.8)

The basis function (wire segment) in NEC is represented by,

anIn = An + Bnsink(s− sn)+Cncosk(s− sn) with |s− sn|< ln/2 (3.9)

where s is the position along the nth segment with center sn and length ln. NEC then solves for the

3N coefficients in equation 3.9 by applying boundary-conditions between segment junctions and at the

ends. A set of basis functions sum to give the overall current, as shown for a four-segment wire, e.g., in

Fig. 3.3.

NEC (TARA uses both NEC-2 and NEC-4) requires that the length of wire segments should not

exceed 1/10 th of the wavelength of operation to ensure accuracy. Tab. 3.2 gives the number of wire

segments used for each wire and the corresponding length per wire segment.

Element Segs Length (in)1 11 0.332 11 1.643 11 3.244 11 5.185 13 6.406 15 7.67

Table 3.2 Number of wire segments for each wire and the corresponding length per segment.

36

3.3 Antenna Characterization

Fig. 3.3 Current Basis Functions and their sum on a four wire segment. Taken from Ref. [10]


Frequency-independent performance requires that the antenna appear electrically the same to the excit-

ing wave at all frequencies. TARA expects cosmic rays in the 50 to 80 MHz band, and correspondingly

the antenna has the minimum number of elements required in their respective active regions for response

over the desired passband. The length of the LPDA elements constrains the maximum and minimum

possible frequencies.

fmin ≈c

4 · lmax≈ 50MHz and fmax ≈

c4 · lmin

≈ 135MHz (3.10)

The impedance of the antenna to a 50 Ω transmission line was measured in an anechoic chamber

and simulated using NEC4. The voltage standing wave ratio (VSWR) in terms of the magnitude of the

complex reflection coefficient (S11) is given by,

V SWR =1 + |s11|1−|s11|

, (3.11)

and is shown as a function of frequency in Fig. 3.4. A VSWR of 3.0 implies that 75% signal power is

transmitted from the antenna to the environment at a given frequency.

At time t, the electric field arriving at the antenna from some direction (θ ,φ ) is a transverse two-

component vector given by,

E(t) = eθ Eθ (t)+ eφ Eφ (t) (3.12)

37


50 60 70 80 90 100Frequency (MHz)

0.0

0.5

1.0

1.5

2.0

2.5

3.0

3.5

4.0

V.S.W.R

Measurement

Simulation

Fig. 3.4 Overlay of SWR of a horizontally polarized TARA LPDA as measured in the KU anechoicchamber compared with NEC simulation.

where θ and φ denote independent polarizations (the radiation electric field is only transverse and for

this reason has no component Er). Fig. 3.5 depicts the coordinate system. The field can be linearly,

circularly or elliptically polarized.

Of particular interest at some time t is the electric field in the antenna’s far field relative to that of

some isotropic antenna and as a function of angular space. The antenna power measured in dBi (decibels

w.r.t isotropic) is described as,

Pi(θ ,φ) = 10log10S(θ ,φ)

S(θ ,φ)i(3.13)

where,

S(θ ,φ) = [E2θ (θ ,φ)+ E2

φ (θ ,φ)]/Z0 (3.14)

is the Poynting Vector of an isotropic antenna in its far field and Z0 ∼ 377Ω the intrinsic impedance of

free space.

Fig. 3.6 shows measured points overlaid on a NEC model for the radiation pattern of the antenna in

the XY-plane. Transmission coefficient (S21) measurement give a measure of the total transmitted power

of a wave relative to an incident wave. These measurements were made using two identical horizontally

polarized LPDA’s at a separation of 5λ at the carrier frequency of 54.1 MHz. As the measurements

were relative, not absolute, a uniform scale factor was determined by minimizing χ2 between the model

and data.

Lowering a horizontally polarized dipole antenna from some height, again at the carrier frequency,

and measuring the received power at a horizontally polarized LPDA on the ground enabled confirmation

of the radiation pattern in the YZ-plane. As previously, a χ2 minimization between measurement and

simulation was performed. (See Fig. 3.7)

38


dφ

dθ

φ

θ

ρr

y

z

x

er

eφ

eθ

Fig. 3.5 Spherical Coordinate System used for defining antenna geometry.

0°

15°

30°

45°

60°75°90°105°

120°

135°

150°

165°

180°

195°

210°

225°

240°255° 270° 285°

300°

315°

330°

345°

Horizontal Gain (dBi)

129

50

510

Fig. 3.6 Radiation pattern in the XY-plane of a horizontally polarized TARA LPDA at the transmittersounding frequency of 54.1 MHz. Beamwidths (3 dB below peak gain) are shown with red lines. Blackpoints show field measurements.

39

3.4 Effective Height

180° 165°150°

135°

120°

105°

90°

75°

60°

45°

30°15°0°-15°

-30°

-45°

-60°

-75°

-90°

-105°

-120°

-135°

-150°-165°

Vertical Gain (dBi)

12

9

5

0

5

10

Fig. 3.7 Radiation pattern in the YZ-plane of a horizontally polarized TARA LPDA at the transmittersounding frequency of 54.1 MHz. Beamwidths (3 dB below peak gain) are shown with red lines. Blackpoints show field measurements.

In the case of the radiation pattern plot in the YZ-plane, a ground plane effect is seen due to the

superposition of the direct wave on the antenna and that of the wave reflected from the ground.

For a vertically polarized antenna, the radiation pattern is expected to be similar but switched in the

E(XZ) and H(YZ) planes.

A property associated with the radiation pattern is the half-power beamwidth, φHPBW and θHPBW in

the antenna E and H planes respectively, defined as the points where the power decreases by 3 dB. The

relationship between the half-power beamwidths and power is given by,

θHPBW φHPBW ∼= ΩA =∫

φ=2π

φ=0

∫θ=π

θ=0Pi(θ ,φ)sinθdθdφ , (3.15)

where ΩA is the beam area, defined as the solid angle through which most of the power radiates.

Fig. 3.8 and Fig. 3.9 show the horizontal half-power beamwidth and the front-to-back ratio of the

antenna measured in the KU anechoic chamber.


The relationship between the terminal voltage at the antenna feeds and the incident electric field [51] is

given by,

V = h · E = hEcosθ , (3.16)

where h is the effective height and θ the angle between the electric field polarization and the dipole.

40


50 60 70 80 90 100Frequency (MHz)

30

40

50

60

70

80

90

100

HorizontalBeamwidth

()

Fig. 3.8 Beamwidth of a horizontally polarized TARA LPDA, as measured in the KU anechoic chamber.

50 60 70 80 90 100Frequency (MHz)

5

0

5

10

15

20

25

30

35

Front/Back

Ratio

(dB

)

Measurement

Simulation

Fig. 3.9 Front-to-Back ratio of a horizontally polarized TARA LPDA, as measured in the KU anechoicchamber and an overlay of NEC simulation. The discrepancy is likely due to differences in the heightof the antenna while taking the measurement (∼ 6 f t.) vs. simulation (12 f t.).

41

3.5 Galactic Floor

A detailed discussion of the effective height is presented in [52] and is given by (see Fig. 3.10),

h( f ) = 2×√

G4π

c2

f 2ZZ0

(3.17)

where G is the unit-less gain, c the speed of light, f the frequency, Z the impedance at the antenna

terminals and Z0 the impedance of free space.

50 60 70 80 90 100Frequency (MHz)

0

1

2

3

4

5

6

7

8

Effective

Height

(m)

Measurement

Simulation

Fig. 3.10 Effective Height of the TARA LPDA.

3.5 Galactic Floor

In any RF receiver system, sensitivity is limited by the combination of external noise entering through

the antenna and internal noise from various sources like low-noise amplifiers and other resistive losses

from filters, cables, and couplers. The sky, earth, and antenna resistive loss generate noise entering the

antenna. Diffuse radio noise from the galactic plane is non-polarized and is the dominant noise source

in the TARA frequency band.

After accounting for the amplifier and instrumental gains and losses, the observed noise background

can be compared with the irreducible galactic noise background [53] across the passband. Specifically,

toward the South Galactic Plane the spectrum of specific intensity in units of Wm−2Hz−1sr−1 is given

by,

Iν = Igν−0.52 1− exp[−τ(ν)]

τ(ν)+ Iegν

−0.80exp[−τ(ν)], (3.18)

where ν is in MHz, the two terms Ig and Ieg correspond to galactic and extragalactic contributions

42

3.5 Galactic Floor

respectively and τ(ν) is the opacity in the polar direction. The parameter values are

Ig = 2.48×10−20, Ieg = 1.06×10−20,τ(ν) = 5.0ν−2.1. (3.19)

The apparent brightness temperature, Tb (figure 3.11), is then related by the Rayleigh-Jeans Law:

kTb = Iνc2/2ν2. (3.20)

50 60 70 80 90 100Frequency (MHz)

102

103

104

105

ApparentBrightnessTem

perature

(K)

Fig. 3.11 Apparent brightness temperature of the sky.

