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Using the period search method described. contact uuaries will eeneralle eive an ideutified perio that corresponds to the orbital period of the svsteun. | Using the period search method described, contact binaries will generally give an identified period that corresponds to the orbital period of the system. |
Hence. the period range chosen above covers a range from about 250 nuuutes (0.175 d) to 333 nüuutes (0.23 d) for eclipsing binary orbital periods. | Hence, the period range chosen above covers a range from about 250 minutes (0.175 d) to 333 minutes (0.23 d) for eclipsing binary orbital periods. |
This is therefore the period range of interest. spannius from just lonecr thu the suggested period cut-off (0.22 d) to below the period of the shortest known svstem (0.19 d). | This is therefore the period range of interest, spanning from just longer than the suggested period cut-off (0.22 d) to below the period of the shortest known system (0.19 d). |
Each of the 5.600 light. curves with periods outside the three ‘contaminated’ period regions were examined by eve to pick out candidate eclipsing binaries. | Each of the 5,600 light curves with periods outside the three `contaminated' period regions were examined by eye to pick out candidate eclipsing binaries. |
Most of these helt curves displaved broadly siuusoidal modulation so are likely to represent some fori of rotational variability. | Most of these light curves displayed broadly sinusoidal modulation so are likely to represent some form of rotational variability. |
AMauy others showed asvinmetric ight curves (with narrow niaxinia and broad minima) characteristic of pulsating stars. | Many others showed asymmetric light curves (with narrow maxima and broad minima) characteristic of pulsating stars. |
These objects will be reported elsewhere. | These objects will be reported elsewhere. |
LO unique objects. reported here. displaved light curves that are characteristic of W UMGa stars. namely xoad maxima with narrow muna (see Figure 2. Figure 3 aud Table 1). | 40 unique objects, reported here, displayed light curves that are characteristic of W UMa stars, namely broad maxima with narrow minima (see Figure 2, Figure 3 and Table 1). |
The large umber of apparently periodic light. curves close to periods of 1/9 d. 1/10 d and 1/11 d were not all eve-balled. as the vast majority of these are expected to be spurious signals due to systematic noise. | The large number of apparently periodic light curves close to periods of 1/9 d, 1/10 d and 1/11 d were not all eye-balled, as the vast majority of these are expected to be spurious signals due to systematic noise. |
However. the period distribution iu Figure 1 shows hat ~1 of these are likely to be genuine. so the ~1 of thei showing the strongest signals exanuned iudividually. | However, the period distribution in Figure 1 shows that $\sim 1\%$ of these are likely to be genuine, so the $\sim 1\%$ of them showing the strongest signals examined individually. |
The strongest signals were chosen because the signal we are looking for is geucrally of lugher amplitude than our typical systematic noise signal. | The strongest signals were chosen because the signal we are looking for is generally of higher amplitude than our typical systematic noise signal. |
This allowed a further 13 caudidate eclipsing svstenis to be identified in he set of objects with periods close to 1/9 d (sce Figure 2 ix Table 1). | This allowed a further 13 candidate eclipsing systems to be identified in the set of objects with periods close to 1/9 d (see Figure 2 and Table 1). |
This exercise also revealed the known sdB|dM eclipsing binary. NY Vir (sce Figure 1). with a period of 115.163 minutes, within the set of objects with periods close to 1/10 d. No candidate W UMa stars were found in this latter set or in the set of objects close to periods of l/ll1 d. but the clear detection of NY Vir shows that such binaries would have been found. had thev existed. despite the systematic noise close to these periods. | This exercise also revealed the known sdB+dM eclipsing binary, NY Vir (see Figure 4), with a period of 145.463 minutes, within the set of objects with periods close to 1/10 d. No candidate W UMa stars were found in this latter set or in the set of objects close to periods of 1/11 d, but the clear detection of NY Vir shows that such binaries would have been found, had they existed, despite the systematic noise close to these periods. |
