Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://www.naic.edu/~astro/spectral_line/examples/HI_writeup_new.ps Äàòà èçìåíåíèÿ: Mon Jun 18 18:26:16 2001 Äàòà èíäåêñèðîâàíèÿ: Sat Dec 22 23:16:14 2007 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: ngc 7139

NAIC/NRAO Single Dish Summer School
Hands­On Project
HI Spectral­Line Observations of Galaxies
Karen O'Neil
For help with this project please contect Karen O'Neil or Jo Ann Eder

2 HI Spectral­Line Observations of Galaxies

K. O'Neil & J. Eder 3
Contents
1 Introduction 7
2 The Sources 7
3 Blank Sky Observation 8
4 System Temperature 8
5 Telescope Gain 9
6 Putting it all together -- Data Calibration 10
7 Finishing Up -- Converting your Data into Useful Numbers 11
8 Running the Telescope 13
Appendix I -- Source Lists 17

4 HI Spectral­Line Observations of Galaxies

K. O'Neil & J. Eder 5
CHECKLIST
Before coming to the control room to start your hands­on project, you should:
ffl Read this document and (at the least) skim through the appendices
ffl Have discussed, and come to a consensus, with your group as to what your source(s) and
calibrators will be
ffl Have discussed, and come to a consensus, with your group as to what your correlator set­up
will be
ffl Have come to an understanding with your group as to what analysis package(s) will be used
for your data reduction
ffl Have a reasonable idea of what will you need to do to complete your observations
ffl Know what time you are expected in the control room, and fully understand the
transportation schedule which will take you to and from Arecibo Observatory and El Buen
Cafe
ffl Be aware that there is limited food available in the control room at night, and come
prepared with your own food if you will be hungry
ffl Bring this write­up, your list of sources, something to write with, and paper to the control
room to take notes

6 HI Spectral­Line Observations of Galaxies

K. O'Neil & J. Eder 7
1. Introduction
This project is designed to demonstrate the fundamentals of spectral line observing with the
position switching technique. Position switching is the technique of observing first the position of
interest and then a blank sky (off­source) observation. The blank sky observation is then used to
remove atmospheric and systematic telescope noise. At Arecibo, position switching provides the
simplest means to eliminate standing waves in your baseline. This is good, as standing waves are
a large problem at Arecibo due to the unique geometry of the Gregorian system which allows for
considerable reflection from sources inside and outside the telescope beam.
2. The Sources
The choices for you project are up to you (and the members of your group) to decide. Appendix I
lists a variety of both HI emission and absorption sources from which to choose. The choice of
source, or sources, is up to you. There are, though, a few items to keep in mind when deciding
what to observe.
First, you want a source which you are able to detect within your alloted time (about 20 minutes
of on source observing). To decide if you will be able to detect a source in that amount of time,
you should be aware that five minutes of observation on a source gives you a 1oe r.m.s of 1.4 -- 1.9
mJy, after combining the two polarisations. (And then recall that (a) you want at least a 3oe
detection, and (b) doubling the time reduces the noise by 1
p
2.)
Second, you need to be aware of whether or not there is any RFI (Radio Frequency Interference)
at the frequency at which you wish to observe and, if using the L­Wide receiver, whether any
`trapped modes' exist at that frequency. The easiest means for finding out this information is
though the ``Spectrum Management'' web page at the Arecibo web site (http://www.naic.edu).
Alternatively, you can ask Karen O'Neil or Jo Ann Eder, who will be helping with your observing
project.
Third, you need to determine which sources will be in the sky accessible to the Arecibo dish (the
`Arecibo sky') during your observations. A reasonable estimate for the length of time an object
will be in the Arecibo sky can be obtained from the following: A source at Arecibo lattitude
(18.3 ffi declination) takes approximately 3 hours to transit the Arecibo sky, while a source 17 ffi
from the telescope latitude has only about 20--30 minutes while it transits the dish. Sources more
than ¸18.degree from the AO lattitude will never be visible to the AO system.
Finally, you should decide how you wish to observe your source(s). The primary objective of this
project is to obtain the HI spectra of at least one object. With this in mind you have the
following options: You can choose your set­up so that you have all four correlator boards centered
on the 21­cm line of your object, at a variety of resolutions; You can simply span the entire
velocity space over which your object should be, and observe in `search mode'; Or you can decide