The system noise was first measured using the LPDA, 100ft of LMR-400 cable, a front end module

comprising a 61 dB amplifier (MITEQ AU - 1525), FM notch filter (Mini-Circuits NBSP - 108+), a

high pass and low pass filter and a HP 54542C oscilloscope sampling at 250 MSa/s with 32768 samples

being taken in so-called snapshots every 30 seconds. See Fig. 3.12.

+V

-V

250 MSa/s , 32768 Sa

100 ft LMR 400 Cable

Fig. 3.12 An experimental setup to measure average system noise. Snapshots were taken every 30seconds for 1 week.

43

3.5 Galactic Floor

The average measured system noise is calibrated by removing the effects of individual components

in the receiver RF chain from the average snapshot spectra to determine the absolute received power.

Without any scaling, the corrected received power compares well with the Galactic expectation [54]

(see Fig. 3.14). Principal components, for which adjustments are made, include filters and amplifiers

measured via the transmission coefficient, and also the transmission line using the manufacturer’s at-

tenuation per unit-length data (Fig. 3.13). Transmission line attenuation is found via,

Attenuation(dB)/100 f t = c·√

(F [MHz])+(0.000260)·F [MHz] (3.21)

where the constant c = 0.075550 for LMR - 600 and c = 0.122290 for LMR - 400 cable.

50 60 70 80 90 100Frequency (MHz)

10-1

100

101

Attenuation

(dB

per

100ft.)

LMR − 600

LMR − 400

Fig. 3.13 Attenuation in dB per 100 ft of the LMR - 600 and LMR - 400 transmission lines.

Anthropogenic noise sources are transient and noise is absent in the measurement band due to the

receiver site’s remote location. In this frequency region, galactic noise dominates thermal and other

noise sources.

44

3.5 Galactic Floor

40 50 60 70 80 90 100Frequency (MHz)

180

170

160

150

140

130

120

110

100

90

Pow

er (d

Bm/H

z)

CarrierFM Radio

MeasurementSimulation

Fig. 3.14 Average receiver system noise floor (green) Power Spectral Density (PSD) in dBm/Hz super-imposed with a fit to the known galactic background noise (red dashed line). System attenuation, filters,and amplifiers were accounted for in calculating the absolute received power.

45

Chapter 4

Cosmic Ray Detector

In order to detect cosmic ray chirps while also adding stereoscopic measurement capabilities in a

radio-quiet environment, receiver stations have been developed. These remote stations have a mostly

analog data acquisition system that has low power consumption and at a cost that is comparatively in-

expensive. This chapter includes discussions on the triggering logic, powering, and the communication

system and some specific details of hardware components.

4.1 Station Structure

The station is designed to detect chirp echoes that have low signal-to-noise ratios. A dual-polarized

LPDA (Chapter 3) feeds the two station channel inputs. The station consists of four basic components:

the Mixer Module, Chirp Acquisition Module (CAM), Transient Detector Apparatus (TDA), Current

Voltage Temperature (IVT) Board and the System Health Monitor (SHM). Fig. 4.1 shows these com-

ponents that are powered via solar panels, while a GHz transceiver provides communications with the

Longridge fluorescence detector station, which houses Internet communications. A discussion of these

components, as well as the powering and communications systems follow.

4.2 Mixer Module

The signal of interest is a chirp signal that has duration Tc seconds, a changing Amplitude A(t) with

start phase φ0, start frequency f0, stop frequency f1 and chirp rate κ Hz/s. Both up and down chirps are

treated identically in the detector. Fig. 4.2 is an example of the signal of interest. Assuming that it is

46

4.2 Mixer Module

Fig. 4.1 Construction schematic showing the different Components of the TARA detector. CourtesyKenneth Ratzlaff, Instrumentation Design Lab, KU.

centered around time t = 0, such a chirp signal can be written as,

s(t) = A(t)sin(φ0 + 2π f0t + πκt2) (4.1)

To detect the presence of the signal s(t) without prior knowledge of the chirp rate κ , the signal is

first down-converted to a monotone. To achieve this, the signal is mixed with a delayed copy of itself,

i.e s(t)⊗

s(t − τ) [55], as depicted in the radar block diagram in Fig. 4.3.

For an incident chirp signal, the non-linear components in the mixer result in a product term that

yields a monotone at a beat frequency

f = κτ, (4.2)

where τ is the delay time. This delay is created with LMR-600 cable, which introduces negligible losses.

This is illustrated in Fig. 4.4, where a -10 MHz/µs chirp is down-converted to a 1 MHz monotone (100

ft of LMR-600 cable was used here).

Signal "de-chirping" is done in the mixer module consisting of Radio Frequency (RF) components

as shown in Fig. 4.5. This mixer module receives both antenna channels.

The mixer module components for channel 1 are: a splitter (ZMSC-4-1-BR, Mini-Circuits), mixer

(ZX05-1L-S+, Mini-Circuits), filter (SLP-1.9+, Mini-Circuits), amplifier (ZFL-500LN+B, Mini-Circuits)

and bias tee (ZFBT-4RG-FT, Mini-Circuits). In addition to copies of the incoming signal used in de-

47

4.2 Mixer Module

Time (s)

Am

plitu

de (V

)

t0 t0+Tc

Time (s)

Freq

uenc

y (H

z)

t0 t0+Tc

f1

f0

slope=

Fig. 4.2 A linear down-chirp in the time domain (top) and frequency-timedomain (bottom).

48

4.2 Mixer Module

TOFPGA

+V

-V

+V

-V

BPF

63 - 77 MH

z

BPF

DC - 1.9 M

Hz

LO

RF

IF

Delay

Antenna

200 MSa/s

AD

C

AD

C

AD

C

AD

C

AD

C

4 MSa/s

4 MSa/s

4 MSa/s

4 MSa/s

100ns

BPF A

BPF B

BPF C

BPF D

Fig.4.3R

adarblock

diagramshow

ingthe

down-conversion

processalong

with

thefiltering

andenveloping

process.

49

4.2 Mixer Module

0 20 40 60 80 100Frequency(MHz)

180

160

140

120

100

80Po

wer

(dBm

)

0 20 40 60 80 100Frequency(MHz)

200

180

160

140

120

100

80

Pow

er(d

Bm)

Fig. 4.4 Top: Power spectrum of a -10 MHz/µs chirp created by a signal generator, prior to mixing.Bottom: power spectrum of a 1 MHz monotone signal after signal mixing and passing through alow pass filter. The chirp is evident as the left-most peak in this distribution. The 24 MHz peak andthe 48 MHz harmonic is likely due to a Serial Peripheral Interface(SPI).

50

4.3 Triggering Mechanism

chirping, additional copies are sent to a TDA and a High-Speed ADC.

For Channel 2, the mixer module consists only of a bias tee (ZFBT-4RG-FT, Mini-Circuits) leading

into a TDA.

Fig. 4.5 Block diagram showing the different components of the mixer module. Courtesy Kenneth Rat-zlaff, Instrumentation Design Lab, KU.

4.3 Triggering Mechanism

The expected value of chirp rates from EAS echoes are typically between -10 to -1 MHz/µs [9] (see

Chapter 2). Consequently, with 95 ns delay, the down-converted signal has a frequency between ∼100

kHz and 1 MHz. To trigger, the mixed signal is split into multiple copies. Each copy is then passed

through a bandpass filter and an envelope detector. Different frequency bands are then compared by

majority logic in a Field Programmable Gate Array (FPGA) , requiring no more than two bands to form

a trigger in order to suppress impulsive noise. Each of the frequency banded outputs corresponds to a

separate range of chirp rates.

To illustrate the triggering mechanism, in the oscilloscope-based example in Fig. 4.6, after mixing

and filtering, the signal is passed through a power detector (8471D; Agilent, Inc.). Here, a chirp with

51

4.4 Chirp Acquisition Module (CAM)

0 dB SNR at a rate of -1 MHz/µs is first band-pass filtered (41–100 MHz) and then amplified by 20

dB. Next, the signal is mixed and low-pass filtered (DC-1.9 MHz) and then passed through the Agilent

power detector.


The CAM is an embedded system with a modular design that provides hardware and software integra-

tion for chirp detection. It consists of five basic parts: Trigger Board, High-Speed Board, FPGA, Global

Positioning System (GPS) and a Single Board Computer (SBC) as shown in Fig. 4.7. A description of

each of these subsystems follows.

4.4.1 Trigger Board

Fig. 4.8 shows a general layout of the Triggering board.

The input to the trigger board is an SMA female connection. The input signal is first split into four

copies, and then each passed through a bandpass filter followed by an envelope detector (see Fig. A.2).