Figure 5 shows the orbital period distribution of the 53 candidate eclipsing binaries found aud listed in Table 1. | Figure 5 shows the orbital period distribution of the 53 candidate eclipsing binaries found and listed in Table 1. |
Further evidence for the fact that these candidates are all W UMa stars is provided by their V.AK colours which are each siguificautlv redder than the majority of the rest of the objects detected in this period rauge. as shown iu Figure 6. | Further evidence for the fact that these candidates are all W UMa stars is provided by their $V-K$ colours which are each significantly redder than the majority of the rest of the objects detected in this period range, as shown in Figure 6. |
Iu fact. thei Iv colours of ~1.5.3.5 are typical of stars of spectral type Ix. The 53 candidate eclipsing binaries include the known W UMa stars: VLO67 Πα. VI1OL Πο. ROTSEL J1702[0.11|151122.7 aud ASAS J1119323950.8. as well as the shortest period object previously known. GSC23110530. described carlier (Dimitrov Kjurkchieva 2010). | In fact, their $V-K$ colours of $\sim 1.5 - 3.5$ are typical of stars of spectral type K. The 53 candidate eclipsing binaries include the known W UMa stars: V1067 Her, V1104 Her, ROTSE1 J170240.11+151122.7 and ASAS J111932–3950.8, as well as the shortest period object previously known, $2314-0530$, described earlier (Dimitrov Kjurkchieva 2010). |
The SuperWASP light curve of this latter object is shown separately in Figure 3. folded at its confined period of 0.1926 d. Whilst we agree with the periods previously reported for ASAS J111932.3950.8 (0.2201 d. Pavihar et al 2009) aud V11041 Ier (0.2279 α. Akerlof et al 2000). we find better periods for V1067 Ier and ΠΟΤΟ J170210.111151122.7. | The SuperWASP light curve of this latter object is shown separately in Figure 3, folded at its confirmed period of 0.1926 d. Whilst we agree with the periods previously reported for ASAS J111932–3950.8 (0.2294 d, Parihar et al 2009) and V1104 Her (0.2279 d, Akerlof et al 2000), we find better periods for V1067 Her and ROTSE1 J170240.11+151122.7. |
The ROTSE survey had previously suggested periods of 0.2581 d and 1.210 d respectively for these two objects (Akerlof et al 2000). but we show them to have periods of 0.2285 d and 0.2312 d respectively. | The ROTSE survey had previously suggested periods of 0.2581 d and 1.210 d respectively for these two objects (Akerlof et al 2000), but we show them to have periods of 0.2285 d and 0.2312 d respectively. |
For a muuber of vears the shortest period main sequence eclipsing binary known was the faint (V. — 18.3) object OGLE BW03 Vo3s. with an orbital period of 0.1981 d | For a number of years the shortest period main sequence eclipsing binary known was the faint (V = 18.3) object OGLE BW03 V038, with an orbital period of 0.1984 d |
where b;=7/2 0r 3/2 when planet j is an internal or external perturber respectively. | where $b_j=7/2$ or $-3/2$ when planet $j$ is an internal or external perturber respectively. |
The stirring time {ώμος for the Sun-Jupiter-Saturn system is plotted in Figure (135. | The stirring time $t_{\mathrm{cross}}$ for the Sun-Jupiter-Saturn system is plotted in Figure \ref{fig:tcross-jupsat}) ). |
We also show /.4,, with Saturn's mass reduced to that of Earth. and for the single-planet ease with Jupiter alone perturbing the dise. | We also show $t_{\mathrm{cross}}$ with Saturn's mass reduced to that of Earth, and for the single-planet case with Jupiter alone perturbing the disc. |
With Saturn at its true mass. we see that the crossing timescale is greatly reduced close to the planets. | With Saturn at its true mass, we see that the crossing timescale is greatly reduced close to the planets. |
However. beyond the outermostsecular resonance. the dependence of foros. on e Steepens. | However, beyond the outermostsecular resonance, the dependence of $t_{\mathrm{cross}}$ on $a$ steepens. |
Specitically. for large e. foros.xa> rather than ot”. | Specifically, for large $a$, $t_{\mathrm{cross}}\propto a^8$ rather than $a^{4.5}$. |
This is because we now have [4|«|g;| in the forced eccentricity term in Equation (32)). and so the @ dependence of eLand D; no longer cancels. | This is because we now have $|A| < |g_i|$ in the forced eccentricity term in Equation \ref{eq:tcrossmulti}) ), and so the $a$ dependence of $A$ and $B_j$ no longer cancels. |
