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These samples certainly emphasize (he need [or more extensive detailed tests not dominated by small number statistics. | These samples certainly emphasize the need for more extensive detailed tests not dominated by small number statistics. |
Another issue in measuring cluster abundance or evaluating the cluster mass function [rom a weak lensing survey is. of course. the low efficiency of the weak lensing maps. | Another issue in measuring cluster abundance or evaluating the cluster mass function from a weak lensing survey is, of course, the low efficiency of the weak lensing maps. |
At least e50% of the weak lensing peaks with v>3.7 in both the DLS and GTO2dee? fields do nol correspond to individual rich clusters (massive halos): (μον are either superpositions in the redshift survey or there is no candidate svstem at all. | At least $\sim 50$ of the weak lensing peaks with $\nu > 3.7$ in both the DLS and $^2$ fields do not correspond to individual rich clusters (massive halos); they are either superpositions in the redshift survey or there is no candidate system at all. |
In other words. the efficiency of | In other words, the efficiency of |
obtain the reddenine-corrected fluxes elven in Table 3. | obtain the reddening-corrected fluxes given in Table 3. |
The [5 i| doublet ratio. Is UJAGTIT) 1ί [5 UJAGT31) ~0.06. eives an electron density of iczT0 7. | The [S ] doublet ratio, $I$ ([S $\lambda$ $/I$ ([S $\lambda$ 6731) $\simeq 1.35 \pm
0.06$ , gives an electron density of $n_{\rm e} \simeq 70$ $^{-3}$. |
The [N u[A6583/IIo. |S n](AGT1T. |. A6731)/IIo. aud |O II /TE. line ratios indicate that massive stars are the source of the ionizing photons (Veilleux Osterbrock 1987). | The [N $\lambda$ $\alpha$, [S $\lambda$ 6717 + $\lambda$ $\alpha$, and [O $\lambda$ $\beta$ line ratios indicate that massive stars are the source of the ionizing photons (Veilleux Osterbrock 1987). |
For a distance of 558 Mpc (2:=0.091). the reddenine-corrected IL) Iuniuositv. LOT) =9.93.«Lo? ere 1. yields a production rate of Lyman continua photous of N(Lve) ~2.081003 photous |. equal to the output of ~GLO ΟΥ stars (Vacca. Carman, Shull 1996). | For a distance of 558 Mpc $z = 0.091$ ), the reddening-corrected $\beta$ luminosity, $L$ $\beta$ ) $\simeq 9.93 \times 10^{39}$ erg $^{-1}$, yields a production rate of Lyman continuum photons of $N$ (Lyc) $\simeq 2.08 \times 10^{52}$ photons $^{-1}$, equal to the output of $\sim 610$ O5V stars (Vacca, Garmany, Shull 1996). |
Finally. the ratio Ro,= ΠΟ 11|A3727) FUO uyAl959) | FO 1]A50073] /Z(HE 2): can be used to estimate the metallicity of the galaxy. | Finally, the ratio $R_{23} \equiv$ $I$ ([O $\lambda$ 3727) + $I$ ([O $\lambda$ 4959) + $I$ ([O $\lambda$ $I$ $\beta$ ) can be used to estimate the metallicity of the galaxy. |
Using the corrected fluxes elven in Table 3. we obtain Ro;zL604 0.06. | Using the reddening-corrected fluxes given in Table 3, we obtain $R_{23} \simeq 4.60 \pm 0.06$ . |
The eiupirical calibration of Eduimuds Pagel (198L. see also Vacca Conti 1992) then vields 12 | log (O/T) =SOL. or an oxvecn abundance with respect to solar of [O/T =0.29. ( | The empirical calibration of Edmunds Pagel (1984, see also Vacca Conti 1992) then yields 12 + log (O/H) $\simeq 8.64$, or an oxygen abundance with respect to solar of [O/H] $\simeq -0.29$. ( |
We adopted the solar oxygen abundance value (Ο/Π). = 8.303410.1 of Meyer 1985.) | We adopted the solar oxygen abundance value $_\odot$ = $\times 10^{-4}$ of Meyer 1985.) |
The derived properties of galaxy D are typical of those of star-forming ealaxies. | The derived properties of galaxy B are typical of those of star-forming galaxies. |
Iu the two-dimensional spectrum taken at PA 1567 on 1997 April 15. we serendipitously found. coutiunua iud lue enüssion from three additional faint galaxies. | In the two-dimensional spectrum taken at PA $156^\circ$ on 1997 April 15, we serendipitously found continuum and line emission from three additional faint galaxies. |