The bandpass filters are Butterworth Pi filters with passbands as given in Table. 4.1 and Bode plot in

Fig. 4.9.

Channel 3dB Low Pass Cut Off (KHz) 3dB High Pass Cut Off (KHz) 3dB Band-Width (KHz)A 344 18 326B 570 240 330C 864 549 315D 1117 818 308

Table 4.1 Passband of the Butterworth Pi Filters and the corresponding Channels.

The envelope detector consists of a rectifying diode followed by an RC combination that allows the

output waveform to follow the envelope of the signal [56]. Fig. 4.10 shows a simulated 50 KHz sine

wave signal passed through an envelope detector. The Figure illustrates the subsequent response in the

four different signal paths.

Next, the four channels are passed into the ADC (AD80066) [57], operating from a 5V supply.

The ADC is packaged in a Small-Shrink-Outline-28-Package (SSOP-28 package) [57] and nominally

consumes 490 mW. It runs in the Sample and Hold Amplifier (SHA) mode with the 16-bit output

multiplexed into 8-bit words and accessed in two read cycles clocking at 3MHz per channel. Fig. 4.11

shows the read operation timing diagram.

The ADC register is programmed via the Serial Peripheral Interface (SPI) Bus at 24MHz as shown

in the timing diagram in Fig. 4.12.

52


0 20 40 60 80 100Time(µs)

0.6

0.4

0.2

0.0

0.2

0.4

0.6

Am

plit

ude(V

)

0 20 40 60 80 100Time(µs)

0.15

0.10

0.05

0.00

0.05

0.10

0.15

Am

plit

ude(V

)

0 20 40 60 80 100Time(µs)

0.005

0.000

0.005

0.010

0.015

0.020

Am

plit

ude(V

)

Fig. 4.6 Top: 0 dB SNR and 1 MHz/µs chirp embedded in noise prior to "de-chirping". Second from top: The monotone signal after input chirp is mixedwith delayed copy of itself and passed through a low-pass filter. Bottom:Monotone passed through the Agilent 8471D power detector.

53


To Comms

SD StorageSBC

Triggering Board

GPS

FPGA

High Speed ADC Board

SPI

SPI

SPI

UART

PPS

Trig. SignalSignal

Fig. 4.7 Elements of the CAM unit showing the communications protocols.

Fig. 4.8 Trigger Board Schematics. Courtesy Rob Young, Instrumentation Design Lab, KU.

54


100 101 102 103

Frequency(KHz)

50

40

30

20

10

0

Magnit

ude (

dB

)

ABCD

Fig. 4.9 Bode plot for the four bandpass filters.

0 5 10 15 20Time(µs)

0

20

40

60

80

100

Am

plit

ude (

mV

)

50 KHz signal

ABCD

Fig. 4.10 A 50 KHz sine wave signal response from the envelope detector in each signal path.

The data are transferred via the FPGA Peripheral module (Pmod) interface for triggering (see

Sec. 4.4.3).

The GPS module (i-Lotus M12M) is commercially available and mounted on the Triggering Board

as well. Fig. A.3 shows the connections between the FPGA, SBC, and the GPS and Triggering Board.

55


Fig. 4.11 Read operation timing for the Analog Devices AD80066 ADC. 16 bits are multiplexed as two8-bit words at 3 MHz per channel.

Fig. 4.12 SPI Timing for the Analog Devices AD80066 ADC. SCLK is the 24MHz transfer rate clock;16 bits are transferred while valid programming i.e while SLOAD is pulled low.

4.4.2 High-Speed Board

The TARA High-Speed Board is a commercially available AD9634 Evaluation Board [58]. The board

consists of a 12-bit ADC sampling up to 250 MHz, with a total power consumption of 360mW. The

ADC clocks at 200 MHz and uses a 1.8V SPI port at 24 MHz for register programming and read-back

(see Fig 4.13).

Fig. 4.13 SPI Timing for the Analog Devices AD9634 ADC. SCLK is the 24MHz transfer rate clock; 8bits are transferred while valid programming, i.e. while SLOAD is pulled low.

Differential signaling provides superior common-mode noise rejection. 12-bit words are transferred

using a custom adapter (Fig. A.4) between the Evaluation Board’s Low Voltage Differential Signaling

(LVDS) parallel output port and the Very High Density Cable (VHDC) on the FPGA. In the next section,

we discuss the timing details.

4.4.3 FPGA

The CAM uses the Xilinx Spartan 6 FPGA based Nexsys 3 digital system board [11].

To configure the FPGA, the configuration file saved with a ’.mcs’ extension is first stored in a non-

volatile parallel Phase Change Memory (PCM) device. This is then transferred to the FPGA on power

up via the BPI-UP port. On the development board, this is one of the four possible configuration modes

and is achieved by removing all connections on the J8 jumper [11] (See Fig. 4.14)

56


Nexys3 Reference Manual

Doc: 502-182 page 2 of 22

Configuration After power-on, the Spartan-6 FPGA board must be configured (or programmed) before it can perform any functions. The FPGA can be configured in one of four ways: a PC can use the Adept "USB Prog" port to program the FPGA any time power is on; a configuration file stored in the non-volatile parallel PCM device can be transferred to the FPGA at power-on using the BPI-UP port; a file stored in the non-volatile serial (SPI) PCM device can be transferred to the FPGA using the SPI port; or a programming file can be transferred from a USB memory stick attached to the USB HID port. An on-board "mode" jumper (J8) selects between the programming modes as shown in the J8 Mode legend in the figure below. JTAG Mode can be accessed at any time without changing jumpers.

M0M1

JTAGPort

USB Controller

Micron SPI Quad mode PCM (P5Q)

1x6 JTAGHeader

SPIPortMicro-AB USB

Connector

Adept “USB Prog” Port

Spartan6Done

PIC24Type A USB Connector

Host PortSerialProg. Port

2

Micron Parallel PCM (P8P)

BPIPort

J8

Programming Mode

SLV Serial

SPI

BPI UPM0 M16-pin JTAG

Header (J7)

Prog

Programming files are stored in SRAM-based memory cells within the FPGA. This data defines the FPGA’s logic functions and circuit connections, and it remains valid until it is erased by removing board power, by pressing the reset button attached to the PROG input, or by writing a new configuration file using the JTAG port. FPGA configuration files transferred via the JTAG port use the .bin or .svf file types, files transferred from a USB stick use the .bit file type, and BPI or SPI programming files can use .bit, .bin, or .mcs types. The ISE/WebPack or EDK software from Xilinx can create bit, svf, bin, or mcs files from VHDL, Verilog, or schematic-based source files (EDK is used for MicroBlaze™ embedded processor-based designs). Digilent's Adept software or Xilinx's iMPACT software can be used to program the FPGA or ROMs using the Adept USB port. During JTAG programming, a .bit or .svf file is transferred from the PC to the FPGA using the Adept USB port. When programming a non-volatile PCM device, a .bit, .bin, or .mcs file is transferred to the in a two-step process. First, the FPGA is programmed with a circuit that can program PCM devices, and then data is transferred to the PCM device via the FPGA circuit (this complexity is hidden from the user – a simple “program ROM” interface is presented by the programming software. Note the PCM devices are next-generation Flash ROM devices, and they are often referred to as "Flash" or "ROM" memory). After the PCM device has been programmed, it can automatically configure the FPGA at a subsequent power-on or reset event as determined by the J8 jumper setting. Programming files stored in the PCM devices will remain until they are overwritten, regardless of power-cycle events. The FPGA can be programmed from a memory stick attached to the USB-HID port if the stick contains a single .bit configuration file in the root directory, the J8 Programming Mode jumper is set to

Fig. 4.14 Spartan 6 Programming via BPI port (BPI prom 28F128P30). Taken from [11].

The Nexys3 board includes a single 100MHz Complementary metal–oxide–semiconductor (CMOS)

oscillator connected to pin V10 [11](V10 is the GCLK0 input in bank 2). Phase Locked Loop (PLL) ,

and Digital Clock Manager (DCM) features on the board can be used to synthesize other frequencies.

The DCM is configured to produce two 48 MHz clocks (see Fig. 4.15).

CLK_IN

RESET

CLK_OUT1

CLK_OUT2

LOCKED

FPGA CLK (100 MHz)

48 MHz, 0° phase, 50% duty cycle616.667 ps pk-to-pk jitter, 150.000 ps phase error

48 MHz, 0° phase, 50% duty cycle616.667 ps pk-to-pk jitter, 150.000 ps phase error

Fig. 4.15 Nexsys 3 board’s 100 MHz CMOS oscillator and two 48 MHz clocks synthesized in theFPGA.

These are then further sub-divided to synthesize clocks for transfer and synchronization between

peripheral boards and other functions, as shown in Fig. 4.16 and Fig. 4.17.

