Introduction

Relativistic jets in active galactic nuclei (AGN) are believed to originate from the vicinity of a supermassive black hole (SMBH), which is located at the center of the galaxy. Understanding the detailed physical processes of jet formation, acceleration, collimation, and subsequent propagation has been one of the major quests in modern astrophysics (see, e.g., Boccardi et al. 2017;Blandford et al. 2019 and references therein for recent reviews).

Extensive studies on these topics have been carried out over the last several decades, in particular by using the technique of millimeter-wave (mm) very long baseline interferometry (VLBI), which provides especially high angular resolution and can penetrate regions that are opaque at longer wavelengths. Notably, recent Event Horizon Telescope (EHT) observations of M 87 at 1.3 mm (230 GHz) have revealed a ring-like structure on event horizon scales surrounding the SMBH, interpreted as the black hole "shadow" (Event Horizon Telescope Collaboration et al. 2019a,b,c,d,e,f; hereafter Papers I-VI). Although the EHT results for M 87 provide an important step toward understanding jet formation near a BH and in AGN systems in general, the first EHT images of M 87 do not yet provide a direct connection between the SMBH and the large-scale jet. Therefore, imaging of fine-scale structures of AGN jets close to the SMBHs still remains crucial in order to better understand the accretion and outflow activities. Also, a more comprehensive understanding of AGN jet formation will require systematic studies over a wider range of AGN classes, given intrinsic differences such as luminosity, accretion rate, and environmental effects (e.g., Yuan & Narayan 2014). We also note that M 87 and the Galactic Center SMBH Sagittarius A* are relatively weak sources of γ-ray emission (e.g., Lucchini et al. 2019), while many other AGN produce prominent and variable high-energy emission, often from compact regions in their jets (e.g., Madejski & Sikora 2016). Therefore, studies of the high-power, high-luminosity AGN also provide more clues regarding γ-ray emission mechanisms (see, e.g., Blandford et al. 2019 for a review).

Unfortunately, most high-power AGN are located at much larger luminosity distances than M 87 and Sgr A*. Observing frequencies up to 86 GHz have thus limited us in the past to studying relatively large-scale jet morphology and evolution in many different types of AGN. However, it is only with the EHT at 230 GHz and beyond that the finest details at the base of those gigantic dynamic structures become accessible. Combined with other VLBI arrays, for example the Very Long Baseline Array (VLBA) or Global Millimeter VLBI Array (GMVA) at 86 GHz, the EHT can also connect the innermost regions of jets with the downstream sections, revealing detailed profiles of the jet collimation and locations of the collimation profile changes to better constrain jet collimation and propagation theories (e.g., Asada & Nakamura 2012;Hada et al. 2013).

The blazar 3C 279 (1253-055) is one of the sources that provided the first evidence of rapid structure variability (Knight et al. 1971) and apparent superluminal motions in compact AGN jets (Whitney et al. 1971;Cohen et al. 1971). Since the discovery of the apparent superluminal motions, the detailed structure of the radio jet in 3C 279 has been imaged and its properties have been studied by a number of VLBI observations until the present day. The 3C 279 jet consists of a compact core and straight jet extended from subparsec (sub-pc) to kiloparsec (kpc) scales. The compact core has high apparent brightness temperature at centimeter wavelengths (T B,app 10 12 K; see, e.g., Kovalev et al. 2005). Both the core and the extended jet show high fractional linear polarization ( 10%), and strong circular polarization on the order of ∼ 1% is also detected in the core region at ≤ 15 GHz (e.g., Homan & Wardle 1999;Homan & Lister 2006;Homan et al. 2009b) and≤ 43 GHz (Vitrishchak et al. 2008). The extended jet components show various propagation speeds (bulk Lorentz factor Γ ∼ 10 -40; e.g., Bloom et al. 2013;Homan et al. 2015;Jorstad et al. 2017), indicating the presence of not only underlying bulk plasma motions, but also patterns associated with propagating shocks or instabilities. Interestingly, the inner jet components of 3C 279 often display various position angles (see, e.g., Homan et al. 2003;Jorstad et al. 2004 and references therein), but later on such components tend to align with the larger-scale jet direction while propagating toward the jet downstream (e.g., Kellermann et al. 2004;Homan et al. 2009a). Based on the small viewing angle of the 3C 279 jet of θ ∼ 2 • (Jorstad et al. 2017), the misaligned jet components are often modeled as spatially bent (and perhaps helical) jet structures, in which the jet Lorentz factor is constant along the outflow but the jet viewing angle changes (e.g., Abdo et al. 2010;Aleksić et al. 2014). We also note that jet bending on VLBI scales is common in many blazar jets (e.g., Hong et al. 2004;Lobanov & Roland 2005;Zhao et al. 2011;Perucho et al. 2012;Fromm et al. 2013). For the innermost region of the 3C 279 jet ( 100 µas ∼ 0.65 pc projectedfoot_0 ), earlier pilot VLBI studies at 230 GHz revealed a complex microarcsecond-scale substructure within the nuclear region of the milliarcsecond scale jet (Lu et al. 2013;Wag-ner et al. 2015). However, the (u, v) coverage, and therefore the imaging fidelity, of these observations was very limited. We also note that 3C 279 is well known for its highly time-variable flux densities, from radio to γ-rays (e.g., Chatterjee et al. 2008;Abdo et al. 2010;Aleksić et al. 2014;Kiehlmann et al. 2016;Rani et al. 2018;Larionov et al. 2020), while the exact locations of the gamma-ray emission zones are often controversial (e.g., Patiño-Álvarez et al. 2018Patiño-Álvarez et al. , 2019)). In particular, 3C 279 shows flux density variations down to minute timescales, which are often difficult to interpret given the size scales and Doppler factors inferred from radio VLBI observations (e.g., Ackermann et al. 2016).