8 HI Spectral­Line Observations of Galaxies
to look for more than one spectral line (i.e. HI and OH), and center two boards on each line.
(Note, though, that if you decide to look for both HI and OH you must use the L­wide receiver.)
3. Blank Sky Observation
The unique geometry of the Arecibo system allows for considerable reflection from sources both
inside and outside the beam. Additionally, the imperfect shape of the Arecibo dish and
irregularity of the surrounding countryside results in a telescope gain which varies across the dish.
As a result, the blank sky (off­source) observations taken in a typical position switching sequence
at Arecibo are not `randomly' chosen blank sky positions. Instead we force the off source
observation to trace the same pattern across the dish as the on source observation. As the
telescope is given one full minute between the on and off source observations (to allow the
telescope to slew to the correct position), a five minute on source observation will have its off
source observation taken at the same declination but at an RA position which is six minutes (in
time) later than that of the source. In this way the same azimuth and zenith angle positions are
traced for both the on and off source observations.
Potential Problems: The off source observation is supposed to be a blank sky observation.
This means you need to be sure the sky area which will be used for this observation does not
contain any significant continuum sources or any other galaxies. If continuum, or another galaxy,
is present at the location of your off source observation, you need to reconsider your observing
strategy (i.e. choose another source or choose a different integration time). A great resource for
determining the location on continuum near your object is the Northern VLA Sky Survey
(NVSS), available online through http://www.cv.nrao.edu/¸jcondon/nvss.html (For the purpose
of this project the postage stamp server is good tool.) To search an area of sky for other galaxies,
the NASA Extra­galactic Database (NED) at http://nedwww.ipac.caltech.edu has a ``near
position'' search that is very complete.
4. System Temperature
Receiver temperatures are obtained at Arecibo through observations of noise diodes with known
(pre­measured) temperatures. Typically, the telescope is moved to a blank sky position and a 10
second observation is done with the diode on, and then 10s with the diode off. To insure that
these measurements accurately reflect the system temperature at the time of the source
observation, the noise diode (or cal) measurements are usually done at the end of the on+off
source observations.
A complete listing of the available noise diodes is given below, and the values for the diode
temperatures for the past three years can be found in the control room. Additionally, the last
measured values are available on the web at http://www.naic.edu/¸operator/tsyslast.htm. The

K. O'Neil & J. Eder 9
typically measured values for an individual receiver can be found at
http://www.naic.edu/¸operator/normaltemp.htm.
Noise Diode Configuration Options:
nocal: run no calibration scans. This saves time but is rarely a good idea
lcal: low uncorrelated cal. diodes A,B straight through (D1!Pol A, D2!Pol B)
hcal: high uncorrelated cal. diodes A,B straight through (D1!Pol A, D2!Pol B)
lxcal: low uncorrelated cal. diodes A,B with crossover (D2!Pol A, D1!Pol B)
hxcal: high uncorrelated cal. diodes A,B with crossover (D2!Pol A, D1!Pol B)
l90cal: low correlated cal with 90 deg phase shift (D2+90 ffi !Pol A, D2!Pol B)
h90cal: high correlated cal with 90 deg phase shift (D2+90 ffi !Pol A, D2!Pol B)
lcorcal: low correlated cal with 0 deg phase shift (D1!Pol A, D1!Pol B)
hcorcal: high correlated cal with 0 deg phase shift (D1!Pol A, D1!Pol B)
Potential Problems: To achieve the best signal­to­noise, a high cal (high temperature diode)
should be used whenever possible. Additionally, careful consideration should be taken when
choosing the type of cal (correlated or uncorrelated, phase shifted or not, crossed­over or not).
Note that these considerations only apply, for our observations, if you are using the L­wide
receiver, as the L­narrow receiver has only one cal, the lcorcal, available.
5. Telescope Gain
Like all astronomical observations, one of the most important steps in spectral line observing is
converting your observations from local units (e.g. a local temperature) into universal units. For
HI observations a typical unit of measurement is the Jansky (Jy). Janskies are units of flux
density, with 1 Jy = 10 \Gamma23 ergs cm \Gamma2 s \Gamma1 Hz \Gamma1 . Conversion from measured source temperatures
(what you have, after taking T sys into account) into Janskies (what you want), is done through
observing a number of sources with well established fluxes, and then determining a conversion
between measured temperature and flux density. This conversion factor is known as the telescope
gain. (Technically, telescope gain is simply the telescope response to a point source. At Arecibo,
we express the gain in terms on antenna temperature due to a point source of flux density 1Jy
which lies in the center of the beam. With this definition, then, the gain is in units of K/Jy).
A good calibration source is a point source with a strong, non variable flux. For L band
observations (the regime in which HI observations are done), this means you want a source which
is considerably smaller than the 3.6 0 size of the beam. Additionally, the flux should be at least
20­30 mJy but not larger than a few Janskies, or you will exceed the dynamic range of the
telescope. A list of possible flux calibrators is given in Appendix I.
At Arecibo you use neither all of the dish, nor identical parts of the dish for every observation.
Additionally, the Arecibo dish (like most things in reality) is not perfect -- the dish, as used, is not
100% spherical, there are cables and supports to hold the Gregorian and receivers in place, etc.
As a result the gain at Arecibo has a dependence both on azimuth and zenith angle. The ideal