The triggering board connects to the FPGA development board via the Pmod connector. These are

2x6 right-angle, 100-mil female connectors that mate with standard 2x6 pin headers [11]. Each set of

12-pin connector comprise eight logic signals and a pair of 3.3V VCC and ground signals (4.18) [11].

Fig A.3 shows the connections.

Words from the triggering board are de-serialized from the parallel input ports in the FPGA via the

implementation of a Finite State Machine (FSM). Once shifted in, a comparator is implemented, and

57


divide by 2divide by 2

24 MHz CLK 12 MHz CLK (3 MHz / Channel for AD80066)

4 bit Counter

48 MHz CLK

CDS CLK2 (CLK For AD80066 to Sample DataSimultaneously)

24 MHz CLK(FSM CLK For Triggering Data)

Inv.

Inv.

CDS CLK2 (CLK For Synchronizing Triggering Data And For Comparator)

Fig. 4.16 Schematic illustrating synthesis of secondary triggering clocks from a 48 MHz clock.

divide by 2

48 MHz CLK

Inv.

24 MHz CLK(FSM CLK For Configuration Of Triggering & High Speed Board)

24 MHz CLK(SPI CLK For AD80066 & AD9634)

Fig. 4.17 Synthesis of registry configuration and data transfer clocks from the 48 MHz clock.

8 signalsGNDVCC

Pin 1Pin 6

Pin 12

Fig. 4.18 Pmod connector with eight logic signals, two grounds and two 3.3V VCC per Pmod.

each channel triggers only if at least one of the following conditions are satisfied,

A>C and A> D (4.3)

B> D and (B> A or B>C) (4.4)

C > A and (C > B or C > D) (4.5)

58


A B C D (((¬A)∧ (¬B))⊕ ((¬C)∧ (¬D)))⊕ ((B∧C)⊕ (¬((B∧C)|(A∨D))))

0 0 0 0 00 0 0 1 10 0 1 0 10 0 1 1 10 1 0 0 10 1 0 1 00 1 1 0 10 1 1 1 01 0 0 0 11 0 0 1 01 0 1 0 01 0 1 1 01 1 0 0 11 1 0 1 01 1 1 0 01 1 1 1 0

Table 4.2 Truth Table for the four input channel triggering (see Fig. 4.9).

D> A and D> B (4.6)

along with the Threshold condition,

A> T hreshold or B> T hreshold or C > T hreshold or D> T hreshold (4.7)

which is imposed both for redundancy and also to ensure that no more than two channels initiate an event

trigger. All four channels are assessed as given in the truth table in Tab. 4.2 and shown in Fig. 4.19.

A snapshot trigger used to measure ambient noise is also implemented on 0xFF clock cycles of the

GPS Pulse Per Second (PPS) clock. A trigger can, for this reason, be registered either on a snapshot

trigger or an event trigger based on the four channel comparator, as shown in Fig 4.20.

To support LVDS signaling, signals to the VHDC connector are routed as matched pairs with the

corresponding fpga pins located in I/O bank0. These are powered at 2.5V [11]. There are twenty

matched pairs of data signals, twenty ground signals, and eight power signals on the VHDC connector

[11] (see Fig. 4.21).

The 200 MHz differential clock outputs from the High-Speed Board are brought into the FPGA and

used in frequency synthesis. The DCM is configured to produce two clocks at 200 MHz from these

input clocks (see Fig. 4.22).

The AD9634 ADC sends even/odd bits on the rising/falling edge of the sampling clock. A Double

59


A

B

C

D

B

C

B

C

A

D

Th

Trigger

Fig. 4.19 Event triggering logic for the four input signals (A,B,C,D) and Threshold (Th).

Snapshot

Event

Trigger

Fig. 4.20 Triggering logic for a comparator based event and a GPS PPS based snapshot trigger


Doc: 502-182 page 20 of 22

asserted, then a “7” will be displayed in digit position 2. If AN0 and CB, CC are driven for 4ms, and then A1 and CA, CB, CC are driven for 4ms in an endless succession, the display will show “17” in the first two digits. An example timing diagram for a four-digit controller is provided. Expansion Connectors The Nexys3 board has a 68-pin VHDC connector for high-speed/parallel I/O, and an 8-pin Pmod connector for lower speed and lower pin-count I/O. VHDC Connector The VHDC connector includes 40 data signals (routed as 20 impedance-controlled matched pairs), 20 grounds (one per pair), and eight power signals. This connector, commonly used for SCSI-3 applications, can accommodate data rates of several hundred megahertz on every pin. Both board-to-board and board-to-cable mating connectors are available. Data sheets for the VHDC connector and for mating board and cable connectors can be found on the Digilent website, as well as on other vendor and distributor websites. Mating connectors and cables of various lengths are also available from Digilent and from distributors. All FPGA pins routed to the VHDC connector are located in FPGA I/O bank0. The FPGA's bank0 I/O power supply pins and the VHDC connector's four Vcc pins are connected together by a small, segregated power supply plane in the PCB. This sub-plane can be connected to 2.5V or 3.3V, depending on the position of jumper JP8. This arrangement allows peripheral boards and the FPGA to share the same Vcc and signaling voltage across the connector, whether it be 3.3V or 2.5V. The unregulated board voltage VU5V0 (nominally 5V) is also routed to four other VHDC pins, supplying up to 1A of additional current to peripheral boards. A second jumper (JP4) allows the unregulated board voltage to be disconnected from the VHDC connector if desired. All I/O’s to the VHDC connector are routed as matched pairs to support LVDS signaling, commonly powered at 2.5V. The connector uses a symmetrical pinout (as reflected around the connector's vertical axis) so that peripheral boards as well as other system boards can be connected. Connector pins 15 and 49 are routed to FPGA clock input pins.

VUPin 34

Pin 68

Pin 1:IO1-P

Pin 35:IO1-NVCC

10 Matched Pairs 10 Matched Pairs

Clock Inputs


Doc: 502-182 page 20 of 22

asserted, then a “7” will be displayed in digit position 2. If AN0 and CB, CC are driven for 4ms, and then A1 and CA, CB, CC are driven for 4ms in an endless succession, the display will show “17” in the first two digits. An example timing diagram for a four-digit controller is provided. Expansion Connectors The Nexys3 board has a 68-pin VHDC connector for high-speed/parallel I/O, and an 8-pin Pmod connector for lower speed and lower pin-count I/O. VHDC Connector The VHDC connector includes 40 data signals (routed as 20 impedance-controlled matched pairs), 20 grounds (one per pair), and eight power signals. This connector, commonly used for SCSI-3 applications, can accommodate data rates of several hundred megahertz on every pin. Both board-to-board and board-to-cable mating connectors are available. Data sheets for the VHDC connector and for mating board and cable connectors can be found on the Digilent website, as well as on other vendor and distributor websites. Mating connectors and cables of various lengths are also available from Digilent and from distributors. All FPGA pins routed to the VHDC connector are located in FPGA I/O bank0. The FPGA's bank0 I/O power supply pins and the VHDC connector's four Vcc pins are connected together by a small, segregated power supply plane in the PCB. This sub-plane can be connected to 2.5V or 3.3V, depending on the position of jumper JP8. This arrangement allows peripheral boards and the FPGA to share the same Vcc and signaling voltage across the connector, whether it be 3.3V or 2.5V. The unregulated board voltage VU5V0 (nominally 5V) is also routed to four other VHDC pins, supplying up to 1A of additional current to peripheral boards. A second jumper (JP4) allows the unregulated board voltage to be disconnected from the VHDC connector if desired. All I/O’s to the VHDC connector are routed as matched pairs to support LVDS signaling, commonly powered at 2.5V. The connector uses a symmetrical pinout (as reflected around the connector's vertical axis) so that peripheral boards as well as other system boards can be connected. Connector pins 15 and 49 are routed to FPGA clock input pins.

3.3V 2.5V

Spartan 6 VHDC

VCCO

20 Matched Pairs

JP8

VU5V0

Bank0

JP4

Fig. 4.21 Right: All FPGA pins routed to the VHDC connector are located in FPGA I/O bank0. Left:40 data signals, 20 ground signals, and 8 power signals are found on the VHDC connector. Takenfrom [11].

60


CLK_P

CLK_N

RESET

CLK_OUT1

CLK_OUT2

LOCKED

ADC CLKs (200 MHz)

200 MHz, 50% duty cycle300 ns pk-to-pk jitter, 150 ps phase error

180° Phase

0° Phase

Fig. 4.22 High-Speed Board’s 200 MHz differential clocks and subsequent FPGA synthesis of 200 MHzclocks.

Data Rate (DDR) interface [59] is used between the FPGA and the ADC, as illustrated in Fig. 4.23.