In April 2017, 3C 279 was observed with a significantly expanded EHT array over four nights. The EHT 2017 observations result in new and more detailed maps of the core region of 3C 279, providing an angular resolution of 20 µas, or ∼ 0.13 pc (corresponding to ∼ 1700 R s for a SMBH of mass M BH ∼ 8 × 10 8 M ; Nilsson et al. 2009). This paper presents the main results from the EHT observation in 2017 and their scientific interpretations. In §2 we briefly describe the observations, imaging procedures, and model-fitting techniques. In §3 the source images and model-fit parameters are presented. In §4 we discuss some physical implications of the peculiar compact jet structure, in relation to the observed rapid variation of the source structure and brightness temperature. §5 summarizes our results. Throughout this paper we adopt a cosmology with H 0 = 67.7 km s The (u, v) coverage is shown in Fig. 1. The high data recording rate of 32 Gbps (corresponding to a total bandwidth of 2 GHz per polarization per sideband) allowed robust fringe detections up to a ∼ 8.7Gλ baseline length, including the SPT, which significantly improved the fringe spacing toward 3C 279 in the northsouth direction. The correlated data were then calibrated using various radio astronomical packages and validated through a series of quality assurance tests (see Paper III for details). The fluxcalibrated visibility amplitude distributions are shown in Fig. 2. Notes. The beam sizes were obtained using Difmap and uniform weighting. We adopt a 20µas circular Gaussian beam for all 3C 279 CLEAN images.

Imaging and model-fitting analysis

For imaging, we used frequency-averaged visibility data from the EHT-HOPS pipeline (see Paper III and Blackburn et al. 2019).

We note that image reconstruction with 1.3 mm wavelength EHT data is particularly challenging because of the sparse (u, v) coverage, total loss of absolute atmospheric phase, and large gain fluctuations at some stations. In addition, the 2017 EHT observations lack relatively short baselines at 1 Gλ to robustly recover extended emission structure on VLBI scale at 100 µas (Paper IV). To ensure that the features we identified in our reconstructed images are robust, the source images were generated by both traditional CLEAN and newer regularized maximum likelihood algorithms implemented in the following programs: Difmap (Shepherd et al. 1994), eht-imaging (Chael et al. 2016(Chael et al. , 2018)), and SMILI (Akiyama et al. 2017a,b). We used imaging pipelines for these three programs (see Paper IV) to generate a total of 12 images of 3C 279 (i.e., one per epoch per imaging method) within a limited field of view of ∼ 100 µas due to lack of short EHT 2017 baselines (Paper IV). In all methods, emission from the further extended milliarcsecond-scale jet (Fig. 4), which lies beyond the compact EHT field of view, was represented by a single large-scale Gaussian (see Paper IV for details).