Their spectra are shown in Fieure 3. where the galaxies lave been denoted as 1. IL aud TW (brightest to faintest). | Their spectra are shown in Figure 3, where the galaxies have been denoted as I, II, and III (brightest to faintest). |
Contin oenüsson from another ναν faint ealaxy (denoted as IV) is also scen in Figure 3. | Continuum emission from another very faint galaxy (denoted as IV) is also seen in Figure 3. |
The spectra of galaxies L IL aud III were extracted from three rows (0.6173) centered on the observed peak brightness of each ealaxy. | The spectra of galaxies I, II, and III were extracted from three rows ) centered on the observed peak brightness of each galaxy. |
Iu order to identify these galaxies. we show a close up of our AR-baud CCD nage in Figure 1. | In order to identify these galaxies, we show a close up of our $R$ -band CCD image in Figure 4. |
The faintest ealaxy IV in Figure 3 appears to be as bright as galaxy I in this amaee. | The faintest galaxy IV in Figure 3 appears to be as bright as galaxy I in this image. |
Thus the faintuess iu Figure 3 may be due to our slit covering only a portion of this ealaxy. | Thus the faintness in Figure 3 may be due to our slit covering only a portion of this galaxy. |
Since there are no conviuciug enuission lines in this spectrum. we will uot discuss this galaxy further iu this paper. | Since there are no convincing emission lines in this spectrum, we will not discuss this galaxy further in this paper. |
The coordinates of galaxies 1, IL. aud IIT are eiven iu Table 1. | The coordinates of galaxies I, II, and III are given in Table 4. |
We also eive rough estimates of R magnitudes which were measured using the spectra shown iu Figure 3. | We also give rough estimates of $R$ magnitudes which were measured using the spectra shown in Figure 3. |
The observed magnitudes were corrected for 0.1 mag of Galactic extinction iu the BR baud (0.15 mag in the D baud. Burstein Iciles 1981) usine the average Galactic extinction law of Seaton (1979). | The observed magnitudes were corrected for 0.1 mag of Galactic extinction in the R band (0.15 mag in the B band, Burstein Heiles 1984) using the average Galactic extinction law of Seaton (1979). |
Note that we do not apply I. correction in the estimates of absolute R magnitiucles. | Note that we do not apply K correction in the estimates of absolute $R$ magnitudes. |
Since it is unlikely that our slit covered the cutive images of the galaxies. the maeuitudes given in Table { iav be upper limits. | Since it is unlikely that our slit covered the entire images of the galaxies, the magnitudes given in Table 4 may be upper limits. |
Identifications. fluxes. and equivalent widths of the Cluission lines. as well as the redshifts derived frou. the observed wavelengths of the lines. are sumunarizect in table 5. | Identifications, fluxes, and equivalent widths of the emission lines, as well as the redshifts derived from the observed wavelengths of the lines, are summarized in table 5. |
All the ealasies are located at τσ0.5. | All the galaxies are located at $z
\simeq 0.5$. |
Note that ealaxy III has ouly oue obvious emission line which we identify as [ο 11|A3727. | Note that galaxy III has only one obvious emission line which we identify as [O $\lambda$ 3727. |
We sce a faint cmiission feature which may be ILj in Figure 3 but we cannot confirm it unanibiguously. | We see a faint emission feature which may be $\beta$ in Figure 3 but we cannot confirm it unambiguously. |
The widths of the emission lines are uot resolved by our spectroscopy. | The widths of the emission lines are not resolved by our spectroscopy. |
Our resolution (13 A)) vields iui upper Iinüt to the line widths of = 700 ku + for the redshifted |O Π) A3727 line. | Our resolution (13 ) yields an upper limit to the line widths of $\simeq$ 700 km $^{-1}$ for the redshifted [O ] $\lambda$ 3727 line. |