10 HI Spectral­Line Observations of Galaxies
calibrator at Arecibo would be a good calibrator (as defined above) which lies only a few beam
widths away in declination, and at the same RA, from the source in which you are interested. An
observer would then perform a number of observations on source in which they are interested and
then observe the calibrator source, repeating this sequence throughout the night.
Unfortunately, reality is rarely nice enough to grant every object in the Arecibo sky its own
calibrator source. (Or perhaps fortunately, since otherwise the sky would be too filled with strong
continuum sources to perform any decent spectral line observations.) To accommodate this fact,
and to save telescope time by avoiding each observer needing to repeat the same observations,
each observing band at Arecibo has a `gain curve' -- a two­dimensional fit of the telescope gain
versus azimuth and zenith angle. Ideally, this curve would be applied directly to the observed
data to obtain flux densities in Jy. Like with any experiment, though, blind trust in the absolute
values of a standard curve could readily land an observer into trouble. Instead it is usually a very
good idea to observe a few good calibration sources located near (in azimuth and declination) to
where your object will be located. These observations can then be used to confirm (or if necessary
correct) the standard calibration curves, and should be repeated periodically throughout your
observing run. (Copies of the curves can be found through the web pages of the respective
receivers -- http://www.naic.edu/menuimag/astronomy.htm.)
6. Putting it all together -- Data Calibration
Once all your observations have been taken, its time to sit down with your favorite data reduction
program and analyse your data. A description of the commands and methods for using the
various data reduction packages available at Arecibo is given in the next section. Rather than bog
you down with various commands, this section is designed to give an overview of the steps
necessary to turn your data into measurements of gas mass, velocity, etc. A more detailed
description of the calibration steps is included in `Notes Relevant to Spectral­Line Calibration'' by
C. Salter, available on the web.
The first step in your data reduction is to remove the background noise from your data. For
position switching, this is done using the off source observations taken during your observing run.
Removing background noise from your source involves subtracting out the contributions due to
the receiver, extraneous noise from the ground due to spillover, vignetting, etc, atmospheric noise,
and background radiation (e.g. the galactic background emission and the Cosmic Background
Radiation). Conveniently, all of this is readily done by simply subtracting your of source
observation from your on source observation.
At this point your data is simply in units of accumulated antenna temperature. To apply your
calibration corrections you need to convert it into independent units -- in this case, percentage of
system temperature. As you already have a blank sky observation taken for the same length of
time, this conversion is done by dividing your data by your off source observation.