We then averaged the three pipeline images to obtain a representative image of the source at each epoch. We refer to Paper IV for the details of the imaging pipelines and image averaging procedures. In order to illustrate the EHT angular resolution toward 3C 279, we show in Table 1 the CLEAN beam sizes of the EHT 3C 279 data calculated by Difmap.

In order to parameterize bright and compact features in the source, we also performed Gaussian model-fitting analyses in two distinct ways. The first is the traditional VLBI model-fitting procedure (DIFMAP modelfit, which employs the Levenberg-Marquardt algorithm for non-linear fits) to reconstruct a static model with more than six components on each observation day. Related components were then identified and the evolution in their relative positions measured.

The second method utilizes Themis, an EHT-specific analysis package, using a parallel-tempered, affine invariant Markov chain Monte Carlo sampler (Broderick et al. in prep., and references therein). In this case, a fully time-variable, ten-component (nine compact and one large-scale) Gaussian component model was reconstructed to naturally facilitate the identification of features in subsequent observations and directly reconstruct their evolution. From this time variable model, component parameters and uncertainties are reconstructed for individual days. Additional descriptions of the underlying model and Themis analysis can be found in Appendix A (also see Paper VI for more general details for the EHT model-fitting and model-comparison analysis).

-50 0

Relative RA (μas) Relative Dec (μas) 7. Same as Fig. 6, but using C0-1 (top) and C0-2 (bottom) as kinematic references. We note the more complicated motions of other jet features in both panels compared to Fig. 6. when C0-2 is chosen as the reference. Therefore, choosing C0-0 as the kinematic reference provides a smoother transition of the kinematics from the inner EHT scale to the outer large jet (see Fig. 3), and also helps avoid unnecessary complexity in the interpretation given the limited available data, although this choice alone does not allow us to determine which of the three C0 subcomponents remains more stationary in time (see §4.3 for more discussions from a physical perspective).

We also note that adopting C0-0 as the kinematic reference helps avoid false identification of the other C0 subcomponents, such as counterjet features. The expected jet-to-counterjet ratio of discrete emission features in 3C 279 can be computed as ((1 + β cos θ)/(1 -β cos θ)) m-α , where β is the jet speed in units of c; m = 2 or 3 for a continuous jet or a single component, respectively (see, e.g., Urry & Padovani 1995); and α is the optically thin spectral index (i.e., flux density S ∝ ν +α ). If we adopt α = -0.7, θ = 2 • , m = 3, and β = β app /(sin θ + β app cos θ), where β app 10 based on the observations, the expected brightness ratio is 10 10 ; however, the observed brightness ratios of the C0 subcomponents are within an order of magnitude (Table D.1). Therefore, we should expect to find no counterjet features situated to the north of the VLBI "core" (see Fig. 7), although emitting features moving in a helically bent jet could perhaps produce this apparent backward motion if the jet is closely aligned to the line of sight (see §4 for a discussion).

In addition, we further note that the VLBI core is usually defined as the most compact and brightest jet feature in the obtained images, and thus has the highest brightness temperature. It is interesting to note in Fig. 8 that the brightest component is not C0-0, but either C0-1 or C0-2, depending on the observing epochs. With this criterion, C0-1 and C0-2 might be still classified as the VLBI core. However, long-term and high-resolution observations of blazar jets find that compact and bright jet components near the VLBI core often have higher brightness temperatures than the cores determined by the jet kinematics (see, e.g., Lisakov et al. 2017;Bruni et al. 2017;Jorstad et al. 2017). Thus, identifying C0-1 and C0-2 as the potential VLBI core based on the flux density and brightness temperature may not be strongly supported in our observations. Therefore, we adopt C0-0 as the VLBI core of 3C 279 in the following analysis.

Discussion

Elongated nuclear structure

The nuclear (C0 region) structure of 3C 279 resolved at the highest 20 µas angular resolution is elongated perpendicular to the large-scale jet. This structure is seen in both independent imaging and model-fitting methods, and can be modeled as three bright features separated by ∼ 30 -40 µas. This corresponds to a projected spatial scale of ∼ 2500-3400 R s for M BH = 8×10 8 M . This morphology has not been commonly seen for 3C 279 by VLBA at 15 and 43 GHz (Jorstad et al. 2017;Lister et al. 2018). If the jet emission represents distribution of underlying synchrotron-emitting plasma, this peculiar structure can be interpreted in various ways. Below we provide four possible interpretations.