Since galaxy I shows several strong cinission lines. we Investigate its ionization propertics iu nore cetail. | Since galaxy I shows several strong emission lines, we investigate its ionization properties in more detail. |
Iu Figure 5. we show the location of galaxy I in the |O H[A3727/|O. ni[ADO07. versus [O. jA5007/ID> excitation diagram (Baldwin. Philips. Terlevich 1981). | In Figure 5, we show the location of galaxy I in the [O $\lambda$ 3727/[O $\lambda$ 5007 versus [O $\lambda$ $\beta$ excitation diagram (Baldwin, Phillips, Terlevich 1981). |
For coniparison. we plot a solid curve represcuting a sequence for II regions photoionized by massive stars. as well the data for nearby iregulaur and spiral galaxies taken frou Ixeunicutt (1992). | For comparison, we plot a solid curve representing a sequence for H regions photoionized by massive stars, as well the data for nearby irregular and spiral galaxies taken from Kennicutt (1992). |
Since galaxy Ties iu the same area of the excitation diagram as these data. its excitation propertics are typical of local star forming galaxies. | Since galaxy I lies in the same area of the excitation diagram as these data, its excitation properties are typical of local star forming galaxies. |
Iowewer. it should be noted that we have made no correction for the (unknown) mterual extinction. which could sieuifcantlv change the [O 113727] line flux. | However, it should be noted that we have made no correction for the (unknown) internal extinction, which could signifcantly change the [O $\lambda$ 3727] line flux. |
We also see uo evideuce for a broad component at the base of the TD} line. but our S/N is not sufficient to allow a definitive statement reearding the preseuce or absence of such a feature. | We also see no evidence for a broad component at the base of the H $\beta$ line, but our S/N is not sufficient to allow a definitive statement regarding the presence or absence of such a feature. |
The spectimm of galaxy Dexlibits a very blue contin as well as stroug [Ne HYA3869 and Mg11 A2798 eimission ines. | The spectrum of galaxy I exhibits a very blue continuum as well as strong [Ne $\lambda$ 3869 and Mg $\lambda$ 2798 emission lines. |
The iutensities relative to that of I> are 0.51 aud 1.53. respectively Gvith uo reddenius corrections). | The intensities relative to that of $\beta$ are 0.51 and 1.53, respectively (with no reddening corrections). |
The Ne HJA3869 cussion liue is contaminated by a strong right skv line (see Figure 3). | The [Ne $\lambda$ 3869 emission line is contaminated by a strong night sky line (see Figure 3). |
Although the [Ne 11|A3869 cluission line has been detected iu a number of star-ornüue galaxies. Meg A2798 in enission 1s more typical of ACN (ce... Storchi-Beremann. Riuney. Challis 1995). | Although the [Ne $\lambda$ 3869 emission line has been detected in a number of star-forming galaxies, Mg $\lambda$ 2798 in emission is more typical of AGN (e.g., Storchi-Bergmann, Kinney, Challis 1995). |
However, these emission lues may be generated i shock-weated reeions (Dopita Sutherland 1996). aud it is xossible that the ionized eas in galaxy Iis affected by some wind activity by massive stars and their desceudeuts (c.e.. WolfRavet stars aud supernovae). | However, these emission lines may be generated in shock-heated regions (Dopita Sutherland 1996), and it is possible that the ionized gas in galaxy I is affected by some wind activity by massive stars and their descendents (e.g., Wolf-Rayet stars and supernovae). |
Furthermore. although Me A2798 emission is ecucrally rare m starburst svstenis. some starbursts do in fact exhibit this cmission feature (c.g. the blue compact dwarf galaxies NGC 1510 and Tol 1921-116: RKiuuev et al. | Furthermore, although Mg $\lambda$ 2798 emission is generally rare in starburst systems, some starbursts do in fact exhibit this emission feature (e.g. the blue compact dwarf galaxies NGC 1510 and Tol 1924-416; Kinney et al. |