K. O'Neil & J. Eder 11
The next step is to convert your data from %T sys into units of absolute temperature. To
accomplish this, you need to determine the system temperature using your observations of the
standard noise diodes (cals). Following the same logic as for your source, and knowing the
absolute temperature for the observed noise diode (TCAL ), the noise diodes can be used to
determine the system temperature via
T SYS = TCAL \Theta
Ÿ
CALON \Gamma CALOFF
CALOFF
–
The antenna temperature of your source is then
TSOURCE = T SYS \Theta
Ÿ
ON \Gamma OFF
OFF
–
= TCAL \Theta
Ÿ
CALON \Gamma CALOFF
CALOFF
–
\Theta
Ÿ
ON \Gamma OFF
OFF
–
As a last step you should convert your measured source temperature into units of Janskies by
applying the relevant gain curve. Remembering that the gain is in units of K/Jy, the final
equation is then
Signal(Jy=beam) = TCAL
G(za; az) \Theta
Ÿ
CALON \Gamma CALOFF
CALOFF
–
\Theta
Ÿ
ON \Gamma OFF
OFF
–
7. Finishing Up -- Converting your Data into Useful Numbers
Once you've successfully calibrated your data, you need to turn your spectra into useful
(publishable) facts -- HI mass, velocity widths, etc. Fortunately, this step is extremely easy. Using
your favorite data reduction software, you need to (a) sum all the relevant data (if you haven't
already); (b) perform a baseline subtraction on your data; (c) integrate the total flux of the
galaxy (to obtain the flux and, ultimately, the total HI gas mass), and (d) fit Gaussian(s) to your
spectra to obtain the central velocity and the FWHM at 50% and 20% of the Gaussians (that is,
obtain the velocity widths).
Baseline Subtraction: The region chosen for your baseline subtraction should be large enough
that sufficient channels are available to get a reasonable fit, yet small enough that the baseline
you are fitting will be accurate for the spectral line in which you are interested. Typically, at
Arecibo, this means no more than a few hundred channels on either side of the spectral line, and
sometimes much less (depending, of course, on the width of the observed line). The other factor
to consider when fitting a baseline is that you do not want to over­ or under­compensate for the
true baseline. The interim correlator at Arecibo (the one which you are using) is extremely good,
for many reasons, including the fact that it gives remarkably flat baselines. As a result, assuming
you choose a baseline which is limited to the region reasonably near your observed line, the
baseline which you fit to your data should not be of order larger than two.
Determining the HI mass: First, as already mentioned, you need to use you favorite data
reduction program to integrate and obtain the total flux of your source. When performing this

12 HI Spectral­Line Observations of Galaxies
integration, you want to be sure that your y­axis (or intensity) is in units of Janskies, and your
x­axis is in terms of velocity. The flux you obtain will then be in the very useful units of
Jy km s \Gamma1 . This can be converted to, i.e.total gas mass (in units of M fi ) through the conversion:
M(HI) = 2:356 \Theta 10 5 \Theta
h
D(Mpc) 2
i
\Theta
Z
(Jy km s \Gamma1 ) M fi
where R
is the measured flux of your source and D is the distance to your object in Mpc (for
distant objects -- objects where v sys is considerable larger than v random -- a reasonable distance
estimate is D=v(km s \Gamma1 )/H 0 (km s \Gamma1 Mpc \Gamma1 ) ).
Determining v sys , v 50 and v 20 : Both the systematic velocity and the velocity widths of HI
absorption and emission lines are typically obtained by fitting one or two Gaussians to the
absorption/emission spectra (the number of Gaussians is chosen base on whether the spectra
shows a one or two `horned' profile). The systematic velocity is then the velocity at the peak of
the Gaussian (or the center of the peaks of the two Gaussians). V 20 and v 50 are the velocity
width at 20% and 50% of the Gaussian peak(s).
When no line is seen (setting upper limits): If no line is detected during your observation,
it is still useful to obtain an upper limit to the flux density of your object.
In order to set flux limits you first need to decide on what a likely velocity width (v 20 ) for your
object is. The easiest means to do this is to go to the literature and determine common velocity
widths for similar objects. When choosing the velocity width to assume for your (undetected)
source, remember that you are setting upper limits on the flux and mass. This means that you
should not choose a small velocity width simply because it gives an impressively low flux limit.
Instead, if anything, you should err on the side of overestimating your object's velocity width.
Once you've decided on a reasonable velocity width, you can determine your flux, and mass,
upper limits. There are two common methods for setting these limits. In the first method, you
need only to determine the r.m.s. noise of the region in your spectra which should contain the line
of interest. Your flux upper limit would then be ! # oe ? \Thetaoe rms \Theta v 20 , where ! # oe ? represents
the signal­to­noise level, in number of oe rms , at which you would expect to see your spectral line.
(This is typically of order 3 or higher.) You can now convert that flux upper limit to a mass limit
via the conversion equation given earlier.
The other common method for determining upper flux limits for non­detections at Arecibo is by
integrating the noise over the region where your source should lie. This gives you a 1oe upper limit
to you flux. If there is any non­linear components to your baseline, this method will provide a
more realistic (albeit higher) flux limit.
One final note in this discussion is the possibility that you do not know the exact red shift of your
source. If you know an approximate velocity, or a range, you can apply either of the above
techniques repeatedly across the velocity range of interest, and then use the highest value as your
flux upper limit. If, though, you have no idea what the red shift of your source is, you are
probably best off just reporting your r.m.s. values.