Standard jet formation scenarios suggest relativistic jet launching by either angular momentum extraction from the spinning SMBH (Blandford & Znajek 1977) or magneto-centrifugal acceleration by an accretion disk (Blandford & Payne 1982), or by both mechanisms at the same time. In this context, a spatially resolved jet base, similar to the jet base morphology found in several nearby radio galaxies, in particular with limb-brightened jets (e.g., 3C 84; Giovannini et al. 2018, Cygnus A;Boccardi et al. 2016) is also possible. However these are viewed at a much larger angle to the line of sight than for 3C 279 and could provide an edge-on view of the limb-brightened jet base or the disk (thus thin elongated geometry if the accretion flow is not a sphere but has a finite height-to-radius ratio of, e.g., H/R 1; see, e.g., Yuan & Narayan 2014). However, for 3C 279 a nearly face-on view (θ ∼ 2 • ) and thus a more rounded, thick emission geometry is expected on the sky for the base of a circular jet or the accretion flow, in contrast to the observed images which show a narrow width along the global direction of the jet. 4Table 2. Summary of geometric and dynamical properties of the jet components discussed in §4.2.1 and §4.2.2.

Notes. (a) The same Lorentz factor but different viewing angles for the jet features. For C0 subcomponents we presume the small θ case, while for C1 we presume the large θ case (see §4.2.1 and §4.2.2). (b) Assumes a constant fixed viewing angle of θ = 2 • .

C1 region

In contrast to the C0 region, the three subcomponents in C1 have comparable apparent speeds (β app ∼ 13 -15), and their position angles with respect to C0-0 are all in a narrow range of ∼ -(173 • -178 • ), which are aligned to the directions of their motion vectors (PA∼ -(160 • -180 • )). Therefore, we can reasonably presume that these components share common kinematic and geometric properties.

We could extend the analysis in §4.2.1 to the C1 region, that is assuming a constant Γ = 20 for all the components to estimate their different viewing angles. We show in Fig. 10 the same β app and δ plane but for the C1 subcomponents, which is used to constrain reasonable ranges of θ and the corresponding δ. For β app = 13 -15, there are two possible sets of parameters, which are (i) θ ∼ 6 • -8 • and δ ∼ 5 -7, and (ii) θ ∼ 1.0 • -1.5 • and δ ∼ 33-35. We note that there is a similar ambiguity in determining whether the jet bends closer to or away from the line of sight. Nevertheless, we could consider that weaker time variability of the C1-0 and C1-1 components might prefer smaller Doppler factor values (i.e., larger θ), while C1-2 shows stronger variability and thus could have larger Doppler boosting (i.e., smaller θ).

Alternatively, dynamical properties of C1 could be better estimated by simply adopting the viewing angle of the larger-scale jet (θ ∼ 2 • ; Jorstad et al. 2017) because the motions of the C1 subcomponents are nearly parallel to the jet downstream (Fig. 3). Using θ = 2 • , we obtain Γ ∼ 16 -17 and δ ∼ 24 -25.

Taken all together, the ranges of Lorentz factors for C1 are comparable to those found from the 3C 279 jet on larger scales and at longer wavelengths ( 10 3 pc or 10 5 R s projected; Bloom et al. 2013;Lister et al. 2016;Jorstad et al. 2017;Rani et al. 2018), and also those estimated from radio total flux variability (e.g., Hovatta et al. 2009). The values of θ, Γ, and δ for C1 are also summarized in Table 2.