C0C1

d d+1 d+2 d+3 d+4 d+5 d+6 d+7D

d d+2 d+4 d+6Q0

d-1 d+1 d+3 d+5Q1

Fig. 4.23 Input Double Data Rate (IDDR2) is implemented to set even/odd bits on the clock ris-ing/falling edge.

The bits are aligned per the C0 non-inverted clock, synthesized by use of an extra flip-flop as in

Fig. 4.24.

Input BUFfer Differential Signaling (IBUFDS) is a differential I/O primitive [59] that is instantiated

for the signals from the ADC and have two pins for the P and N channels of the differential signal (see

Fig. 4.25).

Bits are then serialized via an FSM and then stored in a circular buffer. The buffer has Random

Access Memory (RAM) depth of 213(8192) words and width of 16 bits. The actual width is 12 bits but

padded with four bits (bits 0,7,8,15 are redundant). Buffer read/write operations are at 200 MSa/s. The

Xilinx Block Memory Generator Wizard [60] generates the writing, reading and memory resources,

where one 9K Block RAM (BRAM) and seven 18K BRAMs are used (see Fig. 4.26).

Once a trigger is received, the 16-bit words of 213(8192) bit depth are written into a First In First

61


D D Q

D Q D Q

Q0

Q1

C0

C1

Fig. 4.24 IDDR2 when DDR alignment is C0.

I

IBO

Fig. 4.25 Differential Input Buffer Primitive.

Out (FIFO) buffer at 200 MSa/s. The SBC reads the 8-bit words of 214(16384) bit depth at 8 MSa/s (SPI

clock), i.e. it has a non-symmetric aspect ratio. The Xilinx FIFO Generator Wizard [61] instantiates

the writing, reading and memory resources, where eight 18K BRAMs and First Word Fall Through

(FWFT) are used (see Fig. 4.27).

Fig. 4.28 depicts the circular buffer to FIFO process.

In addition to the signal captured by the FPGA from the High-Speed Board, the Triggering infor-

mation is also saved and forms the header of the data (see Fig 4.29). The header replaces the first 8 bits

of raw data and subsequently invalidates the next 8 bits.

The SBC (Raspberry Pi rev 2.) sets the triggering threshold with transactions initiated using a

different chip select line than the one activated once an interrupt is initiated on the FPGA side. This

interrupt initiates when the FIFO buffer is full. However, due to latencies arising due to the crossing of

clock domains in the FIFO, three synchronous D flip-flops as shown in Fig. 4.30 are instantiated after

the assertion of full on the FIFO. An FSM controls the transfer logic between the FPGA and SBC.

In all transfers to/from the SBC(Master) and FPGA(Slave), an SPI clock at 8 MHz is used (see

Fig. 4.31).

62


CLK_A

ADDRB[12:0]

DOUTB[15:0]FPGA CLK (200 MHz)

CLK_B

FPGA CLK (200 MHz)

WEA[0:0]

DINA[15:0]

ADDRA[12:0]

Fig. 4.26 Block Memory Signals.

WR_EN

FULL

FPGA CLK (200 MHz)

DIN[15:0]

WR_CLK

RESET

RD_CLK

SBC CLK (SPI*)

DOUT[7:0]

RD_EN

EMPTY

Fig. 4.27 FIFO Signals.

63


Read Ptr. (200 MSa/s,depth : 8192 words, width : 16 bits)

Write Ptr. (200 MSa/s,depth : 8192 words, width : 16 bits)

Write Ptr. (200 MSa/s,depth : 8192 words, width : 16 bits)

Read Ptr. (8 MSa/s,depth : 16384 words, width : 8 bits)

FIFO write on Trigger

True Dual Port RAM

First In First Out

Fig. 4.28 Simple Dual Port RAM and FIFO Implementation

Triggered Data(16382 words of 8 bit width)

Header(8 Bits)

01234567

Header - Bit 0 : Forced, Bit 1 : D, Bit 2 : C, Bit 3 : B, Bit 4 : A, Bit 5 : Event Trigger

Data - 16382 words (of 8 bit width) where bits 0 & 7 are redundant. 8191 words (of 12 bit width) are valid ADC data.

Invalid Data(8 Bits)

16384 Words (8 bit width)

Fig. 4.29 Triggering Information and ADC Data.

4.4.4 GPS

The GPS unit (i-Lotus M12M) [62] is mounted on the Triggering Board for stability and power (see

Fig A.3). A Universal Asynchronous Receiver/Transmitter (UART) is used to transfer data from/to the

SBC. The GPS PPS is used as a counter to obtain snapshot triggers as described earlier in Sec. 4.4.3.

Fig. 4.32 shows the Functional Block diagram of the GPS to SBC and FPGA connections.

Timing information is periodically queried from the GPS unit and used to update the time on the

SBC.

D

CLK

D

CLK

D

CLK

Q Q QFULL INTERRUPT

Fig. 4.30 Implementation of D Flip Flop to assert an Interrupt signal.

64


MIS

OM

OSI

INTER

RU

PTNex

sys

3 Spa

rtan

6 (

Sla

ve)

Raspberry Pi (M

aster)

Ban

k 3

(PM

OD

B)

Ban

k 2

(PM

OD

A)

CS0

CS1

CLK

RST

JA1

JB1

JB2

JB3

JB4

JA4

JA3

Fig. 4.31 FPGA - SBC Interface

TxD

RxD

SDA

PPS

TxD

RxD

SDA

Bank 2 (PMOD A)

JA2

i-Lotus M12MRaspberry Pi

Nexsys 3 Spartan 6

Fig. 4.32 GPS to FPGA and SBC Interface.

4.4.5 Single Board Computer (SBC)

The Single Board Computer used in the CAM is the commercially available Raspberry Pi Model B

Rev. 2 (RPi - B Rev 2). The RPi has the Broadcom BCM2835 SoC (System on Chip) that includes

the ARM1176JZF-S 700 MHz processor, VideoCore IV GPU, 512 MB RAM and access to other I/O

peripherals, and a separate three port USB Hub [63] (see Fig 4.33).

The RPi has a 26 (2× 13) pin 2.54 mm expansion header with 8 General Purpose Input/Output

(GPIO) pins. It also has dedicated peripherals such as SPI, UART along with 3.3V, 5V and GND

65

4.5 System Monitoring and Powering

RAM

I/O CPU/GPU USB HUB

2 X USBEthernet

Fig. 4.33 Functional Block Diagram of the Raspberry Pi Model B Rev. 2

GND

MOSIMISOSCLKCS0CS1

RxD

TxD

SD

A

INTE

RRU

PT(#

23)

RST(

#24

)

GN

D

GPS

SPI

Fig. 4.34 Adafruit Prototyping Pi Plate showing the corresponding GPIO pins to FPGA and GPS

supply lines.

The board consists of a 32 GB SD card to which data are stored and also where a minimal Rasp-

bian Wheezy (118 MB image) is compiled with hard float support (3.6.11+ hardfp kernel). WiringPi

written in C for the BCM2835 provides access to the GPIO. The Adafruit Prototyping Pi Plate provides

convenient pin access to connect with the FPGA and GPS (see Fig. 4.34).


In addition to the CAM, the Remote Station comprises several components to power the station and

monitor environmental variables. These components are the TDA Board, IVT Board, and the SHM.

In the Summer of 2013, a prototype station comprising just these components had been deployed to

66


acquire environmental and powering information. In the current station, modifications were made, and

the system was updated to integrate the CAM. A detailed description of the current powering and

operational systems follow.

4.5.1 TDA

The TDA continuously monitors ambient noise at the site by counting threshold-crossings. The thresh-

olds are set with Digital-to-Analog Converters that are remotely controlled through the System Health

Monitor (SHM). The measurement period, usually 10 seconds, is also controlled by the SHM. Fig. 4.35

shows a functional block diagram of the TDA. The TDA communicates via the Local Interconnect

Network (LIN) bus, designed with the robustness needed for automobile environments.

Fig. 4.35 Transient Detector Apparatus (TDA). Courtesy Kenneth Ratzlaff, Instrumentation Design Lab,KU.

4.5.2 IVT Board

The IVT board provides power measurements to monitor the Photo-Voltaic (PV) battery-charging

system. Voltage and current are measured for both the PV power input and the Load power output

(Fig. 4.36). On the board, a temperature sensor is located to enable correction for the current sensors.

Heavy screw lugs accommodate the heavy cables that are required for currents that typically reach 10

Amps. As was the case for the TDA, a micro-controller provides analog-to-digital conversion, averages

over the measurement period and communicates with the SHM over the LIN bus.

67

4.6 Chirp Calibration Unit (CCU)

Fig. 4.36 Current Voltage Temperature (IVT) Board. Courtesy Kenneth Ratzlaff, Instrumentation De-sign Lab, KU.