K. O'Neil & J. Eder 13
8. Running the Telescope
Detailed information on available telescope parameters, using the telescope control gui, etc. are
given in the Spectral Line User's Manual, an abbreviated version of which is available on the web.

14 HI Spectral­Line Observations of Galaxies

K. O'Neil & J. Eder 15
APPENDICES

16 HI Spectral­Line Observations of Galaxies

K. O'Neil & J. Eder 17
Appendix I
Source Lists
Table 1: Possible HI Emission Sources
Object Name RA Dec v rad w 20 Flux S/N
(B1950) (B1950) (km/s) (km/s) (Jy km/s)
UGC 2602y 03 11 21 +16 18 03 10098\Sigma 8 305\Sigma 25 5.1\Sigma .8 7.7
UGC 2716y 03 21 11 +17 34 00 379\Sigma 2 67\Sigma 7 4.8\Sigma .6 19.9
UGC 2840y 03 41 00 +22 29 00 3650\Sigma 9 174\Sigma 28 3.3\Sigma .7 7.1
UGC 2862y 03 44 51 +13 05 56 6703\Sigma 18 466\Sigma 55 2.6\Sigma .7 4.2
UGC 2904y 03 54 12 +16 20 00 8014\Sigma 4 312\Sigma 13 5.0\Sigma .5 12.2
UGC 2927y 03 58 41 +22 58 00 6251\Sigma 8 251\Sigma 24 5.6\Sigma .8 8.4
UGC 3094 a 04 32 38 +19 04 07 7408\Sigma 7 609\Sigma 21 7.4\Sigma 1.4 4.7
UGC 3135y 04 39 11 +19 57 00 7399\Sigma 8 436\Sigma 25 6.7\Sigma .9 8.0
UGC 3157y 04 43 39 +18 22 13 4619\Sigma 8 314\Sigma 25 3.3\Sigma .7 5.5
UGC 3312y 05 25 36 +22 04 00 5654\Sigma 32 428\Sigma 96 2.4\Sigma .9 2.5
UGC 3356y 05 44 15 +17 32 42 5613\Sigma 11 197\Sigma 34 3.3\Sigma .7 7.7
UGC 3498y 06 32 29 +15 01 13 3800\Sigma 6 356\Sigma 18 5.4\Sigma 1.1 7.6
UGC 3503y 06 35 00 +22 40 59 1389\Sigma 4 142\Sigma 12 7.0\Sigma .8 16.8
UGC 3524y 06 41 12 +12 27 15 3930\Sigma 5 431\Sigma 16 6.6\Sigma .7 10.5
UGC 3534y 06 42 36 +22 28 00 4465\Sigma 3 239\Sigma 10 5.8\Sigma .5 17.3
MCG +4­16­ 3y 06 46 54 +25 41 00 4835\Sigma 8 592\Sigma 23 14.4\Sigma 1.6 8.7
UGC 3702y 07 06 00 +28 47 00 4980\Sigma 5 88\Sigma 16 2.9\Sigma .5 11.6
UGC 3726y 07 08 02 +25 59 57 7647\Sigma 12 511\Sigma 35 3.4\Sigma .8 4.0
UGC 3736y 07 09 19 +23 27 55 9886\Sigma 12 505\Sigma 37 5.2\Sigma .9 4.6
UGC 3784y 07 14 54 +26 44 00 2587\Sigma 5 205\Sigma 15 4.0\Sigma .9 6.5
MCG +5­18­11y 07 21 55 +27 25 27 7784\Sigma 14 248\Sigma 42 .9\Sigma .5 2.4
UGC 11672z 21 02 07.9 +15 53 33 9449 465 1.10\Sigma .1 6.2
UGC 11676z 21 03 31.5 +18 16 00 4891 247 5.28\Sigma .1 19.9
UGC 11677z 21 03 44.3 +15 46 13 4964 349 10.40 18.3
UGC 11681z 21 05 27.6 +16 08 03 4968 605 2.00 12.2
UGC 11683z 21 06 08.7 +17 59 20 5029 326 2.59 12.7
UGC 11749z 21 26 30.0 +20 17 33 5312 239 7.14 26.6
UGC 11768z 21 32 13.9 +18 43 40 6608 439 1.45 11.0
UGC 11774z 21 33 49.2 +16 50 10 6873 228 3.92 12.5
UGC 11811z 21 45 06.1 +18 30 06 6656 253 2.94 8.3
UGC 11821z 21 47 14.3 +15 33 20 7773 624 2.11 11.3
IC 5177y 22 09 06 +11 33 00 8103\Sigma 12 437\Sigma 35 4.0\Sigma .7 5.6
UGC 11964y 22 13 06 +18 58 00 1420\Sigma 5 153\Sigma 14 4.0\Sigma .9 6.3
NGC 7321y 22 34 03 +21 21 38 7139\Sigma 6 446\Sigma 18 5.7\Sigma .8 7.8
UGC 12250y 22 53 06 +12 31 25 7254\Sigma 18 425\Sigma 53 1.9\Sigma .6 3.7
UGC 12303y 22 58 21 +26 28 16 7962\Sigma 5 274\Sigma 14 3.1\Sigma .6 7.0
UGC 12308y 22 58 48 +14 04 15 2225\Sigma 2 236\Sigma 5 16.8\Sigma 1.0 23.5
yProfiles are in Theureau, et.al (1998) A&AS 130, p. 1.
zProfiles are in Giovanelli & Haynes (1993) AJ 105, p. 1271
a Profile in Bottinelli, et.al (1992) A&AS 102, p. 57.