Physical implications

In conclusion, it appears that the peculiar C0 structure could be described by a jet closely aligned to the line of sight, but bent by small angles, and the projection of the overall bent geometry to the sky. In this perspective, it is also worth noting that in a previous 230 GHz VLBI experiment on 2011 Mar 29 -Apr 4, a similar nuclear morphology was found in 3C 279 based on a model-fitting approach (Lu et al. 2013). After 2011 December, this structure became resolved by the VLBA at 43 GHz as a bright moving feature situated at a position angle of initially ∼ 150 • , and later at ∼ -170 • relative to the 43 GHz core (Alek-sić et al. 2014;Jorstad et al. 2017), confirming the jet bending scenario (the VLBA 7mm kinematics is shown in Fig. E.1 in Appendix E). Notably, the overall situation of the source in 2011 is similar to the jet geometry we discussed in §4.2.1. This suggests that the inner jet bending may commonly occur in 3C 279. In this respect it is interesting to note that a similar extremely bent jet morphology is sometimes observed in several AGN on small angular scales, especially when the object is in a flaring state at multiple wavelengths (e.g., 1156+295 - Hong et al. 2004;Zhao et al. 2011;PKS 2136+141 -Savolainen et al. 2006;OJ 287 -Agudo et al. 2012;Hodgson et al. 2017;3C 345 -Lobanov & Roland 2005;CTA 102 -Fromm et al. 2013;Casadio et al. 2015;0836+710 -Perucho et al. 2012). The flare is often interpreted as the result of an increase in Doppler beaming of the emission due to the jet bending closer to the line of sight.

There are several possible explanations for the physical origin of the jet bending. First, precession of a jet nozzle, which is induced by propagation of perturbations originating from the accretion disk and BH due to the Lense-Thirring effect (Bardeen & Petterson 1975) or even binary black holes, may display somewhat periodic jet wobbling over time. Abraham & Carrara (1998) and more recently Qian et al. (2019) suggest such a physical model for 3C 279 with a precession period of ∼ 22 yrs. However, we note that the similar erratic inner jet position angle in 2011 and 2017 seen by the EHT implies a precession period of 6 yrs if the jet wobbling is periodic. The mismatching periods would exclude this possibility. Second, it should be noted that the C0-1 component is moving toward C1 and the jet downstream (Fig. 3), and thus the component is being aligned to the larger scale jet during the observing period. The above-mentioned time evolution of the 3C 279 jet structure during 2011 also suggests that the initially bent jet component in the source later aligned with the downstream emission. The jet alignment in a single preferred direction could indicate that the outflow is being actively collimated to a pre-established channel on these small spatial scales, as similarly observed in other sources as well (see discussions in Homan et al. 2015). Third, an internally rotating jet, in which emission regions are located along strong toroidal magnetic field lines, can also reproduce gradual jet bending features in the images (e.g., Molina et al. 2014). Such a scenario is supported by theoretical studies of jet launching and propagation (see, e.g., Tchekhovskoy 2015 and references therein), and also observations of inner jet dynamics in nearby radio galaxies (e.g., Mertens et al. 2016) and smooth variation of linear polarization of many AGN jets in time and space (e.g., Asada et al. 2002;Marscher et al. 2008;Hovatta et al. 2012;Kiehlmann et al. 2016). Whether one of these scenarios is more favored than others is difficult to determine, however. Joint constraints on the model parameters with additional data, for instance with linear polarization time variability information (Nalewajko 2010), should prove fruitful.

We also note that the apparent jet speed and Lorentz factor of C1 are comparable to those in the outer jet ( §4.2.2). This suggests that intrinsic acceleration of the jet (i.e., increasing Γ) would occur upstream of C1. This puts upper limits on the spatial extension of the intrinsic jet acceleration zone of 3C 279 to be within 100 µas from the core, C0-0, which is 0.65 pc ∼ 8500 R s projected distances (∼ 19 pc ∼ 2.4 × 10 6 R s deprojected with θ = 2 • ). If the observed motions of C0-1 and C0-2 can be described by similar bulk Lorentz factors as C1, the intrinsic acceleration zone should be located at much more upstream of the jet, that is within 30 -40 µas core separation ∼ 0.20 -0.26 pc∼ 2600 -3400 R s projected distances (∼ 6 -7 pc ∼ (7.3 -9.7) × 10 4 R s deprojected).

Low brightness temperature at 230 GHz

The innermost jet brightness temperature provides us with insight about the jet plasma acceleration and radiative evolution further downstream (e.g., Readhead 1994;Marscher 1995;Schinzel et al. 2012;Fromm et al. 2013). The observed brightness temperatures of the subnuclear components within C0 are in the range of T B ∼ 10 10 -10 11 K (Fig. 8). We note that these measurements are made in the observer's frame, while the intrinsic brightness temperature of the plasma in the fluid frame, T B , is lowered by the Doppler factor δ, that is T B = T B (1+z)/δ. Considering Doppler factors of δ ∼ 20 or even larger values due to possibly curved jet geometry ( §4.2), an order of magnitude lower intrinsic brightness temperature of T B ∼ 10 9 -10 10 K is possible. This is a significantly low value compared to the long millimeter or centermeter wavelength VLBI core T B (e.g., T B > 10 12 K and T B ∼ 10 11-12 K; Kovalev et al. 2005;Jorstad et al. 2017) and also the inverse Compton limit, T B ∼ 5 × 10 11 K (Kellermann & Pauliny-Toth 1969).