For planetesimals beyond AAU. introducing another perturber has the time-scale for orbit crossing. | For planetesimals beyond AU, introducing another perturber has the time-scale for orbit crossing. |
When Saturn's mass is reduced to that of Earth. we see a large region between Saturn and the outer secular resonance where the the time-scale is the same as for the case with Jupiter alone: because of the large disparity in masses. the perturbations are dominated by Jupiter. | When Saturn's mass is reduced to that of Earth, we see a large region between Saturn and the outer secular resonance where the the time-scale is the same as for the case with Jupiter alone: because of the large disparity in masses, the perturbations are dominated by Jupiter. |
As an example of a multi-planet system with a debris dise. consider HD 38529. | As an example of a multi-planet system with a debris disc, consider HD 38529. |
This star hosts a 0.8 Jupiter mass planet on a AAU orbit and a 12.2 Jupiter mass planet on a AAUorbit. | This star hosts a 0.8 Jupiter mass planet on a AU orbit and a 12.2 Jupiter mass planet on a AUorbit. |
The secular dynamics of the system. including both planets and massless planetesimals. were modelled by ?.. who concluded rom the dynamical analysis and SED fitting that the planetesimals reside in a dynamically stable region at AAU between secular resonances. | The secular dynamics of the system, including both planets and massless planetesimals, were modelled by \cite{2007ApJ...668.1165M}, who concluded from the dynamical analysis and SED fitting that the planetesimals reside in a dynamically stable region at AU between secular resonances. |
Figure (143) shows the crossing time-scale or HD 38529, | Figure \ref{fig:hd38529}) ) shows the crossing time-scale for HD 38529. |
Within the region AAU. the crossing scale is close to that achieved by planet ¢ alone. due to its higher mass and larger semi-major axis. | Within the region AU, the crossing time-scale is close to that achieved by planet c alone, due to its higher mass and larger semi-major axis. |
Within this region of the disc. he planets” secular perturbations induce crossing of neighbouring initially circular orbits on time-scales of 1 MMyr. | Within this region of the disc, the planets' secular perturbations induce crossing of neighbouring initially circular orbits on time-scales of $\lesssim 1$ Myr. |
It may well be he case then that there are no bodies larger than several kilometers in radius in the disc (but see refs:ic)). | It may well be the case then that there are no bodies larger than several kilometers in radius in the disc (but see \\ref{s:ic}) ). |
Note that /.4,,, is formally infinite when ατα=0. | Note that $t_{\mathrm{cross}}$ is formally infinite when $\mathrm{d}A/\mathrm{d}a=0$. |
This would appear to mean that there is a particular semi-major axis between the planets where the perturbations can never induce orbit-crossing. | This would appear to mean that there is a particular semi-major axis between the planets where the perturbations can never induce orbit-crossing. |
However. this singularity is merely a mathematical artefact (see refs:timescale) in reality. this region can still be stirred by secular perturbations. given sufficient time. | However, this singularity is merely a mathematical artefact (see \\ref{s:timescale}) ): in reality, this region can still be stirred by secular perturbations, given sufficient time. |
We have presented a simple picture of the effect of secular perturbations on a planetesimal disc. | We have presented a simple picture of the effect of secular perturbations on a planetesimal disc. |
Here we clarify the assumptions and limitations attached to our model. | Here we clarify the assumptions and limitations attached to our model. |
To determine the outcome of collisions. we have chosen one particular scaling law for threshold collision energy. | To determine the outcome of collisions, we have chosen one particular scaling law for threshold collision energy. |
Other scaling laws differ in the planetesimal radius at which the minimum of (Qj is attained. and the value of the minimum itself. | Other scaling laws differ in the planetesimal radius at which the minimum of $Q_{\mathrm{D}}^*$ is attained, and the value of the minimum itself. |
Given the very large values of « for mm bodies. this is unlikely to be important for this size of planetesimal. but may be important if the planetesimals are larger. | Given the very large values of $a^*$ for m bodies, this is unlikely to be important for this size of planetesimal, but may be important if the planetesimals are larger. |