4.5.3 SHM

The SHM has the functions of reporting on the environmental and operational status of the system and

of controlling the charging of the batteries (Fig. 4.37). The SHM gathers ambient noise data from the

TDA and power data from the IVT over a period that is programmable from 5 seconds to at least 10

minutes. Similarly, it measures the battery voltage using the ADC on the micro-controller. A real-time

clock with battery backup provides a time-stamp for each measurement cycle.

The SHM also controls the charging of the battery. A high-current solid-state relay (SSR) can switch

PV current off when the battery reaches capacity, and another SSR is used to turn the load off if the

battery voltage is too low. As illustrated in Fig. 4.38, hysteresis levels minimize oscillations. The thresh-

olds are remotely programmable. If, however, the voltage level becomes too low, the communications

channel also shuts down until the battery becomes sufficiently charged.

Data are stored on board using an SD card and are also sent to a communication channel. An

Ethernet appliance provides both raw TCP connection and a website. The data are streamed via the

raw TCP port, and the website provides a data display and interactive control of operational variables

including the battery charging thresholds, the TDA thresholds, the measurement period, and calibration

values.


In order to calibrate the remote stations independently, a Chirp Calibration Unit has been deployed in

the field. Fig. 4.39 shows this unit that comprises a ’fat - dipole’, an Arduino with a monostable 555

68


Fig. 4.37 System Health Monitor (SHM). Courtesy Kenneth Ratzlaff, Instrumentation Design Lab, KU.

circuit, and a signal generator (SF1020). The output of the control circuit controls an SSR, which in

turn then triggers the signal generator. Fig. A.5 shows the control circuit for the chirp generator.

The dipole is tuned to a frequency of 70 MHz, however, in order to transmit a broadband chirp-like

signal the radius of the dipole was optimized to transmit efficiently up to 50 MHz. Stable chirps at

1 PPS, spanning 80 to 50 MHz and of a 20 µs duration every 2 hours for 10 seconds, i.e., ten chirp

signals, are produced.

69


Fig. 4.38 Threshold Settings on the System Health Monitor (SHM).

20 dB Attn SF 1020 Signal Generator

Battery

10 WPV Panel

Controller Circuit Arduino

Fig. 4.39 Schematic of the Chirp Calibration Unit (CCU) Courtesy Steven Prohira, Dept. of Physicsand Astronomy, KU.

70

Chapter 5

TARA Remote Station Detection ofUHECR’s

In this chapter, we discuss measurements taken using the Telescope Array Radar Remote Stations.

5.1 Snapshot Triggers

It is essential to understand the environment into which an observable cosmic ray may be embedded.

As mentioned previously, this was accomplished using snapshot triggers with the station prototype, de-

ployed in 2013. Snapshot triggers are taken every 255 seconds to sample the ambient noise at the remote

station location. These triggers are instantiated after 255 continuous pulses from the GPS clock pulsing

at 1 PPS. Since these snapshot triggers are unassociated with actual high-amplitude backgrounds, they

offer the opportunity to probe otherwise weak, but constant radio-frequency signals, such as that of the

Milky Way.

As the Earth rotates, the location of the Galactic center relative to the antenna heading (78 East of

North) changes, leading to a smooth variation in the measured noise floor. This variation is depicted

in Fig. 5.1, where both a fit to the measured signal centered at 70 MHz over a 5 MHz wide band is

superimposed along with the expected modulation by the Galactic Center on the measured signal at 70

MHz. We observe an apparent correlation in the phase of the known Galactic motion with bore-sight

of the receiver antennas, indicating good sensitivity to otherwise-weak signals. The measurements

made were taken from the horizontally polarized channel during the period from the 1st to the 14th of

December, 2014.

The behavior is seen at other frequencies as shown in Fig. 5.2. Here the relative angle is the differ-

ence between the azimuthal angle of the galactic center and the heading of the TARA LPDA.

71

5.2 Expected Event Rate

Dec 03 2014 Dec 05 2014 Dec 07 2014 Dec 09 2014 Dec 11 2014 Dec 13 2014Date (UTC)

166.5

166.0

165.5

165.0

164.5

164.0

163.5

163.0

PSD

(dBm/Hz)

τ (fit) =86176.71s, φ0 (fit) =−37.6

τgal (fit) =86146.09s, φ0, gal (fit) =−54.54

RS2 (70 MHz, 2Hr Resampling)

Fig. 5.1 Modulation of the RS2 antenna at 70 MHz. The green line is a least squares fit to the measuredpower while the orange line is a least squares fit to the expected modulation due to the galactic center,chosen to be Sagittarius A*. The measurements were taken from the horizontally polarized channel.

The measured noise level of Remote Station 2 superimposed on the expected Galactic noise floor

is shown in Fig. 5.3. The measurement again confirms that the Remote Stations are functioning well

enough to observe a high-energy cosmic ray if our putative radar model is correct.

5.2 Expected Event Rate

To quantify our expected signal event rate using the technique of radar reflections, an extensive library of

CORSIKA simulated showers were generated to calculate the lengths and radii of plasma. These were

estimated based on the energy deposited by the showers and the subsequent ionization electron density

in Chapter 2. Determining the RCS of a thin wire at various angles and obtaining the received power

from the bi-static radar equation using numerous locations within the baseline provides a distribution

of RCS vs. received chirp power. Given the TARA geometry and detector design the efficiency for

≥ 1019 eV has been calculated to be ∼ 10% based on the cumulative distribution function shown in

Fig. 5.4. Here a cut was made at 0 dB SNR based on the received signal power relative to the receiver

noise floor.

Fig. 5.5 shows the expected integrated event rate for the TARA Remote Station volume from 10 to

100 EeV. Here, we take dN/dE/dt/dΩ from the Telescope Array measurements and take into account

the efficiency of the detector and the flux of UHECR.

72

5.3 Calibration Events

02Dec2014

03 04 05 06 07 08 09 10 11 12 13

Date (UTC)

170

168

166

164

162

PSD

(dBm/Hz)

Remote Station 2

70 MHz

77 MHz

50

0

50

100

150

200

250

Relative

Angle

()

Galactic Center

Fig. 5.2 Modulation of the RS2 antenna at 70 MHz and 77 MHz. The orange line is a least squares fitto the expected modulation due to the galactic center, chosen to be Sagittarius A*. The measurementswere taken from the horizontally polarized channel.

5.3 Calibration Events

The Chirp Calibration Unit (CCU) generates ten chirps every 2 hours. These chirps span from 80 to 50

MHz in 20 µs for an average slope of 1.5 MHz/µs. The calibration unit was placed ∼ 140 f t from the

two remote stations and subsequently used to trigger them. Fig. 5.6 shows the spectrogram of such a

chirp.

During the period from 20th to the 24th of January, 2015, both force-triggered and self-triggered

events were collected by the CAM. The CAM tags an event as either a force-triggered or self-triggered

event based on the header information received. The events during this period classified as such are

shown in the SNR vs. Chirp-rate plot in Fig. 5.7.

The SNR here is the ratio of the maximum signal power of an event to the RMS of noise collected

in the last 10µs of the event, so as to prevent signal contamination. The chirp rate is calculated using

the "de-chirping" algorithm described earlier, namely by taking the product of the signal with a delayed

copy of itself. The delay was chosen to be 1 µs, and the bin with the maximum signal power below

20 MHz/µs was selected to obtain the chirp rate for the event. The CCU triggered events are clearly

identifiable as a ∼ 1.48 MHz/µs line in the plot ranging from a SNR of ∼ 25 to 30 dB. However, the

forced triggered data also has comparable SNR due to interference from the Transmitter signal at 54.1

MHz. Applying a 5th order Butterworth bandpass filter (58 - 82 MHz) as shown in Fig. 5.8 mitigates

the problem, clearly indicating the need for further carrier attenuation at 54.1 MHz as part of future

upgrades.

73

5.4 Constraints on the Radar Cross Section

50 55 60 65 70 75 80 85Frequency (MHz)

200

180

160

140

120

100

80

Pow

er (d

Bm

/Hz)

Remote Station 2 (12/01/2014 − 12/14/2014)

MeasurementSimulation

Fig. 5.3 The measured noise floor of the TARA LPDA using the Remote Station 2 DAQ, with theexpected noise floor superimposed (primarily galactic at these energies). The measurements made weretaken from the horizontally polarized channel.


In the case, that there are no cosmic ray detections, we have investigated our numerical sensitivity to

the radar cross-section. The efficiency of the TARA detector can be folded into the known TA flux of

UHECR such that, for a 90% C.L, the upper limit on the Radar Cross Section (σlimit) can be obtained

from σlimit = σpred × Events ObservedEvents Expected , where σpred is the expected RCS, Events Observed are the number

of cosmic ray events detected by the Remote Station and Events Expected are the number of events

expected accounting for detector inefficiencies. Fig. 5.9 shows the RCS limit assuming a particular

number of events detected for a live time of 1 year.