It is challenging to make a straightforward interpretation of the low brightness temperature without knowing the level of synchrotron opacity at 230 GHz on 20 µas scales. Nevertheless, we provide two possible implications below.

First, it is worth noting that a trend of decreasing brightness temperature with increasing observing frequencies was previously seen in a number of AGN jet cores in the frequency range of 2 -86 GHz (e.g., Lee et al. 2016;Nair et al. 2019). This trend is often interpreted as an indication of acceleration of underlying jet outflow, based on the following considerations. In the standard model of relativistic jet (Blandford & Königl 1979), the stationary radio VLBI core structure corresponds to a region with high synchrotron opacity (τ ∼ 1) at the corresponding observing frequency. In multiwavelength VLBI observations, the opacity effect appears as a shift in the apparent core position at different frequencies (i.e., the core located more upstream of the outflow at higher frequencies), which is commonly referred to as a "coreshift" (e.g., Lobanov 1998). In this picture, higher frequency T B measurements reveal physical conditions of the jet closer to its origin, if the core T B represents surface brightness of underlying plasma outflow. In addition, we could further assume that the intrinsic brightness temperature of the plasma underlying the compact core is not frequency-dependent and remains the same over short distances (i.e., the coreshift distances). It then follows that higher T B at lower frequencies could only be explained by higher outflow speed further downstream, and consequent Doppler boosting of the emission to increase the observed T B .

It is tempting to apply this framework to the EHT 230 GHz brightness temperature measurement of 3C 279. The consistent apparent jet speeds of ∼ 10 -20c seen near the EHT core and further downstream in the jet at centimeter wavelengths, however, does not strongly support the jet acceleration scenario. Instead, the brightness temperature can simply decrease with increasing frequency if the observing frequency is higher than the synchrotron self-absorption turn-over frequency (Rybicki & Lightman 1979). The low brightness temperature at 230 GHz can therefore be alternatively understood as a signature of low opacity in the core region at 1.3 mm. The ALMA phased-array data of 3C 279 from our observations show a steep spectral index of α = -(0.6 ± 0.06) at 230 GHz (see Goddi et al. 2019), which supports this conclusion, although the ALMA measurements do not spatially resolve the microarcsecond-scale jet.

If, however, the compact VLBI core region still remains optically thick up to 230 GHz, the observed low T B values could be compared to the energy equipartition brightness temperatures (T B,eq ∼ 5 × 10 10 K; Readhead 1994), which is significantly higher than T B derived from the EHT measurements. The lower T B than the particle-to-magnetic field energy density equipartition T B,eq would then suggest that the innermost jet of 3C 279 may be magnetically dominated, contrary to previous conclusions that the jet plasma has low magnetization in 3C 279 (see, e.g., discussions in Hayashida et al. 2015;Ackermann et al. 2016). While high particle-to-magnetic energy density ratios are seen in other AGN especially during flaring activities (e.g., Jorstad et al. 2017;Algaba et al. 2018), low brightness temperature associated with potentially magnetically dominated jet is also seen in the nuclear region of other nearby AGN jets, such as M 87 (e.g., see discussions in Kim et al. 2018). According to the standard model of jet launching and propagation, magnetic energy density is expected to be dominant in a jet up to central engine distances of ∼ 10 5 R s (see Boccardi et al. 2017 and references therein). Considering the spatial scales of the EHT observations of 3C 279 (20 µas∼ 1700 R s ), it is not impossible that the observed innermost 3C 279 jet is indeed magnetic energy dominated. Nevertheless, future spectral decomposition and polarimetric analysis on the 20 µas scale with multifrequency EHT observations should determine the jet core opacity at 230 GHz, in order to provide an unambiguous interpretation of the remarkably low T B values.