We have assumed that collisions occur as soon as the orbits begin to cross. and accounting for collision rates will slightly increase the stirring time. | We have assumed that collisions occur as soon as the orbits begin to cross, and accounting for collision rates will slightly increase the stirring time. |
In common with other studies (e.g.. 22). we have taken the initial planetesimal orbits to be circular. | In common with other studies \citep[e.g.,][]{2006Icar..183..193T,2008ApJS..179..451K}, , we have taken the initial planetesimal orbits to be circular. |
Despite promising recent progress (e.g. 22). the formation of planetesimals is still not fully | Despite promising recent progress \citep[e.g.,][]{2007Natur.448.1022J,2008ApJ...687.1432C}, , the formation of planetesimals is still not fully |
spectrum emission such as synchrotron radiation will be negligible at 31 GHz. | spectrum emission such as synchrotron radiation will be negligible at 31 GHz. |
No emission is detected within the LDN1621 area. | No emission is detected within the LDN1621 area. |
After smoothing to the CBI resolution, the r.m.s. | After smoothing to the CBI resolution, the r.m.s. |
noise level at 4.85 GHz is z3 mJy beam™! corresponding to a 30 upper limit of 9 mJy beam@!. | noise level at 4.85 GHz is $\approx 3$ mJy $^{-1}$ corresponding to a $3\sigma$ upper limit of 9 mJy $^{-1}$. |
At frequencies =1 GHz, free-free emission will be optically thin,: except {or the densest clouds, such as ultracompact HII regions. | At frequencies $\ga1$ GHz, free-free emission will be optically thin, except for the densest clouds, such as ultracompact HII regions. |
We can therefore reliably extrapolate the upper limit to 31 GHz assuming a flux density spectral index (S« v?) of a=—0.12, appropriate for optically thin free-free emission (Dickinsonetal.2003). | We can therefore reliably extrapolate the upper limit to 31 GHz assuming a flux density spectral index $S \propto \nu^{\alpha}$ ) of $\alpha=-0.12$, appropriate for optically thin free-free emission \citep{Dickinson03}. |
. The 3c upper limit to free-free emission at 31 GHz is then 7.2 mJy beam" | The $3\sigma$ upper limit to free-free emission at 31 GHz is then 7.2 mJy $^{-1}$. |
We can therefore be quite confident that the bulk of the 31 !.GHz emission cannot be due to synchrotron or free-free emission, although a small (up to a maximum of ~30 per cent) contribution⋅ from free-free contribution⋅ cannot be ruled out withith the d.data available. | We can therefore be quite confident that the bulk of the 31 GHz emission cannot be due to synchrotron or free-free emission, although a small (up to a maximum of $\sim 30$ per cent) contribution from free-free contribution cannot be ruled out with the data available. |
ilabl The level of free-free emission can also be estimated | The level of free-free emission can also be estimated |
On the map. the angle rotation aloug the geodesic with au azimuth.0. is calculated by coustructing thin spherical triangle with side-angle-side given by (7/2—©.0.d7). | On the map, the angle rotation along the geodesic with an azimuth,$\theta$, is calculated by constructing thin spherical triangle with side-angle-side given by $(\pi/2 - \phi, \theta,d\tau)$. |
because the geodesic intersects the north-soutl meridian (a vertical line in the Mercator) at au augle of 0 initially. auc at an augle x—9 at the other eud. where 9 is the auele iu the spherical triangle at the other eud of the dr side. | because the geodesic intersects the north-south meridian (a vertical line in the Mercator) at an angle of $\theta$ initially, and at an angle $\pi-\beta$ at the other end, where $\beta$ is the angle in the spherical triangle at the other end of the $d\tau$ side. |
Solving for da using spherical trigonometry in the limit as dT goes to zero. we find that Thus. lor 0=7/2. au east-west geodesic. we find that so that in the northern hemisphere. traveling east one's geodesic is bendiug clockwise witli dafdr=lanó. | Solving for $d\alpha$ using spherical trigonometry in the limit as $d\tau$ goes to zero, we find that Thus, for $\theta = \pi/2$, an east-west geodesic, we find that so that in the northern hemisphere, traveling east one's geodesic is bending clockwise with $d\alpha/d\tau = tan \phi$. |