74


220 200 180 160 140 120 100 80 60 40

Power (dBm)

0.0

0.2

0.4

0.6

0.8

1.0

1.2

Cumulative

Distribution

Power Transmitted ∼40 Kilowatts

1019 eV

1020 eV

noise floor

Fig. 5.4 The cumulative distribution function for the TARA remote station receivers for an ensembleof Monte Carlo CORSIKA-simulated events at 1019 eV (green) and 1020 eV (blue). The pink bandindicates a bandwidth of 24 MHz up to 200 MHz. Cuts at 0 dB SNR are made on the received signalrelative to the receiver noise floor, assuming this is our future operational threshold.

101 102

Energy(EeV)

10-1

100

101

102

103

104

IntegratedEventRate

(yr−

1)

TARA Event Rate

Fig. 5.5 The expected integrated TARA remote station event rate based on the procedure described inthe text.

75


0 5 10 15 20 25 30 35 40Time (µs)

20

40

60

80

100

Frequency

(Mhz)

160

150

140

130

120

110

100

90

80

Power

(db)

Fig. 5.6 Remote Station triggered chirp (data). Calibration chirp sent from CCU in the field.

0 5 10 15 20Chirp − Rate (MHz/µs)

10

15

20

25

30

Signal−to−Noise

Ratio

(dB

)

01/20/2015 − 01/24/2015

self−triggeredforce−triggered

Fig. 5.7 Self and force-triggered events during the period from 20th to the 24th of January, 2015 forRS2, prior to software-filtering contamination by the carrier at 54.1 MHz.

76


0 2 4 6 8 10 12 14Chirp − Rate (MHz/µs)

10

15

20

25

30

35

Signal−to−Noise

Ratio

(dB

)

01/20/2015 − 01/24/2015

self−triggeredforce−triggered

Fig. 5.8 Self and Force-triggered events during the period from the 20th to the 24th of January, 2015for RS2 after passing through a 5th order Butterworth bandpass filter (58 to 82 MHZ) to suppressout-of-band noise.

50 100 150 200Observed No. of Events

0

50

100

150

200

σlim.

(cm

2)

σpred.=200cm2 , live−time=1yr

Energy ≥ 1019 eV

+10 dB

0 dB

−10 dB

Fig. 5.9 The Radar Cross Section limit given a certain number of observed events.

77

Chapter 6

Conclusion

The Telescope Array RAdar (TARA) detector is an ambitious project based on a remote sensing

technique known as bistatic radar that aims to probe large portions of the Earth’s surface in the quest

to measure cosmic ray induced radio echoes. Together with the Telescope Array in radio-quiet western

Utah, TARA’s pilot receiver and transmitter stations were deployed in 2011. This initial deployment

gave us insight into the detectability of air shower radar echoes. The receiver stations comprised an array

of Log Periodic Dipole Antennas. An oscilloscope-based data acquisition system was implemented for

noise calibration, including tracking galactic noise as the galactic plane migrates through the sky. That

experience laid the foundation for upgrades, including the construction of a dedicated transmitter station

in the summer of 2013 and deployment of additional remote receiver stations.

A prototype remote station was deployed in March of 2013, and for a period of one year collected

invaluable data regarding power budgeting and the ambient noise conditions in the Utah desert. These

results led to the eventual deployment of the first Remote Station in June of 2014 and was upgraded via

the addition of a second station in November of the same year.

At the core of these stations, an FPGA-embedded system provides logic for triggering and storing

raw data sampled at 200 MSa/s. These stations reduce costs relative to off-the-shelf commercial options

and add stereoscopic measurement capabilities for UHECR events.

In mid-January of 2015 a Chirp Calibration Unit was installed, and the detection of transmitted

chirps have demonstrated the working principle of the Remote Stations. Valuable insight has been

gained for future improvements on these stations including improvement in timing resolution for stereo-

scopic measurements. Going forward, data taken for a live-time of the order of months may eventually

provide sufficient statistics to answer some of the most intriguing questions in our Universe. Under-

standing the composition, anisotropy and energy of particles continuously bombarding Earth will ulti-

mately lead to an understanding of the most violent processes in the Cosmos.

78

Astroparticle physics, at the intersection of astronomoy, cosmology and particle physics has a con-

tinuosly changing landscape. The TARA project implements a novel technique to detect cosmic ray

chirps and similar techniques are being proposed to detect cosmic neutrinos [64]. The remote stations

implementation for the detection of chirp signals may prove to be useful in similar applications.

79

References

[1] Hanlon, W., 2014. URL http://www.physics.utah.edu/~whanlon/spectrum.html.

[2] J.W. Cronin. The highest energy cosmic rays. Nucl. Phys. Proc. Suppl., 138:465–491, 2005.

[3] Neto, J. (for the Pierre Auger Collaboration). Anisotropy studies with the pierre auger observatory.J. Phys.: Conf. Ser., 409 012108, 2013.

[4] Aab, A. et al. Searches for anisotropies in the arrival directions of the highest energy cosmic rayssearches for anisotropies in the arrival directions of the high energy cosmic rays detected by thepierre auger observatory. arXiv:1411.6111 [astro-Ph.HE], 2014.

[5] Abbasi, R.U et al. Indications of intermediate-scale anisotropy of cosmic rays with energy greaterthan 57 eev in the northern sky measured with the surface detector of the telescope array experi-ment. ApJ, 790(2), 2014.

[6] Kachelreiss, M. Lecture notes on high energy cosmic rays. arXiv:0801.4376 [astro-Ph], 2008.

[7] Stasielak, J. et al. Feasibility of radar detection of extensie air showers. arXiv:1411.7295 [astro-ph.IM], 2014.

[8] Bugallo, M.F. et al. Mariachi : A multidisciplinary effort to bring science and engineering to theclassroom. IEEE International Conference on Acoustics, Speech and Signal Processing., pages2661–2664, 2008.

[9] Abbasi, R.U et al. Telescope Array Radar (TARA) observatory for ultra-high energy cosmic rays.Nucl. Inst. and Meth. in Phys. Research, A, 767:322–338, 2014.

[10] G. Burke and A. Poggio. Numerical electromagnetic code (nec) method of moments, parts i, ii,iii. Technical report, Lawrence Livermore National Laboratory, NEC-1 (1977), NEC-2 (1981),NEC-3 (1983).

[11] Digilent. Nexsys3, 2014. URL http://www.digilentinc.com/Data/Products/NEXYS3/Nexys3_rm_V2.pdf.

[12] Hess, V. F. uber beobachtungen der durchdringenden strahlung bei sieben freiballonfahrten.Physikalische Zeitschrift, 13:1084, 1912.

[13] Smida, R. Cosmic-Ray Physics with the Pierre Auger Observatory. PhD thesis, Charles Universityin Prague, 2009.

[14] Millikan, R. A. and Cameron, G. H. High frequency rays of cosmic origin iii. measurements insnow-fed lakes at high altitudes. Phys. Rev., 28(851), 1926.

80

http://www.physics.utah.edu/~whanlon/spectrum.html

http://www.digilentinc.com/Data/Products/NEXYS3/Nexys3_rm_V2.pdf

http://www.digilentinc.com/Data/Products/NEXYS3/Nexys3_rm_V2.pdf

References

[15] Hoerandel, J.R. Models of the knee in the energy spectrum of cosmic rays. Astropart. Phys., 21:241–265, 2004.

[16] Khan, E. et al. Photodisintegration of ultra-high-energy cosmic rays revisited. Astropart. Phys.,23(2):191–201, 2005.

[17] Barnhill, D.S. Composition Analysis of Ultrahigh Energy Cosmic Rays Using the Pierre AugerComposition Analysis of Ultra High Energy Cosmic Rays Using the Pierre Auger ObservatorySurface Detector. PhD thesis, UCLA, 2005.

[18] Greisen, K. End of the cosmic-ray spectrum? Phys. Rev. Lett., 16(17):748–750, 1966.

[19] Zatespin, G. T. et al. Upper limit of the spectrum of cosmic rays. J. of Exp. and Theor. Phys. Lett.,4:78–80, 1966.

[20] Stanev, T. High Energy Cosmic Rays. Springer, 2003.

[21] Bergstrom, L. and Goobar, A. Cosmology and Particle Astrophysics. Springer, 2006.

[22] Bird, D. J. et al. Evidence for correlated changes in the spectrum and composition at extremelyhigh energies. Phys. Rev. Lett., 71(21), 1993.

[23] Boghrat, P. Search for Ultra High Energy Cosmic Ray Anisotropy with Auger. PhD thesis, UCLA,2007.

[24] Dolag, K. et al. Mapping deflections of extragalactic ultra-high energy cosmic rays in magneto-hydrodynamic simulations of the local universe. JETP Lett., 79:583–587, 2004.