Connection to γ-ray emission in 3C 279

During the EHT observations in April 2017, 3C 279 was in a highly active and variable state at γ-ray energies (see, e.g., Larionov et al. 2020). Here we briefly discuss possible implications of the innermost 3C 279 jet kinematics revealed by the EHT observations on the γ-ray emission of the source. Generally, the jet speeds measured closest to the jet origin are important in order to understand the origin of γ-ray emission in blazars. One of most plausible scenarios explaining γ-ray emission in blazars is inverse Compton (IC) scattering of seed photons from within or around the relativistic jet, while details of the IC models vary depending on the assumptions of the background photon fields (see, e.g., Madejski & Sikora 2016). Observationally, bright γray flares from blazars are often associated with emergence from the VLBI core of new, compact jet features, which travel toward the jet downstream (e.g., Jorstad & Marscher 2016). This association implies that the IC process may occur near (or even upstream of) the VLBI core. Therefore, observational constraints on the innermost jet speed is crucial for an accurate modeling of the IC process.

The EHT measurements of the proper motion suggest a minimum Lorentz factor of Γ 20 at core separations ≤ 100 µas. On the other hand, much higher Lorentz factors of Γ 100 are derived from the observations of rapid γ-ray flares (Ackermann et al. 2016). To accommodate the lower limit of Γ from the EHT observations with the larger Lorentz factors from the jet kinematics and γ-ray variability, viewing angles smaller than θ < 1 • in the region C0-1 and C0-2 may be considered. For such small angles, Doppler factors of ∼ 100 could be reached, which are sufficient to explain the observed γ-ray variability. On the other hand, we note that the continued VLBA 43 GHz monitoring of the source during 2015-2018 now suggests faster motion and higher Lorentz-factors of Γ 37 than in the past (Larionov et al. 2020). As the authors note, the local values of Γ can be even larger (e.g., ∼ 70) if fast "mini-jets" are embedded within the main flow (e.g., Giannios et al. 2009) or if multiple, tur-bulent emitting zones are present (e.g., Narayan & Piran 2012;Marscher 2014). The latter could increase the local Γ values by factors of a few. Future detailed modeling of the broadband spectral energy distribution during the EHT 2017 campaign will provide more detailed tests of the relation between the jet dynamics and the γ-ray emission in 3C 279.

Summary

In this paper, we presented the first 1.3 mm VLBI images of the archetypal blazar 3C 279 at the extreme angular resolution of 20 µas. The sharpest-ever images of 3C 279 obtained at four different epochs within a week reveal (i) peculiar substructures in the millimeter VLBI core, which can be interpreted as a bent jet, or perhaps a linear, knotty structure that could result from large-scale magnetic reconnection or plasma instabilities; (ii) rapid day-to-day closure phase variations pronounced in the longest baselines, which are associated with proper motions of ∼ 1.1 -1.7 µas day -1 and rapid flux variability; and (iii) low apparent (T B 10 11 K) and intrinsic brightness temperatures (T B 10 10 K) after correcting for Doppler boosting of at least δ ∼ 10 -20. This suggests that either the jet core is optically thin at 230 GHz, or that the innermost jet of 3C 279 is dominated by magnetic energy if the synchrotron turn-over frequency were close to 230 GHz.

More details of the source properties, such as the magnetic field configuration and detailed jet energy balance, will be subject to follow-up studies, for example by EHT full-Stokes imaging of the 3C 279 jet (EHT Collaboration et al., in prep.). As mentioned in §4.5, 3C 279 was also in a highly active and variable state at γ-ray energies during the EHT observations in April 2017. Follow-up work (EHT Collaboration et al., in prep.), combining the results from this paper with other multiwavelength data obtained close in time, will provide a more detailed understanding of the physical processes in the jet, allowing detailed tests of the potential curvature in the innermost jet, and possible jet acceleration and alternative physical scenarios, as discussed in §4.3 and §4.5. -92.0 ± 0.3 16.2 ± 0.6 6.9 +0.7 -0.6 0.64 ± 0.02 Apr 11 0.57 ± 0.05 -10.9 ± 0.8 -93.0 ± 0.3 14.6 ± 0.5 9.0 ± 0.8 0.69 +0.03 and Gaussian models (colored diamonds) for all epochs, for both observing frequency bands (HI and LO), and plotted against the baseline length and quadratic sum of the three baseline lengths in triangles (u 1 , u 2 , and u 3 ), respectively. In both panels the bottom subpanels show residuals (i.e., differences between data and model) normalized by the uncertainties of each data point.