Therelore. the east-west geodesic bends downward. | Therefore, the east-west geodesic bends downward. |
For. 0.= a north-south eeodesic. the [lexion ix zero. as we expect. since the merkliaus of lougitude are straight in the Mercator map. | For, $\theta = 0$, a north-south geodesic, the flexion is zero, as we expect, since the meridians of longitude are straight in the Mercator map. |
IL we average over all azimuths at a given poiut. we fin: Now we cau integrate this over all points ou the sphere to produce the average flexion over the whole sphere F. Taking advantage of the sviumetry between the northern aud southern hemisphere we cal integrate ouly over the northern hemisphere (where (4=23cos odo). yielding: The flexiou is less than that of the stereographic because of the 1807 bouudary cut along the longitude line at the international date line. | If we average over all azimuths at a given point, we find: Now we can integrate this over all points on the sphere to produce the average flexion over the whole sphere F. Taking advantage of the symmetry between the northern and southern hemisphere we can integrate only over the northern hemisphere (where $dA = 2\pi \cos\phi d\phi$ ), yielding: The flexion is less than that of the stereographic because of the $180^\circ$ boundary cut along the longitude line at the international date line. |
A geodesic is a great circle on the globe. aud if this is shown as a closed curve ou the map that is always coucave luward (he best possible case) it will always have a total rotation of the velocity vector of 360° aud so will have au average integrated flexion of 1. | A geodesic is a great circle on the globe, and if this is shown as a closed curve on the map that is always concave inward (the best possible case) it will always have a total rotation of the velocity vector of $360^\circ$ and so will have an average integrated flexion of 1. |
I£ there is a bouudary cut. the great circle does not have to close ou the map (it has two loose ends at the boundary eut) aud so need not completely rotate by 3607. | If there is a boundary cut, the great circle does not have to close on the map (it has two loose ends at the boundary cut) and so need not completely rotate by $360^\circ$. |
Acceleration in the direction. parallel to the velocity vector of the truck «|. causes the truck to increase its speed aloug the geodesic curve without causing any rotatiou. | Acceleration in the direction parallel to the velocity vector of the truck $a_\parallel$, causes the truck to increase its speed along the geodesic curve without causing any rotation. |
This causes skewuess. because as the truck accelerates. it covers more distauce on the map on oue side of a point than on the other. so the point in question will uot be at the center of the liue segment of are centered ou that point ou the spliere. | This causes skewness, because as the truck accelerates, it covers more distance on the map on one side of a point than on the other, so the point in question will not be at the center of the line segment of arc centered on that point on the sphere. |
Conskler a segment of a meridian of longitude ou the globe centered at 15° north. latitude. | Consider a segment of a meridian of longitude on the globe centered at $45^\circ$ north latitude. |
Going from South to North aloug that geodesic in the Mercator map the truck is accelerating with a)> 0. because the scale factor is getting larger and larger the further north one goes. so as the | Going from South to North along that geodesic in the Mercator map the truck is accelerating with $a_\parallel > 0$ because the scale factor is getting larger and larger the further north one goes, so as the |
Wide-Angle Tail (WAT) galaxies are radio. galaxies whose radio jets appear to bend in a common direction. | Wide-Angle Tail (WAT) galaxies are radio galaxies whose radio jets appear to bend in a common direction. |
They are generally detected in dynamical. non-relaxed clusters of galaxies (e.g.Burns1990). and may be used as probes or tracers for clusters (Blantonotal.2000. 2001). | They are generally detected in dynamical, non-relaxed clusters of galaxies \citep[e.g.][]{Burns90}
and may be used as probes or tracers for clusters \citep{Blanton00,
Blanton01}. |
.. Clusters of galaxies are the largest eravitationally bound. structures in the Universe and are powerful. testhecs of cosmological mocels (c.g.al. 2009). | Clusters of galaxies are the largest gravitationally bound structures in the Universe and are powerful testbeds of cosmological models \citep[e.g.][] {Borgani04, Sahlen09,