1993: Storchi-Dergniaun. liuuev. Challis 1995). | 1993; Storchi-Bergmann, Kinney, Challis 1995). |
While we have not detected anv high ionization lines such as [Ne v]. nor see any evidence for a broad line component in the detected Ines. our observations still cannot rule out an ACN origin. | While we have not detected any high ionization lines such as [Ne ], nor see any evidence for a broad line component in the detected lines, our observations still cannot rule out an AGN origin. |
Nonetheless. we will proceed under the assuuption that the emission lines are due to a starburst. | Nonetheless, we will proceed under the assumption that the emission lines are due to a starburst. |
The emission line huminosities can be used to estimate the star formation rate (SFR) iu these galaxies. | The emission line luminosities can be used to estimate the star formation rate (SFR) in these galaxies. |
We use the calibration based on [O 11] Iainosity derived by Ieunicutt (19092: SFR 25«10H L(O uj) M. x Ly) | We use the calibration based on [O ] luminosity derived by Kennicutt (1992; SFR $\simeq 5
\times 10^{-41}$ ([O ]) $M_\odot$ $^{-1}$ ). |
The resulta are sunuuarized in Table 6. | The results are summarized in Table 6. |
We obtain SFR ~1M. ! for galaxy I while SFR ~2A. y+ for galaxies IT and ΠΠ. | We obtain SFR $\simeq 4 M_\odot$ $^{-1}$ for galaxy I while SFR $\simeq 2 M_\odot$ $^{-1}$ for galaxies II and III. |
These SFRs are smaller by a factor of a few than that for a normal galaxy (ie. a so-called Z galaxy). cLOAL. ! (e... Cowie. IIu. Soneaila 1995). | These SFRs are smaller by a factor of a few than that for a normal galaxy (i.e., a so-called $L^*$ galaxy), $\simeq 10 M_\odot$ $^{-1}$ (e.g., Cowie, Hu, Songaila 1995). |
The estimated paraimcters for galaxies I. IL. aud III indicate that they are so-called faint blue galaxies. similar to those fouud in the previous deep survey programs (6.8.. Cowie. Songaila. IIu 1991: see for reviews. too Ixrou 1992: Ellis 1997). | The estimated parameters for galaxies I, II, and III indicate that they are so-called faint blue galaxies, similar to those found in the previous deep survey programs (e.g., Cowie, Songaila, Hu 1991; see for reviews, Koo Kron 1992; Ellis 1997). |
Thev may be alsorelated to the class of compact narrow enuüssion line galaxies at iuteriuediate redshifts (soo ct al. | They may be alsorelated to the class of compact narrow emission line galaxies at intermediate redshifts (Koo et al. |
1995: Corian et al. | 1995; Guzmánn et al. |
1996). | 1996). |
The observed separations between the galaxies is 1..correspouding to a linear separation ~ 60 kpe. | The observed separations between the galaxies is ,corresponding to a linear separation $\simeq$ 60 kpc. |
This is typical of the mean separations found iu compact eroups of galaxies (Ilickson et al. | This is typical of the mean separations found in compact groups of galaxies (Hickson et al. |
1992). | 1992). |
Such eroups frequently coutain inultiple star forming members (e.g. Coziol et al. | Such groups frequently contain multiple star forming members (e.g. Coziol et al. |
1998). | 1998). |
Tlowever. since Πο from ealaxy C prohibits | However, since light from galaxy C prohibits |
10 presence or absence of “X-structure’. | the presence or absence of `X-structure'. |
The values of X for 1e galaxies in their sample ranged (rom 0 for galaxies with =10 [ine structure to 7.6 indicating the largest amount of fine garucture observed. | The values of $\Sigma$ for the galaxies in their sample ranged from 0 for galaxies with no fine structure to 7.6 indicating the largest amount of fine structure observed. |
NCC 1700 was assigned a value of 3.70. | NGC 1700 was assigned a value of 3.70. |
-"hey found that X correlated with a galaxys residual fron 10 mean colourmagnitude relation. | They found that $\Sigma$ correlated with a galaxy's residual from the mean colour–magnitude relation. |