[25] Aglietta, M. et al. Anisotropy studies around the galactic centre at eev energies with the augerobservatory. Astropart. Phys., 27:244, 2007.

[26] Ajello, M. et al. The 60 month all-sky burst alert telescope survey of active galactic nucleus andthe anisotropy of nearby agns. ApJ, 749(21), 2012.

[27] Kotera, K. and Olinto, A. The astrophysics of ultrahigh-energy cosmic rays. Annu. Rev. Astron.Astrophys., 49:119–153, 2011.

[28] Apel, W.D et al. The spectrum of high-energy cosmic rays measured with KASKADE-Grande.Astropart. Phys., 36:183–194, 2012.

[29] Abreu, P. et al. Measurement of the proton-air cross section at 57 tev with the pierre auger obser-vatory. Phys. Rev. Lett., 109, 2012.

[30] Abraham, J. et al. Measurement of the depth of maximum measurement of the depth of extensiveair showers above 1018 ev. Phys. Rev. Lett., 104, 2010.

[31] Abbasi, R.U et al. Study of ultra-high energy cosmic ray composition using telescope array’smiddle drum detector and surface array in hybrid mode. arXiv:1408.1726 [astro-ph.HE], 2014.

[32] Lemoine, M and Waxman, E. Anisotropy vs chemical composition at ultra high energies. J.Cosmol. Astropart., 11:9, 2009.

[33] Auger, P. et al. Extensive cosmic-ray showers. Rev. Mod. Phys., 11:288–291, 1939.

81

References

[34] Gaisser, T.k. and Hillas, A.M. Reliability of the method of constant intensity cuts for reconstruct-ing the average development of vertical showers. 15th Int. Cosmic Ray Conf., 8:353, 1977.

[35] Ostapchenko, S. Qgsjet - ii : towards reliable description of very high energy hadronic interactions.Nucl. Phys. B - Proc. Suppl., 151(1):143–146, 2005.

[36] Heck, D. et al. CORSIKA : A monte carlo code to simlate extensive air showers. Technical report,Karlsruhe Institute of Technology, 1998.

[37] Brau, C. A. Modern Problems In Classical Electrodynamics. Oxford University Press, 2004.

[38] Takai, H. et al. Forward scattering radar for ultra high energy cosmic rays. In 32nd InternationalCosmic Ray Converence, 2011.

[39] Gorham, P. On the possibility of radar echo detection of ultra-high energy cosmic ray- andneutrino-induced extensive air showers. Astropart. Phys., 15:177–202, 2001.

[40] Vidmar, R.J. On the use of atmospheric pressure plasmas as electromagnetic reflectors and ab-sorbers. IEEE. Transactions On Plasma Science, 18(4):733 – 741, 1990.

[41] Schneider, B.I and Brau, C. A. Two- and three-body electron attachment in air. Journal of PhysicsB: Atomic and Molecular Physics, 15(10):1601, 1982.

[42] Neunteufel, P. et al. Microwave emission due to molecular bremsstrahlung in non-thermal airshower plasmas. In 33rd International Cosmic Ray Conference, 2013.

[43] Dyrud, Lars P et. al. The anamolous diffusion of meteor trails. Geophysical Research Letters, 28:2775–2778, 2001.

[44] Blackett, P.M.S and Lovell, A.C.B. Radio echoes and cosmic ray showers. Proceedings of theRoyal Society A: Mathematical, Physical and Engineering Sciences, 177(969):183, 1941.

[45] Lovell, F.R.S. The blackett-eckersley-lovell correspondence of world war ii and the origin ofjodrell bank. Rec. R. Soc. Lond., 47(1):119–131, 1993.

[46] Crispin, J. W. and Maffett, A. L. Radar cross - section estimation for simple shapes. Proc. ofIEEE, 8:833–848, 1965.

[47] Bakunov, M.I. et al. Relativistic effects in radar detection of ionization fronts produced by ultra-high energy cosmic rays. Astropart. Phys., 33:335–3040, 2010.

[48] Stasielak, J. et al. Enhancement of the radar signal of air showers due to time compression.arXiv:1310.0743 [astro-ph.HE], 2013.

[49] Constantine A. Balanis. Antenna Theory Analysis and Design. Wiley - Interscience, 2005.

[50] W.C. Gibson. The Method of Moments in Electromagnetics. Chapman and Hall, 2008.

[51] J.D Kraus and R.J Marhefka. Antennas. McGraw - Hill, 2003.

[52] Hanson, J. The Performance and Initial Results of the ARIANNA Prototype. PhD thesis, Universityof California Irvine, 2013.

82

References

[53] G.A Dulk, et. al. Calibration of low-frequency radio telescopes using the galactic backgroundradiation. Astron. Atrophys., 365:294–300, 2001.

[54] H.V. Cane. Spectra of the non-thermal radio radiation from the galactic polar regions. Mon. Not.R. Astron. Soc., 189, 1979.

[55] Zaino, J.C. et al. An fpga based adaptive computing implementation of chirp. In Military andAerospace Applications of Programmable Devices and Technologies (MAPLD), 1999.

[56] Envelope detector. URL http://seniord.ee.iastate.edu/SSOL/RADAR/prjpln99/detector3.html.

[57] Analog Devices. AD80066, Oct 2014. URL http://www.analog.com/static/imported-files/data_sheets/AD80066.pdf.

[58] Analog Devices. AD9634, Oct 2014. URL http://www.analog.com/static/imported-files/data_sheets/AD9634.pdf.

[59] Xilinx. Spartan-6 FPGA SelectIO Resources, 1.6 edition, Feb 2014. URL http://www.xilinx.com/support/documentation/user_guides/ug381.pdf.

[60] Xilinx. Spartan-6 FPGA Block RAM Resources, 1.5 edition, Jul 2011. URL http://www.xilinx.com/support/documentation/user_guides/ug383.pdf.

[61] Xilinx. LogiCORE IP FIFO Generator, 9.2 edition, Jul 2012. URL http://www.xilinx.com/support/documentation/ip_documentation/fifo_generator/v9_2/pg057-fifo-generator.pdf.

[62] i Lotus. i-lotus M12M, Feb 2008. URL http://www.ilotus.com.sg/sites/all/themes/zeropoint/pdf/m12m/M12M%20Timing%20-%20TDS%20(Ver%201.0.0).pdf.

[63] Lantronix xpico wi-fi pi plate for raspberry pi, Oct 2014. URL http://www.lantronix.com/news/2014/10-08_lantronix-xpico-wi-fi-pi-plate-for-raspberry-pi-now-shipping-worldwide.html.

[64] de Vries, et al. On the feasibility of RADAR detection of high-energy induced showers in ice.Astropart. Phys., 60:25–31, 2015.

83

http://seniord.ee.iastate.edu/SSOL/RADAR/prjpln99/detector3.html

http://seniord.ee.iastate.edu/SSOL/RADAR/prjpln99/detector3.html

http://www.analog.com/static/imported-files/data_sheets/AD80066.pdf




http://www.xilinx.com/support/documentation/user_guides/ug381.pdf




http://www.xilinx.com/support/documentation/ip_documentation/fifo_generator/v9_2/pg057-fifo-generator.pdf



http://www.ilotus.com.sg/sites/all/themes/zeropoint/pdf/m12m/M12M%20Timing%20-%20TDS%20(Ver%201.0.0).pdf

http://www.ilotus.com.sg/sites/all/themes/zeropoint/pdf/m12m/M12M%20Timing%20-%20TDS%20(Ver%201.0.0).pdf

http://www.lantronix.com/news/2014/10-08_lantronix-xpico-wi-fi-pi-plate-for-raspberry-pi-now-shipping-worldwide.html



Appendix A

Schematics

CAM

Fig. A.1 shows the schematics of the power supply board for the CAM.

Triggering Board

Fig. A.2 shows the schematics of the Triggering board including the bandpass filters, envelope detectors

and the ADC.

Fig A.3 shows the connections between the FPGA, SBC and the GPS and Triggering Board.

High Speed Board

To provide superior common-mode noise rejection differential signalling is used and the 12-bit words

are transferred using a custom adapter (Fig. A.4) between the Evaluation Board’s LVDS (Low Voltage

Differential Signalling) parallel output port and the VHDC (Very High Density Cable) connector on the

FPGA.

Chirp Calibration Unit (CCU)

To clibrate the remote station a chirp calibration unit was deployed at the station location. Fig. A.5

shows the control circuit for the chirp generator.

84

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D5

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C6

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D6

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C7

C7

D7

D7

C8

C8

D8

D8

C9

C9

D9

D9

C10

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A2

A2

B2

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A3

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A4

A4

B4

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A5

A5

B5

B5

A6

A6

B6

B6

A7

A7

B7

B7

A8

A8

B8

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A9

A9

B9

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A10

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C2

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D2

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D7

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C10

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