Kravtsov09}. |
. Clusters also host diffuse racio emission in the form of radio haloes and relies (Ciovannini&boer-otal. 2009). | Clusters also host diffuse radio emission in the form of radio haloes and relics \citep{Giovannini00, Feretti05, Ferrari08,
Giovannini09}. |
. The bent nature of WATS has commonly been attributed. to strong intra-cluster winds caused by dynamical interactions such as cluster-cluster mergers (Burns1998). | The bent nature of WATs has commonly been attributed to strong intra-cluster winds caused by dynamical interactions such as cluster-cluster mergers \citep{Burns98}. |
. WATs are prelerentially found. in enhanced X-ray regions (Pinkneyetal.2000)/ and are usually associated with dominant cluster galaxies (Owen&Rudnick1976). | WATs are preferentially found in enhanced X-ray regions \citep{Pinkney00} and are usually associated with dominant cluster galaxies \citep{Owen76}. |
. Maoetal.(2009a) found he tailecl radio galaxies. including WATs. to be ocated in the densest regions of clusters in the ocal Universe. consistent with earlier. studies: (e.g.Burns1990:Blantonetal.2000. 2001). | \citet{Mao09a} found the tailed radio galaxies, including WATs, to be located in the densest regions of clusters in the local Universe, consistent with earlier studies \citep[e.g.][] {Burns90, Blanton00, Blanton01}. |
. Phus WAYS represent valuable tracers of high density regions in he intracluster medium (CAL). and this approach has oen used in a number of recent studies (e.g.BlantonOklopcicetal. 2010). | Thus WATs represent valuable tracers of high density regions in the intracluster medium (ICM), and this approach has been used in a number of recent studies \citep[e.g.][] {Blanton00, Blanton03, Smolcic07, Giacintucci09,
Kantharia09, Oklopcic10}. |
. Llere we present the radio properties of six. WATs that we have identified in ATLAS. the Australia | Here we present the radio properties of six WATs that we have identified in ATLAS, the Australia |
by subtracting an average bias [rame and dividing by a normalized Hat field. | by subtracting an average bias frame and dividing by a normalized flat field. |
Phe near-LR. data was reduced using and the specific PANIC package provided by the Las Campanas Observatory. | The near-IR data was reduced using and the specific PANIC package provided by the Las Campanas Observatory. |
The lis a 60-em fast slewing telescope located at la Silla. Chile. which is dedicated to prompt optical/near-Lh follow-ups of CRB afterglows (Zerbi2001:Chincarinictal.2003:C'ovinoetal. 2004). | The is a 60-cm fast slewing telescope located at la Silla, Chile, which is dedicated to prompt optical/near-IR follow-ups of GRB afterglows \citep[][]{zerbi01,chincarini03,covino04}. |
. The aautomatically responded when BAT triggered on the X-ray burst from aand started observing with the ROSS camera ISS s after the BAT trigger. | The automatically responded when BAT triggered on the X-ray burst from and started observing with the ROSS camera 188 s after the BAT trigger. |
A series of 5. H--band. observations with | A series of 5 -band observations with |
zovLl with an intrinsic I. luminosity of about twice that of Arp 220. | $z\sim 4.1$ with an intrinsic IR luminosity of about twice that of Arp 220. |
The uncertainty in the inferred redshift is too large (o0.~0.5) for a CO search using existing instruments. but a spectroscopic verification of ILDFS50.1 and other optically faint subimin galaxies should become possible using future broadband instruments on the Green Dank Telescope or the Large Millimeter Telescope. | The uncertainty in the inferred redshift is too large $\sigma_z \sim 0.5$ ) for a CO search using existing instruments, but a spectroscopic verification of HDF850.1 and other optically faint submm galaxies should become possible using future broadband instruments on the Green Bank Telescope or the Large Millimeter Telescope. |
SAM. J09429--4658 is another historically significant SCUBA galaxy. whose extremely dusty nature and red color (A.=19.4 mag. J—A>6) was clearly demonstrated [or the first lime bv the combined analvsis of the high angular resolution VLA continuum imaging and deep near-IR imaging2000). | SMM J09429+4658 is another historically significant SCUBA galaxy, whose extremely dusty nature and red color $K=19.4$ mag, $I-K>6$ ) was clearly demonstrated for the first time by the combined analysis of the high angular resolution VLA continuum imaging and deep near-IR imaging. |