Using this fact. and relating galaxy colours to ages via a star formation model. rey estimated. the time since the merger event. | Using this fact and relating galaxy colours to ages via a star formation model, they estimated the time since the merger event. |
We quote jer most representative age of 6.0 Gwe with an error of Gyr in Table 1. | We quote their most representative age of 6.0 Gyr with an error of $\pm2.3$ Gyr in Table 1. |
We note however that the mocels usec in ? assumed a solar metallicity starburst. | We note however, that the models used in \scite{ss92} assumed a solar metallicity starburst. |
Wf the stellar »opulation is super-solar. as suggested by Fig. 13.. | If the stellar population is super-solar, as suggested by Fig. \ref{fig:grids}, |
then the 7 age is expected to be a slight over-estimate. | then the \scite{ss92} age is expected to be a slight over-estimate. |
The age5 estimate of ? comes from measurements. of the stellar velocity. Geld of NGC 1700. | The age estimate of \scite{statler96} comes from measurements of the stellar velocity field of NGC 1700. |
Thev found. that within 2.55. (S.S kpe) the galaxy is kinematically well mixed. | They found that within $\sim2.5r_{\mathrm{e}}$ (8.8 kpc) the galaxy is kinematically well mixed. |
Εις constrains the time since the last major merger event to be z2.7 Cyr to allow sullicient time for phase mixing and ciffercntial precession. | This constrains the time since the last major merger event to be $\geq2.7$ Gyr to allow sufficient time for phase mixing and differential precession. |
Their observations also define an upper limit for the time since the merger. | Their observations also define an upper limit for the time since the merger. |
The asvmmetric photometric and kinematic signatures at larger radii preclude a merger age greater than 5.3 Cor. otherwise ese features would. have relaxed. and. disappeared. | The asymmetric photometric and kinematic signatures at larger radii preclude a merger age greater than 5.3 Gyr, otherwise these features would have relaxed and disappeared. |
Thus we quote an age of 2.75.3 vr. | Thus we quote an age of 2.7–5.3 Gyr. |
The same paper argues that 1ο counter rotating core and boxy features (the latter of =vhich we attribute to the tidal tail-like structures) could =100 have been created by the same merger event and. that 10 Observec form o£ NGC 171 must have arisen from the =nereer of at least three separate stellar systems. | The same paper argues that the counter rotating core and boxy features (the latter of which we attribute to the tidal tail-like structures) could not have been created by the same merger event and that the observed form of NGC 1700 must have arisen from the merger of at least three separate stellar systems. |
However. un was not clear whether these events occurred. sequentially or simultaneously. | However, it was not clear whether these events occurred sequentially or simultaneously. |
H£ the two tail-like structures are indeed eenuine tidal tails this would suggest that the history of GC 1700 ias included at least one major merger event involving two approximately equal mass disce. galaxies. | If the two tail-like structures are indeed genuine tidal tails this would suggest that the history of NGC 1700 has included at least one major merger event involving two approximately equal mass disc galaxies. |
Lf he kinematically distinct core (INDCO) was formed prior to his event it would. have probably been disrupted: during he ensuing violent relaxation processes. | If the kinematically distinct core (KDC) was formed prior to this event it would have probably been disrupted during the ensuing violent relaxation processes. |
1 the KDC was indeed. formed. by a separate process it must have resulted rom a subsequent minor merger (e.g. 7)) or an interaction (sce ?)) some time alter the major merger event that created NGC 1700 and its tails. | If the KDC was indeed formed by a separate process it must have resulted from a subsequent minor merger (e.g. \pcite{balcells90}) ) or an interaction (see \pcite{thomson90}) ) some time after the major merger event that created NGC 1700 and its tails. |