18 HI Spectral­Line Observations of Galaxies
Table 2: Possible HI Absorption Sources
Object Name RA Dec z abs w 20 NHI
(J2000) (J2000) (km/s) (km/s) (\Theta10 20 (cm \Gamma2
B0738+313A 07 41 10.7033 +31 12 00.229 0.09123 25 15.
B0738+313B 07 41 10.7033 +31 12 00.229 0.2212 16 7.9
PKS 1413+135 14 15 58.8174 +13 20 23.713 0.24671 39 200
3C196 08 13 36.033 +48 13 02.56 0.4366 250 6.3
3C286 13 31 08.2881 +30 30 32.960 0.69215 18 20.
PKS 0454+039 04 56 47.1749 +04 00 52.950 0.8596 -- 5.0
MC3 1331+170 13 33 35.7838 +16 49 04.033 1.77636 25 76.
PKS 1157+014 11 59 44.817 +01 12 06.18 1.944 60 63.
J2342+342 23 44 51.254 +34 33 48.64 2.9084 -- 20.
PKS 0201+113 02 03 46.6570 +11 34 45.410 3.3875 -- 25.
Reference: Chengalur & Kanekar (2000) MNRAS 318, p. 303

K. O'Neil & J. Eder 19
Table 3: List of Potential Flux Calibrators
Name RA Dec Flux (1420) Size Variable?
(B1950) (B1950) (Jy)
B0300+162 03 00 26.9 16 15 04 2.73 !10 00 N
B0316+162 03 16 09.1 16 17 41 7.79 !10 00 N
B0333+321 03 33 22.4 32 08 36 3.35 !10 00 Y
B0428+205 04 28 06.86 20 31 09.2 3.86 !10 00 N
B0453+227 04 53 42.17 22 44 41.7 3.21 !10 00 N
B0500+019 05 00 45.18 01 58 54.7 2.30 !10 00 N
B0518+165 05 18 16.53 16 35 26.8 8.09 !10 00 N
B0521+281 05 21 14.4 28 09 50 1.77 ¸2.1 0 N
B0640+233 06 40 04.9 23 22 10 2.64 !10 00 N
B0711+356 07 11 05.62 35 39 52.5 1.83 !10 00 Y
B2041+170 20 41 59.4 17 00 00 0.66 !10 00 N
B2153+377 21 53 45.52 37 46 13.9 6.81 !10 00 N
B2209+080 22 09 32.2 08 04 25 1.69 !10 00 N
B2210+016 22 10 05.17 01 37 59.4 2.84 !10 00 N
B2223+210 22 23 14.8 21 02 50 1.97 !10 00 Y
B2247+140 22 47 56.33 14 03 56.7 2.13 !10 00 N
B2249+185 22 49 07.7 18 32 44 1.72 !10 00 N
B2251+158 22 51 29.5 15 52 54 13.2 !10 00 Y
B2314+038 23 14 02.27 03 48 55.2 4.51 !10 00 N
B2338+132 23 38 00.9 13 16 23 1.78 !10 00 N