. Our new photometric redshilt analvsis suggests a moderately high redshilt (2,57 3.9). but the SED model fit (Figure 6)) is quite poor (AZ~4) as our SED template cannot be fully reconciled with the three existing SED measurements. | Our new photometric redshift analysis suggests a moderately high redshift $z_{ph} \sim 3.9$ ), but the SED model fit (Figure \ref{fig:4seds2}) ) is quite poor $\chi^2_n \sim 4$ ) as our SED template cannot be fully reconciled with the three existing SED measurements. |
One possible explanation is (hat this subnim galaxy is underbuminous in radio continuum as in HRIO and CUDSS14.13. and its actual redshift is smaller. | One possible explanation is that this submm galaxy is underluminous in radio continuum as in HR10 and CUDSS14.18, and its actual redshift is smaller. |
SMAL JOO266+1708 is another clear example of an optically faint (A.—22.5 mag) and red submnm galaxy. whose extreme properties are clearly established by (he combined analysis of the high angular resolution millimeter continuum observations al OVRO and deep imaging using the Ixeck telescope2000).. | SMM J00266+1708 is another clear example of an optically faint $K=22.5$ mag) and red submm galaxy, whose extreme properties are clearly established by the combined analysis of the high angular resolution millimeter continuum observations at OVRO and deep near-IR imaging using the Keck telescope. |
The derived photometric reclshilt Of zyy=3.50 makes it one of the highest redshift objects examined in this study. perhaps only the second after IIDES50.1. and it has the second largest intrinsic FIR luminosity after SAM 02399-0136. | The derived photometric redshift of $z_{ph}=3.50$ makes it one of the highest redshift objects examined in this study, perhaps only the second after HDF850.1, and it has the second largest intrinsic FIR luminosity after SMM $-$ 0136. |
CUDSS14.1 is the brightest SCUBA source discovered in the CERS 14hr field by(2000). | CUDSS14.1 is the brightest SCUBA source discovered in the CFRS 14hr field by. |
. Its optical counterpart has been identified using hieh angular resolution imagine αἱ the VLA and IRAM interferometer2000).. but it has not vel vielded a spectroscopic redshift. | Its optical counterpart has been identified using high angular resolution imaging at the VLA and IRAM interferometer, but it has not yet yielded a spectroscopic redshift. |
All four signilicant detection points lie along the moclel SED while all upper limits are consistent with the model. | All four significant detection points lie along the model SED while all upper limits are consistent with the model. |
The derived photometric recdshilt of 2,7=2.06 is in good agreement with its optical photometric redshift of ze21999). | The derived photometric redshift of $z_{ph}=2.06$ is in good agreement with its optical photometric redshift of $z\sim 2$. |
. Using the SED model of and optical to near-IR color. estimate its redshift to lie between 2 and 4.5. | Using the SED model of and optical to near-IR color, estimate its redshift to lie between 2 and 4.5. |
Lockmana50.1 is the brightest SCUBA source detected in the Lockman Hole ISOPHOT survey reeion wilh an extremely red (A20. 1—NK> 6.2) and extended optical counterpart2001). | Lockman850.1 is the brightest SCUBA source detected in the Lockman Hole ISOPHOT survey region with an extremely red $K\sim 20$, $I-K>6.2$ ) and extended optical counterpart. |
. Comparing the observed SED to that of Arp 220. Lutz et al. | Comparing the observed SED to that of Arp 220, Lutz et al. |
estimate its redshift to be around 3. which agrees well with our photometric redshilt τρ—2.72. | estimate its redshift to be around 3, which agrees well with our photometric redshift $z_{ph}=2.72$. |
SAM. J14009+0252 is one of the submm sources identified in the Abell 1835 field by with a faint optical counterpart (A.~ 21). | SMM J14009+0252 is one of the submm sources identified in the Abell 1835 field by with a faint optical counterpart $K\sim 21$ ). |
We derive τμ= 1.30. but the SED fit shown in Figure 6. [or the dust spectrum is poor (42= 2.3). | We derive $z_{ph}=1.30$ , but the SED fit shown in Figure \ref{fig:4seds2}
for the dust spectrum is poor $\chi^2_n = 2.3$ ). |
Ivison et al. | Ivison et al. |
rejected | rejected |
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