Neither of these suggestions are very appealing. | Neither of these suggestions are very appealing. |
Alternatively. the tidal structures seen in Fig. | Alternatively, the tidal structures seen in Fig. |
5r could be interpreted as plumes which have arisen from the infall of a small disc galaxy into a pre-existing elliptical. | \ref{fig:resid} could be interpreted as plumes which have arisen from the infall of a small disc galaxy into a pre-existing elliptical. |
AX recent study by ? showed that a galaxv's devia from the Pundamenal Plane (FP) correlated with its age. albeit with large scatter. | A recent study by \scite{FP1} showed that a galaxy's deviation from the Fundamental Plane (FP) correlated with its age, albeit with large scatter. |
The EP residual is definec Ray.Aly.fie)=2log(ma)|0.286Mp0.255.3.101 (?).. | The FP residual is defined as $R(\sigma_0,M_{\mathrm{B}},\mu_{\mathrm{e}})=2\log(\sigma_0)+0.286M_{\mathrm{B}}+0.2\mu_{\mathrm{e}}-3.101$ \cite{prugniel96}. |
Young ellipticas [al below the FP (ie. have negaive IXiduals) and. evolve towards it until they lie on the pr at an age of about 10 Gar. | Young ellipticals fall below the FP (i.e. have negative residuals) and evolve towards it until they lie on the FP at an age of about 10 Gyr. |
Older ellipticals tend. to lic above the FP. | Older ellipticals tend to lie above the FP. |
LH is hus »ossible to use the FP residual to approximately age date an elliptical galaxy. | It is thus possible to use the FP residual to approximately age date an elliptical galaxy. |
Ehe FP resicual [tx NGC 1700 is lay.Mp.fie)=0.37. | The FP residual for NGC 1700 is $R(\sigma_0,M_{\mathrm{B}},\mu_{\mathrm{e}})=-0.37$. |
This corresponcls ic, an age of 1242.0 Gyr based on the fit in ?.. | This corresponds to an age of $1.2\pm2.0$ Gyr based on the fit in \scite{FP1}. |
This is slightly vounger but comparable to the other age estimates discussed in this section. | This is slightly younger but comparable to the other age estimates discussed in this section. |
2>clore directly. comparing the dillerent age estimates. one 5iould bear in münd that the dating methods may be measuring cdillerent tine-scaes and may have large errors associated with them. | Before directly comparing the different age estimates, one should bear in mind that the dating methods may be measuring different time-scales and may have large errors associated with them. |
As mentioned above. the stellar 5»ectroscopy— methods Measure the central “luminosity ycighted average age’ which means they are. dominated by the last burst of star formation. although the old μαcllar population also contributes. | As mentioned above, the stellar spectroscopy methods measure the central `luminosity weighted average age' which means they are dominated by the last burst of star formation, although the old stellar population also contributes. |
Thus the true age of re starburst may be slightly. less than the spectroscopic ee. | Thus the true age of the starburst may be slightly less than the spectroscopic age. |
The voung GC's were possibly formed. in the same ar formation event as the galaxy. starburst. ancl as such 51ould give a similar age as the stellar spectroscopy. | The young GCs were possibly formed in the same star formation event as the galaxy starburst, and as such should give a similar age as the stellar spectroscopy. |
The age estimate from the ? trend is also based on a star formation time-scale. | The age estimate from the \scite{FP1} trend is also based on a star formation time-scale. |
The stellar dynamics. fine structure ancl tail fraction ages estimate the time since the last merger. event. | The stellar dynamics, fine structure and tail fraction ages estimate the time since the last merger event. |
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